{
    "0403/astro-ph0403538.txt": {
        "abstract": "We present the first sub-kpc $(\\sim0.7''\\approx 850$ pc$)$ resolution \\twco(1-0) molecular line observations of the ISM in the host galaxy of the QSO I~Zw~1. The observations were obtained with the BIMA mm-interferometer in its compact A configuration. The BIMA data are complemented by new observations of the \\twco(2-1) and \\thco(1-0) line with IRAM Plateau de Bure mm-interferometer (PdBI) at $0.9''$ and $1.9''$ resolution, respectively. These measurements, which are part of a multi-wavelength study of the host galaxy of I~Zw~1, are aimed at comparing the ISM properties of a QSO host with those of nearby galaxies as well as to obtain constraints on galaxy formation/evolution models. Our images of the \\twco(1-0) line emission show a ring-like structure in the circumnuclear molecular gas distribution with an inner radius of about 1.2~kpc. The presence of such a molecular gas ring was predicted from earlier lower angular resolution PdBI \\twco(1-0) observations. A comparison of the BIMA data with IRAM PdBI \\twco(2-1) observations shows variations in the excitation conditions of the molecular gas in the innermost $1.5''$ comprising the nuclear region of I~Zw~1. The observed properties of the molecular cloud complexes in the disk of the host galaxy suggest that they can be the sites of massive circumnuclear star formation, and show no indications of excitation by the nuclear AGN. This all indicates that the molecular gas in a QSO host galaxy has similar properties to the gas observed in nearby low luminosity AGNs. ",
        "introduction": "One of the open questions in galaxy evolution is concerned with the formation of active galactic nuclei (AGN) and its relation to star formation (SF) in early type galaxies. The separation of starburst and AGN components in extragalactic objects %and the investigation of the role of star formation % and the properties of the interstellar medium in the -- especially in host galaxies of quasars and QSOs -- is a key problem in the investigation of evolutionary sequences proposed for AGNs \\citep{nor88,san88,rie88,haa03}. Although it is not exactly known how the host galaxy affects the energy release of the QSO, there is statistical evidence for preceding mergers or current interaction of QSO host galaxies with a companion galaxy \\citep{mcl94,lim99}. It is believed that quasar activity is a common, but short lived, phenomenon in galaxy evolution \\citep{mcl99}. Furthermore, host galaxies of QSOs show enhanced SF activities \\citep{cou98}. The transition between Ultra Luminous Infrared Galaxies (ULIRG) and AGN seems to be continuous, all showing signs of enhanced star formation in their nuclear regions \\citep{gen98,mcl99}.  %Observations of abundances, excitation and dynamics  of the molecular %interstellar medium in the central regions of AGN are essential, since %the interstellar matter is the ''fuel'' for star formation and to a certain %amount also for the central engine. %High resolution molecular line observations allow a detailed study of %the kinematics %of the  cold  interstellar medium and the derivation of important %parameters %such  as gas masses, surface mass densities and star formation %efficiencies. Millimeter molecular line observations are ideal for studying the mechanisms which transport the molecular gas into the AGN. Observations of abundances, excitation and dynamics of the molecular interstellar medium in the central regions of AGN are essential, since the interstellar matter provides the ''fuel'' for star formation as well as the central engine. Consequently, there are large numbers of high resolution observations available for nearby objects. These observations have revealed the presence of circumnuclear starburst rings in a large number of (nearby) active and IR luminous galaxies \\citep [e.g. NGC 1068,] [] {pla91, gen95, hel03}. The extension of high angular resolution observations of molecular gas emission lines to QSO hosts is imperative in order to understand the connection between local active galaxies and high-z QSOs. The results can be used to refine model predictions of the physical conditions in high z- QSOs \\citep[such as used e.g. in] [] {com99}.  Due to the limited angular resolution of single dish mm-wavelength telescopes, % molecular line observations, only interferometric observations at these wavelengths allow insight into the morphology and into the kinematics of molecular clouds in the nuclear region of QSO host galaxies with sufficient angular resolution. High resolution molecular line observations allow for detailed kinematic studies of the cold interstellar medium and the derivation of important parameters such as gas masses, surface mass densities and star formation efficiencies. Observations of multiple CO transitions reveal excitation conditions and thus can be used to constrain the physical conditions in an observed region. They can be used to better constrain the  contribution of star formation to the total observed infrared luminosities. % Molecular line observations % , which are practically insensitive to the inclination angle of the %galaxy, %directly trace the distribution of the cold interstellar molecular gas %and its kinematics.  %The %interpretation of the far infrared emission of QSO in the framework of %star formation is problematic, since, without additional molecular %line observations, it is difficult to derive the amount of thermal %emission which contributes to the total observed infrared %luminosities. Since molecular clouds are the ``fuel'' of star %formation, and the excitation temperature of molecular clouds in SF %regions is elevated with respect to dark molecular clouds, strong %molecular line emission goes along with SF.     \\subsection{The nearby QSO I~Zw~1} \\label{IZw1_Intro}  The radio-quiet QSO I~Zw~1 is regarded as the closest QSO which can be used for detailed studies of its host galaxy (Tab. \\ref{tab:izw1}). I~Zw~1 has a  systemic velocity of 18,290 km s$^{-1}$, which corresponds to a redshift of 0.0611 \\citep{con85}, or a distance of 255 Mpc \\citep[$h_0=0.72$,] [] {spe03}. The nucleus of I~Zw~1 is extremely bright in the optical \\citep[$M_B$ of -23.45 mag,] [] {sch83, bar89}) and also has very bright X-ray emission \\citep{kru90}. I~Zw~1 belongs to the class of infrared luminous galaxies \\citep[$L_{FIR}=10^{11.3}L\\sun$,] [] {haa03}, however its QSO nature %thus can also serve as a template for high-z starburst galaxies. is not apparent in the FIR.  There is a  strong indication for interaction between I~Zw~1 and a neighboring companion \\citep[e.g.] [] {lim99,schar03}.  A direct comparison between the FIR emission and CO molecular line emission in the host galaxy reveals that the FIR continuum emission of I~Zw~1 is predominantly thermal in nature \\citep{bar89}. The FIR to \\twco(1-0) ratio shows a star formation efficiency of about 30 $L_\\sun/M_\\sun$ in the nuclear region \\citep{eck94}. \\citet{schin98} show that the molecular gas mass in the central $3.3''$ (3.7 kpc) is 2/3 of the total observed mass of $9 \\times 10^9 M_\\sun$. A kinematic analysis of their 1.9'' resolution PdBI \\twco(1-0) observations suggests that the central molecular gas is distributed in a ring of $1.5''$ ($\\sim 1.9$ kpc) diameter.  However the observations do not spatially resolve this ring. High resolution NIR observations show that the nuclear NIR spectrum is dominated by emission from the AGN, nevertheless about 25\\% of the NIR emission can be attributed to a nuclear starburst \\citep{schin98},  which might be located in the molecular gas ring.  The  star formation activity in the host galaxy of I~Zw~1 underlines the important role of starbursts in the evolution of QSOs.   %in the spatial images, which is consistent with its predicted radius. %of close to 1''. The possible presence of a starburst ring with %sub-kpc radius in the nucleus of I~Zw~1 % and its possible % role for the energy release in the nucleus a the AGN in the %nucleus emphasizes the need for subarcsecond angular resolution %molecular line observations of the QSO host galaxy. %Follow up  PdBI in \\twco(2-1) which have an % angular resolution of 0\".9 however show no spatial details, the predicted ring % is not apparent in the  \\twco(2-1) transition. Only sub-arcsecond molecular line observations would allow actual imaging of this possible starburst ring. The extension of BIMA's longest baselines up to a maximum of almost 2 km makes the array the millimeter instrument with the highest angular resolution currently available. Here we present \\twco(1-0) $\\sim 0.7''$ angular resolution observations of the QSO I~Zw~1. This corresponds to a spatial resolution of roughly 850~pc. We also present new observations with the IRAM PdBI in \\twco(2-1) and \\thco(1-0) line at $0\".9$ and $2.''1$ resolution, respectively.  \\begin{deluxetable}{lcc} %\\tabletypesize{\\scriptsize} \\tablecaption{Adopted Properties of I~Zw~1 (=PG0050+124) \\label{tab:izw1}} \\tablewidth{0pt} \\tablehead{ \\colhead{Property} & \\colhead{Value}} \\startdata Right ascension (J2000.0) & 00$^h$53$^m$34$^s$.9\\\\ Declination (J2000.0) & +12$\\degr 41'36''.2$\\\\ Inclination (deg) & 38\\\\ Position angle (deg) & 135\\\\ Systemic velocity (km s$^{-1}$) & 18,290\\\\ Distance (Mpc)& 255\\\\ $1''$ & 1.24 kpc \\enddata \\tablenotetext{-}{The coordinates refer to the phase center of the millimeter observations. The inclination and position angle are taken from \\citet{schin98}, while the systemic velocity is from \\citet{con85}. } \\end{deluxetable}   %PdBI \\thco(1-0) observations with a resolution of $\\sim 2\"$ angular %resolution. %An interpretation of these %data in the framework of multi wavelengths studies to test starburst %models will be presented in another paper. % (Schinnerer et al., in prep.) ",
        "conclusions": "Our BIMA \\twco(1-0) observations represent the first molecular line observations of the host galaxy of a QSO with sub-kpc resolution. The molecular gas disk of a QSO is for the first time spatially resolved into individual giant molecular cloud complexes. We observe giant molecular clouds with peak molecular hydrogen column densities of $2.4\\times 10^{23}$cm$^{-2}$ at a distance between $1''$ and $3''$ from the nucleus. With such high column densities these clouds could be actively forming stars.  The combination of BIMA observations with new and previously published PdBI CO data allows for the first interferometric multi-transition study in the disk of I~Zw~1. The line brightness ratios in the inner $1.5''$ are consistent with moderately dense cold GMCs, and they are not peaked at the center. This strongly suggests that the AGN has no significant effect on the central molecular material.  The exact location of the circumnuclear starburst very likely seen in nuclear NIR spectra (Schinnerer et al. 1998) is still not identified. Possible candidates are the resolved \\twco(1-0) molecular clouds seen by BIMA as well as the more thermalized material traced by the higher line ratios inside the inner $1.5''$. However, any of these molecular clouds are likely sites of a massive starburst which contributes to the observed far-infrared luminosity of $L_{FIR}=10^{11.3}L\\sun$.  The distribution and properties of the molecular gas in the host galaxy of the nearby QSO I~Zw~1 are quite similar to what is observed in nearby low luminosity AGNs \\citep[e.g.] [] {pag01, gar03}. This suggests that the ISM in QSO at high redshift might be similar to nearby low luminosity AGN as well.  Four additional attempts to re-observe I~Zw~1 in \\twco(1-0) with BIMA in the A configuration all failed to provide useful data, due to insufficient phase coherence of the observations. This demonstrates that the observations presented here are on the limit of what can be done without active phase correction systems, at least from a site such as Hat Creek where the BIMA array is situated.   %The molecular line emission to the south-east of the galaxy center is stronger than the emission to the north-west.  The reason for this %is rather that the emission to the north is resolved out by the long BIMA baselines, %whereas it is not by the lower resolution PdBI observations. %-- Gas masses: Schinnerer et al.: %Schinnerer et al. (1998) observe a \\twco(1-0) line flux of 56 mJy~beam$^{-1}$ %from the %innermost  3.''3  which is the emission from the nuclear region. %They find that the $N_{H_2}/I_{CO}$ %conversion factor in I~Zw~1 is close to $2\\times10^{20} $cm$^{-2}$ K$^{-1}$ km$^{-1}$ s %which is also found for molecular gas in our galaxy and and many %nearby external galaxies. %The derived molecular nucleus hydrogen column density in the inner %3\".3 is $4.2\\times10^{22}$ cm$^{-2}$, the molecular gas mass is $7.5\\times10^9 %M_\\sun$.  % This is  2/3 of the total measured flux. The corrsponding  mass is $9 % \\times 10^9 M_\\sun$.  %--The  kinemattic gas masses derived from the rotation curve indicate a %CO/mass ratio which %is in good agreement with Galactic conversion factors. Our data .....  %-- a comparison with the high resolution \\twco(2-1) data of Eckart %et al. show that %the molecular inside the 1 kpc molecular ring  has %higher temperatures/densities than the molecular gas in the ring. This is consistent %with the finding, that  %-- topology: %Our integrated line maps with different synthesized beam weights. %There is strong clumping of the molecular material, i.e. a majority %(... Jy compared to ... Jy in a 2'' beam) are detected at a %resolution of ...) %emission from small scale structures (structures with high spatial %frequencies)  %--Fig. .. shows an overlay of the integrated CO 1-0 molecular line %emission and an optical image obtained with ......  %-- We compare our fluxes for different synthesized beam sizes to the %Schinnerer %  et al. fluxes:  %The corresponding brightness temperature of $T_B = 6$ K indicates a %large beam filling factor and thermalization of a significant amount %of gas within an area corresponding to roughly 500 pc across. %  kinematische Masse: v^2=GM/R  %-- discuss numbers %convert table from power point talk)  %-- simplest radiative transfer? Wolfire model calcs?   %-- compare with other published numbers   %-- A look at channel maps reveals: channel maps are consistent with %orbital %   motion.  %-- velocity diagrams  %-- masked data pos-vel maps  %-- discuss spectrum %We find a double looped spectrum which is consistent with earlier %published %spectra. %A comparison of our observed spectrum with the lower resolution %spectrum from %single dish observations ..."
    },
    "0403/astro-ph0403697_arXiv.txt": {
        "abstract": "Recently, there has been a great deal of interest in understanding the reionization of hydrogen in the intergalactic medium (IGM).  One of the major outstanding questions is how this event proceeds on large scales. Motivated by numerical simulations, we develop a model for the growth of \\ion{H}{2} regions during the reionization era.  We associate ionized regions with large-scale density fluctuations and use the excursion set formalism to model the resulting size distribution.  We then consider ways in which to characterize the morphology of ionized regions.  We show how to construct the power spectrum of fluctuations in the neutral hydrogen field.  The power spectrum contains definite features from the \\ion{H}{2} regions which should be observable with the next generation of low-frequency radio telescopes through surveys of redshifted 21 cm emission from the reionization era.  Finally, we also consider statistical descriptions beyond the power spectrum and show that our model of reionization qualitatively changes the distribution of neutral gas in the IGM. ",
        "introduction": "\\label{intro}  The reionization of the intergalactic medium (IGM) is one of the landmark events in early structure formation.  It marks the epoch at which radiative feedback from luminous objects impacted the farthest reaches of the Universe -- the point at which structure formation affected every baryon in the Universe, albeit indirectly.  The timing, duration, and character of this event contain an enormous amount of information about the first cosmic structures and also have important implications for later generations of baryonic objects.  For these reasons, a great deal of attention -- both observational and theoretical -- has recently been focused on this process.  Most significantly, reionization is an important signpost that connects several disparate observations.  Currently the data provide tantalizing hints about the ionization history of the Universe but few definitive answers.  Observations of Ly$\\alpha$ absorption in the spectra of high-redshift quasars indicate that reionization ends at $z \\sim 6.5$ \\citep{becker,fan,white03,wyithe04-prox}, although this interpretation is controversial \\citep{songaila04}.  The main difficulty with these measurements is that the Ly$\\alpha$ optical depth is extremely large in a fully neutral medium \\citep{gp}, making it difficult to place strong constraints when the neutral fraction exceeds $\\sim 10^{-3}$. On the other hand, measurements of the cosmic microwave background (CMB) imply a high optical depth to electron scattering, apparently requiring reionization to begin at $z \\ga 14$ \\citep{kogut03,spergel03}.  Unfortunately, the CMB data provide only an integral constraint on the ionization history.  Taken together, these observations rule out simple pictures of fast reionization (e.g., \\citealt{barkana01} and references therein), but it is not yet clear what they do imply about early generations of luminous sources \\citep{sokasian03a,sokasian03b,wyithe03,cen03,haiman03,onken03,fukugita03}.  At the same time, our theoretical understanding of how reionization proceeds, given some source population, has been advancing rapidly. Most models of reionization are based on the growth of \\ion{H}{2} regions around individual galaxies \\citep{arons72,barkana01}.  These models use semi-analytic techniques to compute the evolution of global quantities like the total ionized fraction.  However, they are unable to describe the morphology of reionization, an issue that is both theoretically interesting and observationally accessible.  Morphology is a difficult aspect to address because it depends on such factors as the locations of individual sources, the clumpiness of the IGM, the underlying density distribution (the cosmic web), recombinations, and radiative transfer effects.  Some semi-analytic models have been developed to address these issues.  The most popular assumes that reionization is controlled by recombinations and proceeds from low to high density regions \\citep{miralda00}, but these models are approximate at best.  Fortunately, it has recently become possible to incorporate radiative transfer into numerical simulations of the reionization era \\citep{gnedin00,sokasian03a,sokasian03b,ciardi03-sim}, at least as a post-processing step.  Because simulations can include all of the above processes (with the possible exception of clumpiness, which still requires extremely high mass resolution), they give a more nuanced view of the reionization process.  One of the chief lessons of the simulations is that reionization is significantly more inhomogeneous than expected.  The classical picture of a large number of small \\ion{H}{2} regions surrounding individual galaxies does not match the simulations well; instead, a relatively small number of large ionized regions appear around clusters of sources (see, for example, Figure 6 of \\citealt{sokasian03a}).  Moreover, in the simulations reionization proceeds from high to low density regions, implying that recombinations play only a secondary role in determining the morphology of reionization; instead, large-scale bias plays a dominant role.  This picture suggests a new approach to analytic models of reionization, one that takes into account the large-scale density fluctuations that are ultimately responsible for this ionization pattern.  We describe such a model in \\S \\ref{bubbles}.  We derive the size distribution of \\ion{H}{2} regions in a way analogous to the \\citet{press} halo mass function and show that it has the qualitative features seen in the simulations.  Observing the morphology of \\ion{H}{2} regions requires new observational techniques as well as robust ways to characterize the data.  Among the most exciting approaches to study the reionization process are surveys of 21 cm emission from neutral hydrogen at high redshifts \\citep{field58,scott,mmr,zald04}.  The idea behind these observations is to map the fluctuating pattern of emission (or absorption) from neutral hydrogen in the Universe over a range of frequencies.  This yields a measurement of fluctuations due to both the density field and the \\ion{H}{2} regions.  Because 21 cm emission comes from a single spectral line, such observations allow a three-dimensional reconstruction of the evolution of neutral hydrogen in the IGM and, owing to the design of low-frequency radio telescopes, can probe the large cosmological volumes needed to study the IGM.  Despite numerous technical challenges, instruments able to make the necessary observations will be built in the coming years. These include the Primeval Structure Telescope (\\emph{PAST}),\\footnote{ See http://astrophysics.phys.cmu.edu/$\\sim$jbp for details on PAST.} the Low Frequency Array (\\emph{LOFAR}),\\footnote{ See http://www.lofar.org for details on LOFAR.} and the Square Kilometer Array (\\emph{SKA}).\\footnote{ See http://www.skatelescope.org for details on the SKA.}  The major obstacles to high-redshift 21 cm observations are the many bright foreground sources, which include Galactic synchrotron emission, free-free emission from galaxies \\citep{oh03}, faint radio-loud quasars \\citep{dimatteo02}, and synchrotron emission from low-redshift galaxy clusters \\citep{dimatteo04}.  Fortunately, all of these foregrounds have smooth continuum spectra.  Zaldarriaga et al. (2004; hereafter ZFH04) showed that the foregrounds can be removed to high precision because the 21 cm signal itself is uncorrelated over relatively small frequency ranges.  ZFH04 also showed how to compute the angular power spectrum of the 21 cm sky as a function of frequency, given some model of reionization.  This is the simplest statistical measure of the morphology of \\ion{H}{2} regions, and they showed that it is quite powerful in distinguishing different stages of reionization (at least in a simple toy model).  Of course, predictions for the 21 cm signal, or any other measurement of reionization, rely on an accurate model for reionization.  The initial conditions are straightforward: when the neutral fraction $\\bxh \\approx 1$, the power spectrum follows that of the density \\citep{tozzi}.  This breaks down when the \\ion{H}{2} regions appear, and more sophisticated approaches become necessary.  There have been several recent attempts to predict the signal using numerical simulations \\citep{ciardi03,furl-21cmsim,gnedin03}.  However, the boxes in these simulations have sizes $\\la 20 \\Mpc$, subtending only a few arcminutes on the sky.  Because the angular resolution of low frequency radio telescopes at the required sensitivity is also a few arcminutes (ZFH04), such predictions require extrapolations to larger scales, which may be dangerous given the already large sizes of ionized regions.  With our analytic model for the size distribution of these bubbles, we are able to make the first detailed statistical characterizations of the fluctuating ionization pattern on the large scales most relevant to observations.  We begin with the power spectrum as the simplest description.  In \\S \\ref{ps}, we show how to compute the power spectrum for an arbitrary size distribution, and in \\S \\ref{res} we apply that formalism to our model of \\ion{H}{2} regions.  The power spectrum is only one way to describe the 21 cm field.  It is an excellent approximation if fluctuations in the matter density dominate the signal, because the density field is neary gaussian on the large scales of interest here.  Unfortunately, once the \\ion{H}{2} regions dominate the power spectrum, the probability distribution is no longer gaussian \\citep{morales03}.  There have, however, been no attempts to describe the distribution.  Using our model for reionization, we discuss some of the relevant features in \\S \\ref{nongauss}.  We show that the reionization model qualitatively changes the character of fluctuations in the neutral gas density.  This paper is primarily concerned with developing a general and useful model for the morphology of reionization and with the crucial features of that model.  We will focus here on the physics of reionization rather than on their observable consequences; we consider some of these in a companion paper about the 21 cm signal expected from high redshifts \\citep{furl04b}.  In our numerical calculations, we assume a cosmology with $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, $\\Omega_b=0.046$, $H=100 h \\hunits$ (with $h=0.7$), $n=1$, and $\\sigma_8=0.9$, consistent with the most recent measurements \\citep{spergel03}. ",
        "conclusions": "\\label{disc}  In this paper, we have described a new analytic model for the size distribution of \\ion{H}{2} regions during reionization.  While the overlap process is usually described in terms of Str{\\\" o}mgren spheres around individual galaxies, recent cosmological simulations of reionization have demonstrated that the ionized regions are much larger than naively expected, even early in the reionization process.  Here we have shown that the size distribution can be understood in terms of large-scale features in the density field, and we have constructed the size distribution with an approach analogous to the standard derivation of the \\citet{press} mass function for collapsed halos. The model has only two input parameters (if the cosmology is fixed): the ionizing efficiency of collapsed objects and the minimum mass halo that can host a luminous object.  Interestingly, at a fixed neutral fraction the characteristic size of the bubbles is fairly insensitive to these input parameters, suggesting that the morphology of reionization is close to invariant (at least for simple, single reionization episodes).  While we make several simplifying approximations in the model (see the discussion in \\S \\ref{bubbles}), it provides a self-consistent approach that reproduces the qualitative features of simulations.  In the future, the model must be quantitatively compared to simulations; however, an accurate comparison requires simulations with large ($\\sim 100^3 \\Mpc^3$) volumes because of the effects of the large-scale density field (see Figure \\ref{fig:dndr-comp}, as well as \\citealt{barkana03}).  The morphology of reionization has implications for a number of observables.  Most important, regardless of the technique, we predict large \\ion{H}{2} regions that should be feasible to detect, whether through quasar absorption spectra \\citep{miralda00,barkana02}, Ly$\\alpha$ lines at extremely high-redshifts \\citep{pello04,loeb04,ricotti04,gnedin04,cen04,bart04}, or 21 cm tomography, For example, the noise in a 21 cm map is proportional to the square root of the bandwidth and, more important, to the square of the pixel size.  Our model predicts substantial contrast on relatively large scales of several arcminutes, which should make detections easier.  Note, however, that in many next-generation experiments like LOFAR the ``bandwidth\" can be chosen \\emph{after} the observations are complete, so the expected scale of the bubbles need not determine the experimental design \\citep{morales03}.  Even if high signal-to-noise detections of individual \\ion{H}{2} regions are not available, we have shown that statistical measurements of the size distribution can still strongly constrain reionization. We first constructed the power spectrum of fluctuations in the neutral density, including both density fluctuations and the ionized regions. We found that the ionized regions imprint clear features on the power spectrum and amplify the power by a factor of several during the middle and late stages of reionization.  Most important, the power spectrum evolves throughout reionization, allowing us to map the time history of reionization.  For 21 cm observations, the large-scale \\ion{H}{2} regions predicted by our model put the features at $l \\la 10^4$ (or $\\theta \\ga 2\\arcmin$). This matches well to the scales able to be probed by upcoming experiments like PAST, LOFAR, and SKA (ZFH04).  If, on the other hand, reionization occurred through the overlap of \\ion{H}{2} regions around individual galaxies, the features in the power spectrum would appear at much smaller scales, perhaps beyond the reach of these instruments.  We have also shown explicitly that the ionized bubbles induce qualitative changes to the initially gaussian neutral density distribution.  This suggests that statistical measurements beyond the power spectrum can offer probes of the physics of reionization that may be even more powerful.  Our results also have important implications for other measurements. For example, one of the most successful strategies for targeting high-redshift galaxies is by searching for strong Ly$\\alpha$ emitters. If the galaxy is embedded in a mostly neutral medium, Ly$\\alpha$ absorption from the IGM can have a substantial effect on the line profile \\citep{haiman02,santosm03,loeb04}.  In our model, we expect this absorption to be less significant, because most galaxies are embedded in \\ion{H}{2} regions with large sizes, even relatively early in the reionization process.  Finally, in this paper we have considered only the simplest reionization histories with a single type of source.  Many models for reconciling the quasar and CMB data on reionization require multiple, distinct generations of sources that cause ``stalling\" or even ``double\" reionization \\citep{wyithe03,cen03,haiman03,sokasian03b}.  Such histories will of course change the morphology of reionization and hence modify the 21 cm signal.  Moreover, alternative models of reionization in which (for example) voids are ionized first yield different sets of signatures.  In \\citet{furl04b}, we use the formalism developed here to examine how well 21 cm tomography can distinguish these histories and models."
    },
    "0403/astro-ph0403374_arXiv.txt": {
        "abstract": "The mergers of black hole-neutron star binaries are calcuated using a pseudo-general relativistic potential that incorporates ${\\mathcal O}(v^2/c^2)^3$ post-Newtonian corrections.  Both normal matter neutron stars and self-bound strange quark matter stars are considered as black hole partners.  As long as the neutron stars are not too massive relative to the black hole mass, orbital decay terminates in stable mass transfer rather than an actual merger.  For a normal neutron star, mass transfer results in a widening of the orbit but the stable transfer ends before the minimum neutron star mass is reached.  For a strange star, mass transfer does not result in an appreciable enlargement of the orbital separation, and the stable transfer continues until the strange star essentially disappears.  These differences might be observable through their respective gravitational wave signatures. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403468_arXiv.txt": {
        "abstract": "Although thermal disk emission is suppressed or absent in the hard state of X-ray binaries, the presence of a cold, thin disk can be inferred from signatures of reprocessing in the $\\sim2-50$ keV band. The strength of this signature is dependent on the source spectrum and flux impinging on the disk surface, and is thus very sensitive to the system geometry.  The general weakness of this feature in the hard state has been attributed to either a truncation of the thin disk, large ionization, or beaming of the corona region away from the disk with $\\beta\\sim0.3$.  This latter velocity is comparable to jet nozzle velocities, so we explore whether a jet can account for the observed reflection fractions.  It has been suggested that jets may contribute to the high-energy spectra of X-ray binaries, via either synchrotron from around $100-1000$ $r_{\\rm g}$ along the jet axis or from inverse Compton (synchrotron self-Compton and/or external Compton) from near the base.  Here we calculate the reflection fraction from jet models wherein either synchrotron or Compton processes dominate the emission.  Using as a guide a data set for GX~339$-$4, where the reflection fraction previously has been estimated as $\\sim10\\%$, we study the results for a jet model. We find that the synchrotron case gives $< 2\\%$ reflection, while a model with predominantly synchrotron self-Compton in the base gives $\\sim 10-18\\%$.  This shows for the first time that an X-ray binary jet is capable of significant reflection fractions, and that extreme values of the reflection may be used as a way of discerning the dominant contributions to the X-ray spectrum. ",
        "introduction": "X-ray binaries (XRBs) have been observed in several distinct states, which are characterized by the relative strength of their soft and hard X-ray emission components, as well as by their variability properties \\citep[see, e.g.,][]{McClintockRemillard2003}.  In the ``standard'' models \\citep[see][and references therein]{ReynoldsNowak2003} the soft component is well-explained with thermal emission from a standard thin disk \\citep{ShakuraSunyaev1973}, while the hard power-law component is generally attributed to inverse Compton (IC) scattering processes.  The various models currently in existence often have quite different seed photons and system geometries, yet predict similar results for the broad continuum emission \\citep[see][]{NowakWilmsDove2002,Markoffetal2003}.  In order to discern between the models, therefore, one has to look at finer details which are dependent upon specific elements of the geometry.  In the hard state, the power-law component dominates the thermal disk emission over most of the X-ray range.  The presence of a cold, thin disk can still be inferred via detection of a soft component in the $\\approx$0.3--1\\,keV band, and via spectral components in the $\\approx$2--50\\,keV band  suggesting that a fraction of the hard X-rays is reprocessed or reflected from an optically thick surface. The reflection component is characterized by a flattening of the power-law above $\\sim10$\\,keV \\citep[e.g.,][]{Poundsetal1990,GeorgeFabian1991}, as well as by spectral features such as an Fe K$\\alpha$ fluorescent line and an Fe edge \\citep[see][for a review]{ReynoldsNowak2003}.  The strength of these components is directly related to the spectrum and flux hitting the disk, and is therefore sensitive to assumptions about the system geometry.  For example, in models where the hard X-rays are due to IC in a hot coronal plasma completely ``sandwiching'' the disk, the reflection is easily too high for typical X-ray binary spectra (and the self-consistently derived coronal temperatures are too low; e.g., \\citealt{stern:95a,dove:97a}).  Thus modifications have been proposed, such as a recessed thin disk \\citep{dove:97b,poutanen:97b}, beaming away from the disk \\citep{reynolds:97c,Beloborodov1999}, or large amounts of ionization of the disk \\citep{ross:99a,nayakshin:01a,done:01c}.  In addition to the standard corona models, \\citet{MarkoffFalckeFender2001} proposed that the entire broadband spectrum of hard state XRBs could instead result from synchrotron emission at the beginning of an acceleration region.  While controversial, this model succeeds at explaining the tight correlation of radio and X-ray emission seen in several sources \\cite[e.g., GX~339$-$4,][]{Corbeletal2000,Corbeletal2003,Markoffetal2003}, which, in fact, may be a universal correlation in XRB hard states \\citep{GalloFenderPooley2003}.  It is also the first model to provide a link between the inferred presence of a hot, magnetized electron plasma near the inner regions of the central engine to the hot, magnetized plasma that we know exists via imaged radio jets \\citep[e.g.,][]{Stirlingetal2001}.  A shortcoming of these models, however, is that they have not yet attempted to explain detailed X-ray spectral features such as reflection and iron lines.  In this Letter we calculate the expected reflection fraction from several jet models with parameters that provide good descriptions of broad band features (e.g., overall luminosity, flat radio spectrum, X-ray spectral slope and cutoff) for `typical' XRB data sets, such as those shown here for the Galactic source GX~339$-$4. We then discuss the ensuing constraints on jet models. ",
        "conclusions": "As a first step towards judging the magnitude of the reflection fraction, we can calculate $R(\\mu)$ as a function of energy by substituting $dP/d\\Omega d\\nu$ into eq.~\\ref{eq:def}.  The synchrotron-dominated cases show the smallest overall reflection.  The ratio is only $\\approx$1--2\\% in between energies of 1--100\\,keV ($\\beta_{\\rm s}=0.3, ~0.4$).  The ratio peaks at 0.3\\,keV with value 6\\% and 8\\% for $\\beta_{\\rm s}=0.3$ and $\\beta_{\\rm s} = 0.4$, respectively.  The SSC-dominated jet shows significantly more reflection.  For $\\beta_{\\rm s}=0.3$ and $\\beta_{\\rm s}=0.4$, the ratio is $>10\\%$ in between 0.5--23\\,keV.  The former peaks at 18\\% at 6\\,keV, while the latter peaks at 17\\% at 5\\,keV.  Such values of `reflection fraction' are comparable to the observed range for GX~339$-$4, and in fact are larger than fitted for this particular observation \\citep{NowakWilmsDove2002}.  \\clearpage  \\begin{figure*} \\epsscale{1} \\plottwo{f4a.eps}{f4b.eps} \\caption{Total model spectrum (direct plus reflected - solid line), summed spectrum incident upon the disk (dashed line), and reflected spectrum (lower solid line).  {\\it Left:} SSC-dominated jet with $\\beta_{\\rm s} = 0.4$.  {\\it Right:} Synchrotron-dominated jet with $\\beta_{\\rm s} = 0.4$. \\label{fig:ref}} \\end{figure*}  \\clearpage  As described above, we have also directly calculated the reflection spectrum expected from our models.  We show spectra for the SSC-dominated jet and the synchrotron-dominated jet, both with $\\beta_{\\rm s} = 0.4$, in Fig.~\\ref{fig:ref}.  Qualitative differences are immediately apparent.  The SSC-dominated jet clearly has an overall brighter reflection spectrum, and effects of relativistic smearing are readily visible.  These effects are to be expected given both the lower velocities of the SSC-dominated region ($\\beta \\sim 0.3$) as well as the fact that this region is closer to the central black hole.  We have chosen a slightly face-on orientation, $\\mu=0.77$, since there has been some suggestion that this is indeed the case for the GX~339$-$4 system \\citep{Wuetal2001}.  Most BHC systems should have a higher inclination, which, given eq.~\\ref{eq:def}, allows for even greater reflection fractions.  If we freeze all of our model parameters but instead choose $\\mu=0.5$, reflection fractions, as defined by the ratio of the disk-incident spectrum to the directly viewed spectrum, increase.  Over the 1--20\\,keV interval, this ratio is everywhere $> 26\\%$ for the SSC-dominated jet with $\\beta_{\\rm s} = 0.3$, and $>7\\%$ for the synchrotron-dominated jet with $\\beta_{\\rm s} = 0.4$. The former peaks at 38\\% at 2\\,keV and, again, exhibits clear relativistic broadening of the reflection features.  While still preliminary, this is the first time that the reflection of jet emission off an accretion disk has been calculated for an X-ray binary system.  These results show that jets are indeed capable of producing the reflection fraction inferred from the X-ray data, providing further support for a connection between the base of the jet and the corona.  We have highlighted two extreme cases which might be applicable to different physical situations.  Systems that exhibit very low reflection ($\\lesssim 5\\%$), with sharp (i.e., non-relativistic) features in any reflected spectrum \\citep[e.g., XTE J1118+480,][]{Milleretal2001}, could be synchrotron-dominated, and clearly rule out SSC-dominated jets.  Systems with significantly larger reflection fractions ($\\gtrsim 15\\%$) cannot be synchrotron-dominated, especially if they exhibit features which are unambiguously relativistically smeared.  However even in the Compton-dominated regime, as shown in Fig.~\\ref{gx339}, synchrotron radiation can contribute $\\gtrsim 10\\%$ of the flux which, as we discuss further below, will greatly effect fits to data with corona plus reflection models.  For intermediate values, or values $\\gtrsim 30\\%$, other factors need to be considered.  For instance, these results assume that the disk is perfectly flat, and that the jet is always perpendicular to the disk. Realistically, disks are expected to flared or warped \\citep[e.g.,][]{Dubusetal1999}, and several systems also show evidence for misalignment between the jets and outer disks \\citep{Maccarone2002}.  Both of these effects will serve to increase the reflection fraction from the jet, particularly for the synchrotron component.  Therefore, we treat these numbers as lower limits.  We would like to emphasize, however, that a significant X-ray contribution from jet synchrotron emission can greatly alter how one even defines `reflection fraction' based upon a presumed single-component, underlying continuum.  As shown in Fig.~\\ref{cygx1}, jet radiation alone can provide a good description of the high energy cutoff region.  One can readily imagine a model wherein the soft X-ray region is dominated by SSC emission (as in Fig~\\ref{fig:ref}b) and/or Comptonization of external (disk) photons, each with a large covering factor of the disk (i.e., essentially unity reflection fraction, for that component alone).  The hard X-ray radiation could then be dominated by synchrotron radiation with inherently low reflection fraction.  The net spectrum would have an intermediate fitted value of `reflection fraction' that does not have a `geometric interpretation' entirely appropriate for either the soft or hard emission components. If the broad-band X-ray continua of hard state BHC are in fact comprised of such multiple components, as in some of the jet models presented here, then this calls into question current interpretations of `reflection fractions' based upon single component fits.  Of course, in order to determine whether such multiple spectral components are indeed present in the observations, actual fitting of the combined direct plus reflection spectrum needs to be performed. The calculations and models presented here provide vital clues as to how much each process can contribute for this next step. This work has shown that this type of analysis may hold the key to disentangling the emission processes relevant from the accretion inflow and outflow, and place limits on the synchrotron vs. Compton contributions to the hard state spectrum."
    },
    "0403/astro-ph0403142_arXiv.txt": {
        "abstract": "{ We present the initial results from our \\xmm\\ program aimed at searching for X-ray activity cycles in solar-type stars. \\hd\\ is a G2-type star (somewhat more evolved than the Sun, and with a less massive companion) with a pronounced 8.2 yr chromospheric cycle, as evident from from the Mt.\\ Wilson program data. We present here the results from the initial 2.5 years of \\xmm\\ observations, showing that large amplitude (a factor of $\\simeq 10$) modulation is present in the X-ray luminosity, with a clearly defined maximum in mid 2002 and a steady decrease since then. The maximum of the chromospheric cycle took place in 2001; if the observed X-ray variability is the initial part of an X-ray cycle, this could imply a phase shift between chromospheric and coronal activity, although the current descent into chromospheric cycle minimum is well reflected into the star's X-ray luminosity. The observations presented here provide clear evidence for the presence of large amplitude X-ray variability coherent with the activity cycle in the chromosphere in a star other than the Sun. ",
        "introduction": "\\label{sec:intro}  The 11 year cycle is perhaps the oldest known manifestation of the Sun's magnetic activity, having first been noted by Schwabe in 1843 as a periodic modulation of the number of sunspots. Subsequently, most activity indicators have been observed to follow a similar cyclical variation, with the amplitude of the modulation dependent on the indicator used. On cool stars other than the Sun, the detection of cycles had to wait for the foresight of O. Wilson, who started a long-term monitoring of a Ca\\,{\\sc ii} H\\&K activity indicator (the ``$S$ index'') in a significant number of stars, using the Mt. Wilson 100 inch telescope. An analysis of the vast amount of data from the Mt. Wilson program, now covering nearly 40 years of observations (\\citealp{bds+95}), shows that solar-like cycles are present in many stars, although some stars show no detected variability (perhaps being in a ``Maunder-like'' state) while others show significant non-periodic variability. In the Sun, the amplitude of the cyclical S index modulation is a factor of about 2, while the amplitude in X-rays is much stronger, i.e.\\ a factor of 100 in the Yohkoh 0.73--2.5 keV band.  Yet evidence of cyclical variability in X-rays in stars other than the Sun has only become available very recently. In fact (with the recent exception of 61~Cyg), the X-ray observations of the few stars for which sufficient data exist suggest that their X-ray luminosity is relatively stable over long-term intervals, as discussed e.g. by \\cite*{ste98b}. Observations of homogeneous samples of active stars, e.g.\\ of the Hyades, show little variation in \\lx\\ across the $\\simeq 10$-year separation between the \\emph{Einstein} and ROSAT PSPC surveys, with 90\\% of the stars showing less than a factor of 2 variability. In these stellar samples, however, the median activity level is much higher than the Sun's, by 2 or more dex: stars at this activity level typically do not show cycles in Ca\\,{\\sc ii}, instead varying irregularly, so that perhaps the lack of X-ray modulation is not surprising.  This result was however confirmed by \\cite*{ls2004} on a volume-limited sample of solar-type stars using ROSAT All-Sky Survey (RASS) and pointed observations. Weak statistical evidence for the presence of solar-like cycles in less active stars was derived by \\cite*{hss96} using the RASS observations of the stars in the Mount Wilson program, looking at the deviations from a ``mean'' X-ray luminosity for each star as a function of the Ca\\,{\\sc ii} cycle phase.  In their analysis of X-ray variability properties of solar mass stars \\cite*{mmp+2002} found that, in relatively quiet stars, amplitude variations increase with time scales, and interpreted this as an indication of the presence of solar-like cycles in stars with X-ray activity of the same order of that of the Sun. The comparison between the Sun and nearby stars is consistent with a fraction of moderately active stars ($\\overline{L}_{\\rm X} < 10^{28}$ erg/sec) having X-ray variability similar to the Sun, while more active stars lack solar-like cyclic coronal activity (\\citealp{mm2003}).  More recently, \\cite{hsb+2003} presented evidence of X-ray luminosity modulation in 61 Cyg A and B (K5V and K7V) over 4.5 years, well correlated with their $S$ index, strongly suggestive of an activity cycle in X-rays. The observed modulation amplitude (using ROSAT HRI data, which prevented a study of the spectral evolution) is a factor of 2.5.  In this paper we present the results of the first 2.5 yr of the \\hd\\ \\xmm\\ monitoring program, showing for the first time in a star other than the Sun strong evidence for long-term coherent X-ray variability with an amplitude a factor of about 10, consistent with the cyclic amplitude modulation in the Sun. ",
        "conclusions": "\\label{sec:concl}  We have presented clear evidence of long-term variations of the X-ray luminosity in \\hd, a solar-type star with a well defined cycle in its chromospheric activity. The variations thus far determined have an amplitude of a factor of 10 (over the 2.5 years covered by the observations), comparable to the variations seen in the X-ray luminosity of the Sun during the solar cycle. These variations suggest the beginning of a cycle; while the observed X-ray maximum appears somewhat offset (by about 1 year) from the chromospheric one, the current descending phase of the chromospheric cycle is well reflected in \\hd's decreasing X-ray luminosity over the last two years.  \\hd\\ is the subject of a long-term monitoring program performed with the \\xmm\\ observatory (which also includes $\\alpha$ Cen and 61~Cyg), of which we are presenting here the first (still perforce preliminary) results. A further two years of observations (at six months cadence) are already planned on the same target, and we will re-propose the target for every AO, to ensure the continuous monitoring. While the nature of this program is such that a few more years will be necessary before a detailed analysis can be performed, the initial results presented here show that our initial choice of targets was a good one; we now can state with some confidence that the Sun is not the only solar-type star for which large amplitude X-ray variability coherent with the activity cycle in the chromosphere is present.  The continuation of our monitoring program on \\hd\\ will allow to compare the characteristics of its coronal cycles with the solar one."
    },
    "0403/astro-ph0403232_arXiv.txt": {
        "abstract": "Long-slit stellar kinematic observations were obtained along the major-axis of $30$ edge-on spiral galaxies, $24$ with a boxy or peanut-shaped (B/PS) bulge and $6$ with other bulge types for comparison. B/PS bulges are identified in at least $45\\%$ of highly inclined systems and a growing body of theoretical and observational work suggests that they are the edge-on projection of thickened bars. Profiles of the mean stellar velocity $V$, the velocity dispersion $\\sigma$, as well as the asymmetric ($h_3$) and symmetric ($h_4$) deviations from a pure Gaussian are presented for all objects. Comparing those profiles with stellar kinematic bar diagnostics developed from N-body simulations, we find bar signatures in $24$ of our sample galaxies ($80\\%$). Galaxies with a B/PS bulge typically show a double-hump rotation curve with an intermediate dip or plateau. They also frequently show a rather flat central velocity dispersion profile accompanied by a secondary peak or plateau, and numerous galaxies have a local central $\\sigma$ minimum ($\\gtrsim40\\%$). The $h_3$ profiles display up to three slope reversals. Most importantly, $h_3$ is normally correlated with $V$ over the presumed bar length, contrary to expectations from axisymmetric disks. These characteristic bar signatures strengthen the case for a close relationship between B/PS bulges and bars and leave little room for other explanations of the bulges' shape. We also find that $h_3$ is anti-correlated with $V$ in the very center of most galaxies ($\\gtrsim60\\%$), indicating that these objects additionally harbor cold and dense decoupled (quasi-)axisymmetric central stellar disks, which may be related to the central light peaks. These central disks coincide with previously identified star-forming ionized-gas disks (nuclear spirals) in gas-rich systems, and we argue that they formed out of gas accumulated by the bar at its center through inflow. As suggested by N-body models, the asymmetry of the velocity profile ($h_3$) appears to be a reliable tracer of asymmetries in disks, allowing us to discriminate between axisymmetric and barred disks seen in projection. B/PS bulges (and thus a large fraction of all bulges) appear to be made-up mostly of disk material which has acquired a large vertical extent through bar-driven vertical instabilities. Their formation is thus probably dominated by secular evolution processes rather than merging. ",
        "introduction": "} The number of edge-on disk galaxies with a boxy or peanut-shaped (B/PS) bulge is known to be significant. In fact, the fraction of these objects has continuously increased as sample selection, data, and classification criteria improved \\citep*[see][]{j86,sa87,s87,db88,ldp00a,ldp00b}. Recently, \\citet{ldp00a} classified $\\approx1350$ edge-on disk galaxies and found $45\\%$ of B/PS bulges, with a fraction higher than $40\\%$ for all morphological types from S0 to Sd. This was the first study to report such a large incidence of B/PS bulges in very late type systems, but it is also the largest and most homogeneous. B/PS bulges are therefore common, but they are only seen in highly inclined systems ($i\\gtrsim75^{\\circ}$; e.g.\\ \\citealt*{sdb90}), indicating that their shape is mainly related to the vertical distribution of light.  Several scenarios have been suggested to explain the structure and formation of B/PS bulges. They can be divided into two main categories: axisymmetric accretion and non-axisymmetric bar buckling models. \\citet{bp85} and \\citeauthor{r88} (\\citeyear{r88}; but see also \\citealt{man85}) showed that it is possible to form axisymmetric cylindrically rotating B/PS bulges from the accretion of one or (more likely) many satellite galaxies onto a larger galaxy host. As pointed out by \\citet{bp85}, however, the particular alignment required puts too many constraints on the orbits, angular momenta, and relative masses of the galaxies, and this process appears unlikely in view of the high incidence of B/PS bulges. Moreover, satellite galaxies are not preferentially observed around galaxies with a B/PS bulge \\citep{s87}. Using N-body simulations, \\citet{cs81} first suggested that B/PS bulges could be due to the thickening of bars in the disks of spiral galaxies. They showed that soon after a bar develops through the usual bar instability, it buckles, thickens, and settles with a larger velocity dispersion and thickness. When highly inclined, the bar then appears peanut-shaped when seen side-on (i.e.\\ perpendicular to the bar), boxy-shaped when seen at intermediate angles, and almost round when seen end-on (i.e.\\ parallel to the bar). It is also largely cylindrically rotating. \\citeauthor{cs81}'s (\\citeyear{cs81}) results were later confirmed and improved upon by many studies \\citep[e.g.][]{cdfp90,rsjk91,am02} and they are now universally accepted. The three-dimensional (3D) orbital structure of these bars has just been investigated in a series of papers (\\citealt*{spa02a,spa02b,psa02,psa03}; but see also \\citealt{p84,pf91}), allowing a more detailed comparison of models with high spatial resolution images.  There are several observational studies supporting the connection between B/PS bulges and bars. Most important for this work, \\citet{km95} suggested that bars could be detected in edge-on galaxies from their (projected) kinematic signature. In their pioneering study of two spirals with a peanut-shaped bulge (\\objectname{NGC~5746} and \\objectname{NGC~5965}), \\citet{km95} found a characteristic ``figure-of-eight'' bar signature in their position-velocity diagrams (PVDs; the projected density of material as a function of line-of-sight velocity and projected position). Although the interpretation of this signature was ultimately flawed, hydrodynamical simulations by \\citet{ab99} confirmed its relationship to bars and showed how it can be used to identify and characterize edge-on bars from gas kinematics. In a study of emission-line spectra for $10$ edge-on spirals with various types of bulges, \\citet{mk99} later confirmed the link between B/PS bulges and bars. \\citeauthor{bf99} (\\citeyear{bf99}; see also \\citealt{b98}) also presented ionized-gas PVDs for $23$ edge-on galaxies, $17$ with a B/PS bulge and $6$ with other types of bulges (mostly spheroidal). Using \\citeauthor{ab99}'s (\\citeyear{ab99}) diagnostics, they found bar signatures in $14$ of $17$ galaxies with a B/PS bulge ($\\approx80\\%$) and in none of the galaxies with a more spheroidal bulge.  In spite of the aforementioned studies, our knowledge of the stellar kinematics of B/PS bulges and their hosts remains limited. This probably arises because the gas is a more practical tracer of the barred potential, being both easier to observe and reacting more strongly than the stars to asymmetries (shocks). Studying the stars directly is nevertheless essential as B/PS bulges are stellar structures and, when comparing to N-body simulations, stellar kinematics largely bypasses issues such as the selection of an optimal gaseous tracer (H$\\alpha$, CO, \\ion{H}{1}, etc), star formation, and (for gas-poor objects) the response of the orbital structure to a massive gas disk \\citep[e.g.][]{bhsf98}.  Existing stellar kinematic studies of B/PS bulges are limited to a few objects and have largely focused on the issue of cylindrical rotation, using slits offset from the major-axis or perpendicular to it \\citep*[e.g.][]{bc77,ki82,r86,j87,s93a,swc93}. As noted above, however, this is consistent with both axisymmetric and barred configurations, and so does not discriminate between the two dominant formation scenarios. We advocate here instead the new kinematic bar diagnostics of \\citet{ba04}, which directly probe the shape of the potential \\citep[but see also][]{bg94,km95}.  The importance of B/PS bulges goes beyond their sheer number and interesting morphology, as they may hold vital clues for our understanding of bulge formation. There is growing evidence against significant merger growth in many bulges (e.g.\\ \\citealt*{k93,apb95,mch03,bgdp03}) and a corresponding emphasis on secular evolution processes, largely bar-driven \\citep[e.g.][]{wfmmb95,es02}. Most theoretical models involve the growth of a central mass through bar inflow and/or possibly recurring bar destruction \\citep*[e.g.][]{pn90,fb93,fb95,nsh96}, although the efficiency of bar dissolution mechanisms remains uncertain (\\citealt{ss03} and references therein). Nevertheless, because of their probable relationship to bars, B/PS bulges are likely to play a central role in those scenarios, and a proper understanding of their structure and dynamics is essential to constrain them.  Our primary goal in this work is thus to study the stellar kinematics of a statistically significant number of galaxies with a B/PS bulge, simultaneously and independently verifying their relationship to bars and probing for embedded structures. We study \\citeauthor{bf99}'s (\\citeyear{bf99}) sample of $23$ edge-on spirals, $17$ with a B/PS bulge and $6$ with other bulge types, as well as an additional $7$ B/PS bulges which showed little (or confined) emission and were only briefly discussed by the authors. The latter offer the best comparison with N-body models, being largely gas and dust free.  In \\S~\\ref{sec:diagnostics}, the stellar kinematic bar diagnostics of \\citet{ba99} and \\citet{ba04} are summarized. In \\S~\\ref{sec:obs}, we describe the sample, observations, as well as the reduction and analysis of the data. The stellar kinematics of the $30$ galaxies is presented briefly in \\S~\\ref{sec:results}, while \\S~\\ref{sec:discussion} discusses in greater detail the structure and dynamics of B/PS bulges and how they fit into bar-driven secular evolution scenarios. We summarize our results and conclude briefly in \\S~\\ref{sec:conclusion}. ",
        "conclusions": "} The (projected) morphology of B/PS bulges departs strongly from that of classical, spheroidal bulges \\citep[e.g.][]{s93b}, implying an even more distorted 3D structure, but their high incidence ($\\gtrsim45\\%$ of all bulges; see \\citealt{ldp00a}) makes them more than simple curiosities and a widespread mechanism must be able to account for their structure and dynamics. By studying the stellar kinematics of the $30$ edge-on spiral galaxies of \\citet{bf99}, we have attempted here to discriminate between the two leading formations scenarios: accretion \\citep[e.g.][]{bp85,r88} and bar-buckling \\citep[e.g.][]{cs81,cdfp90}. The latter is particularly attractive in view of currently popular secular evolution scenarios, which argue that bars may drive bulges along the Hubble sequence (e.g.\\ \\citealt{fb95}; \\citealt{nsh96} and references therein).  Our results are summarized in Table~\\ref{tab:class}. $22$ of $24$ galaxies with a B/PS bulge show clear kinematic bar signatures ($\\approx90\\%$), while only $2$ of the $6$ control sample galaxies do and are probably transition objects. We thus confirm the results of \\citet{bf99} for galaxies with extended ionized-gas and, for the first time, extend those results to gas-poor B/PS bulges (S0s). The kinematic evidence in favor of a close relationship between thickened bars and B/PS bulges is thus now overwhelming across all morphological types (S0--Sbc), although we still have not probed the potentials out of the equatorial plane. Zamojski et al.\\ (2004, in preparation) shall discuss the stellar kinematics of a few objects at large galactic heights.  Most objects show kinematic profiles in agreement with those predicted from N-body simulations of barred disks \\citep{ba04}, including the expected relationship between bulge morphology and bar viewing angle. In particular, most galaxies with a B/PS bulge have i) a double-hump rotation curve with a dip or plateau at moderate radii, ii) a rather flat central velocity dispersion profile with a secondary peak or plateau and (in about $40\\%$ of cases) a local central minimum, and iii) an $h_3$ profile correlated with $V$ over the expected bar length, with up to three slope reversals. As expected if caused by a bar, the strengths and extents of those features are correlated. Our $h_4$ profiles are generally too noisy to be of much use.  A number of differences with the simulations also exist, most of which appear related to the past or current presence of gas. In particular, the kinematic profiles of most galaxies show an anti-correlation of $V$ and $h_3$ in the very center, over a region corresponding to the rapidly rising part of the rotation curve (or at most the first plateau) and encompassing the central light peak and central $\\sigma$ minimum (if present). Being cospatial with observed ionized-gas disks in gas-rich galaxies \\citep{bf99}, this kinematic behavior is most likely due to cold quasi-axisymmetric central stellar disks, which we argue have decoupled due to bar-driven gas inflow, and may have been enhanced by subsequent star formation. B/PS bulges are thus entirely consistent with the predictions of most bar-driven secular evolution scenarios \\citep[e.g.][]{hs94,fb95}, whereby bars can simultaneously account for the creation of a large vertically extended component out of disk material and for the accumulation of mass in their centers (possibly associated with the so-called bulges of face-on systems), bypassing the need for merging or accretion of external material.  Lastly, we reiterate that $h_3$ is a very useful tracer of the structure and dynamics of galaxies. As suggested by \\citet{ba04}, the correlation or anti-correlation of $h_3$ and $V$ appears to be a reliable diagnostics of the intrinsic shape, at least for highly inclined disks."
    },
    "0403/astro-ph0403281_arXiv.txt": {
        "abstract": "We confirm the orbital period of WX~Cen $\\equiv$ WR~48c determined by \\citet{dia95} -- DS95 -- and refined its value to $P_{orb} = 0.416~961~5 (\\pm 22)$ d. The light curve of this object has a peak to peak variation of approximately 0.32 magnitudes. It is non-sinusoidal in the sense that it has a V-shaped narrow minimum, similar to the ones seen in V~Sge, V617~Sgr and in Compact Binary Supersoft Sources -- CBSS.  Most of the emission lines in the optical spectrum are due to Balmer, He II, C IV, N V, O V and O VI. An analysis of the He II Pickering series decrement shows that the system has significant amount of hydrogen. The emission lines of He II 4686{\\AA} became weaker between the 1991 and 2000/2002 observations, indicating distinct levels of activity. The spectra of WX~Cen show variable absorption features in the Balmer lines with $V = -2900$ km~s$^{-1}$ and in emission with $V = \\pm 3500$ km~s$^{-1}$. These highly variable events remind the satellites in emission of CBSS.  We estimate the color excess as $E(B-V)=0.63$  on the basis of the observed diffuse interstellar band at 5780{\\AA}. Given the distance-color excess relation in the direction of WX Cen, this implies a distance of $2.8  \\pm 0.3$ kpc. Interstellar absorption of the Na I D lines show components at $-4.1$ km~s$^{-1}$, which corresponds to the velocity of the Coalsack, and three other components a $-23.9$, $-32.0$ and $-39.0$ km~s$^{-1}$. These components are also seen with similar strengths in field stars that have distances between 1.8 and 2.7 kpc. The intrinsic color of WX Cen is $(B-V)_0=-0.2$ and the absolute magnitude, $M_V = -0.5$.  Extended red wings in the strong emission lines are seen. A possible explanation is that  the system has a spill-over stream similar to what is seen in V617~Sgr. We predict that when observed in opposite phase, blue wings would be observed. A puzzling feature that remains to be explained is the highly variable red wing ($V \\sim 700$ km~s$^{-1}$) of the O VI emission lines as well as of the red wings of the H and He lines.  The velocity of the satellite-like feature is consistent with the idea that the central star is a white dwarf with a mass of $M \\sim 0.9 M_{\\sun}$. With the high accretion rate under consideration, the star may become a  SN~Ia in a time-scale of $5 \\times 10^{6}$ years. ",
        "introduction": "WX~Cen was initially identified by \\citet{egg} as a possible optical counterpart of the hard X-ray transient source Cen XR-2, although this identification was latter discarded. Because of its spectral characteristics, the object was then classified as a Wolf-Rayet star of type WN 7 \\citep{huc81}, while \\citet{vogt} classified it as a nova-like variable. \\citet{dia95} -- DS95 -- showed that it is a binary system with an orbital period of 10.0 hr. The photometric orbital variation, determined by these authors, based on their spectroscopic observations, is approximately sinusoidal with an amplitude of $\\sim 0.3$ mag. In spite of the period being relatively long (in the context of Cataclysmic Variables) it was not possible to identify any spectral signature of the secondary star. A distance of 1400 pc was determined from the Na I D line equivalent width, and an $E(B-V) = 0.4 \\pm 0.1$ was suggested from the average of three distinct determinations (4430{\\AA} DIB and Na I D equivalent widths and He II 4686, 10124 {\\AA} lines ratio). Doppler tomography of WX~Cen produced by DS95 suggests that the gas of the secondary has normal chemical abundance, while the wind of the primary star is probably over-abundant in helium (He$^{++}$/H = 0.56). DS95 also showed that, if the primary component of WX~Cen is a white dwarf, then the secondary should have a mass bellow $0.35 M_{\\sun}$ and, therefore, be evolved. If, however, one assumes the hypothesis that the secondary is a main sequence star, it would have $M_2 = 1.16 M_{\\sun}$, and the primary, a mass of $M_1 > 3.5 M_{\\sun} \\pm 0.5 M_{\\sun}$.  \\citet{steiner98} included WX~Cen in a group of 4  galactic binaries defined as the V~Sge stars. The other three objects are V~Sge \\citep{her,dia99}, V617~Sgr \\citep{steiner99,cie} and DI~Cru \\citep{vee02c}. They are characterized by the presence of strong emission lines of O VI and N V. Besides, He II is at least two times more intense than H$\\beta$. The V~Sge stars are similar to the Compact Binary Supersoft Sources (CBSS), seen in the Magellanic Clouds, but not that frequent in the Galaxy. CBSSs are interpreted as suffering hydrostatic hydrogen nuclear burning on the surface of a white dwarf. This burning is due to the high mass-transfer rate, a consequence of the fact that the systems have inverted mass ratios (see \\citet{kah} for a revision and references). \\citet{patt} also considered the hypothesis of Galactic CBSS nature for the stars V Sge, T Pyx and WX Cen, established on their extremely blue colors, high luminosities, orbital light curves and highly excited emission line spectra.  In an observational program to search for galactic WR stars, \\citet{shara} rediscovered WX~Cen as a new WN3, when this star received the WR~48c designation. In our search for V~Sge type stars, we selected WR~48c as a candidate, on the basis of spectroscopic criteria, for a detailed observational study. We, then, realized that WR~48c was the already known star WX~Cen.   \\begin{figure} \\vspace*{10pt} \\centerline{\\includegraphics[width=84mm]{fg1.eps}} \\caption{Average \\textit{V} light curve of the data obtained in 2000, folded with the orbital period and epoch from the photometric ephemeris. \\label{lca}} \\end{figure}    \\begin{figure} \\vspace*{50pt} \\centerline{\\includegraphics[width=84mm]{fg2.eps}} \\caption{The light curve of WX~Cen for the nights 2000 April 24, 25 and 26. The upper light curves are of the comparison star. An increase of 0.08 mag in the mean magnitude is seen in consecutive nights. \\label{lcb}} \\end{figure} ",
        "conclusions": "The main conclusions of this paper are:  \\begin{enumerate}  \\item We confirmed the orbital period of WX~Cen determined by DS95 and refined its value to $P_{orb} = 0.416~961~5 (\\pm 22)$ d, based on spectroscopic observations.  \\item The light curve has total amplitude of approximately 0.32 magnitudes and is non-sinusoidal in the sense of having a narrow, V-shaped, minimum. The object presents flickering with time-scales of tens of minutes. Night-to-night variations of about 0.08 mag are also observed.  \\item We identified most of the emission lines as due to Balmer, He II, C IV, N V, O III, O V, O VI.  \\item The object shows absorption satellites in the Balmer lines with $V = -2900$ km~s$^{-1}$ and in emission with $V = \\pm 3500$ km~s$^{-1}$. These highly variable events remind the satellites in emission in CBSS.  \\item The analysis of the emission lines show that the He II 4686{\\AA} line became weaker between 1991 and 2000/2002 observations. At the same time the emission lines were narrower, suggesting that in 2000/03 the system was in a less active state than in 1991.  \\item The object presents red-shifted emission in high ionization species.  Extended red wing emission is also seen in the strongest lines. The O VI lines show strong variability with velocity of about 700 km~s$^{-1}$.  \\item An analysis of the He II Pickering series decrement shows that the system has significant amount of hydrogen, with an abundance, by number, between 2.1 and 4.5.  \\item We estimate the color excess as $E(B-V)=0.63$  on the basis of the observed DIB at 5780{\\AA}. Given the distance-color excess relation in the direction of WX Cen, this implies a distance of $2.8 \\pm 0.3$ kpc. Interstellar absorption of the Na I D lines show components at $-4.1$ km~s$^{-1}$, which corresponds to the velocity of the Coalsack, and three other components a $-23.9$, $-32.0$ and $-39.0$ km~s$^{-1}$. These components are also seen with similar strengths in field stars that have distances between 1.8 and 2.7 kpc. The intrinsic $B-V$ color index of WX Cen is $(B-V)_0=-0.2$ and the absolute magnitude, $M_V = -0.5$.  \\item A possible explanation for the extended wings in the strong emission lines is that the system has a spill-over stream similar to what is seen in V617~Sgr. We predict that when observed in opposite phase, blue wings would be observed.  \\item The velocity of the satellites is consistent with the idea that the central star is a white dwarf with a mass of $M \\sim 0.9 M_{\\sun}$. With the high accretion rate, it may become a  SN~Ia in a time-scale of $5 \\times 10^{6}$ years.     \\end{enumerate}"
    },
    "0403/astro-ph0403248_arXiv.txt": {
        "abstract": "We have completed an intensive monitoring program of two fields on either side of the center of M31, and report here on the results concerning microlensing of stars in M31. These results stem from a three-year study (the VATT/Columbia survey) of microlensing and variability of M31 stars, emphasizing microlensing events of 3 day to 2 month timescales and likely due to masses in M31. These observations were conducted intensively from 1997-1999, with baselines 1995-present, at the Vatican Advanced Technology Telescope and the 1.3-meter telescope at MDM Observatory, with additional data from the Isaac Newton Telescope, including about 200 epochs total. The two fields monitored cover 560 square arcminutes total, positioned along the minor axis on either side of M31. Candidate microlensing events are subject to a number of tests discussed here with the purpose of distinguishing microlensing from variable star activity. A total of four probable microlensing events, when compared to carefully computed event rate and efficiency models, indicate a marginally significant microlensing activity above that expected for the stars alone in M31 (and the Galaxy) acting as lenses. A maximum likelihood analysis of the distribution of events in timescale and across the face of M31 indicate a microlensing dark matter halo fraction consistent with that seen in our Galaxy towards the Large Magellanic Cloud (Alcock et al.~2000a). Specifically, for a nearly singular isothermal sphere model, we find a microlensing halo mass fraction $f_b=0.29^{+0.30}_{-0.13}$ of the total dark matter, and a poorly constrained lensing component mass (0.02 to $1.5 M_\\odot$, 1 $\\sigma$ limits). This study serves as the prototype for a larger study approaching completion; between the two there is significant evidence for an asymmetry in the distribution of microlensing events across the face of M31, and therefore a large population of halo microlensing dark matter objects. ",
        "introduction": "\\bigskip  The nature of the dark matter in the halo of disk galaxies remains a mystery. While this component of galaxies contributes the majority of their mass, most other data about the dark matter are indirect at best. Some clue as to characteristics of part of the dark matter may be implied by the microlensing detection of masses towards the Large Magellanic Cloud. The MACHO survey revealed a frequency of microlensing that was unanticipated in the context of the visible stellar population alone, but still falls short of accounting for the entire dark matter halo (Alcock et al.~2000a). In particular, the survey indicates a most probable dark matter halo mass fraction of 20\\% (limits of 5\\% to 50\\%, at the 95\\% confidence level), for an indicative spatial distribution model (singular, isotropic, isothermal sphere). Constraints from the EROS magellanic cloud survey (Afonso et al.~2003) yields a constraint consistent with 20\\% microlensing halo fraction, but also consistent with no microlensing halo. There are few independent indications that bear on the validity of the presence of a microlensing halo component. Fluctuations in the brightnesses of images of lensed QSO MG~0414+0534 indicate a significant but sub-dominant halo fraction of microlensing masses (Schechter \\& Wambsganss 2002).  A decade ago we proposed that M31 offers a favorable alternative venue for probing the halo dark matter problem in spiral galaxies, by applying the microlensing techniques to stars in M31 itself, for lenses primarily in M31 but also the Galaxy. In particular, such a signal could be easily distinguished in terms of an asymmetry in microlensing rate across the face of M31, and could be monitored effectively using image subtraction to suppress the severe crowding of M31's stars (Crotts 1992). This required developing techniques for the application of image subtraction to a time series of images (Tomaney \\& Crotts 1996), which led to the first candidate microlensing events in M31 (Crotts \\& Tomaney 1996). At least several more years of such observations have been required to both amass sufficient lensing events for a statistically meaningful sample and to cull out variable stars by extending the baseline. The current work presents the results from this more extended survey and offers our interpretation of the microlensing observations which have resulted. While our survey has found several thousand variable stars, to be reported elsewhere, we have also isolated a sample of events that are more consistent with microlensing events, and imply that a significant fraction of the dark matter halo in M31 is due to objects of stellar mass.  M31 microlensing surveys have also been conducted by the AGAPE groups (Baillon et al.~1993, Calchi Novati et al.~2002), and the survey reported here has been extended by MEGA (Crotts et al.~2001, deJong et al.~2003), the first results of which we mention in conjunction with these in reaching our conclusion. ",
        "conclusions": "How do our results compare to those from Galactic Halo microlensing searches? The most detailed results come from the MACHO survey, reporting (Alcock et al.~2000a) that, during 5.7 years of observations, the detection of 13-17 events towards the LMC when only $\\sim$2-4 events would have been expected due to known, intervening stellar populations. This corresponds to a microlensing optical depth, $\\tau_{LMC} = 1.2_{-0.3}^{+0.4} \\times 10^{-7}$ for $2 < \\hat{t} < 400$ days. The events have timescales of $<\\hat{t}>$ = 34-230 days and the most probable mass for the events is $<m>$ = 0.15-0.9 \\msolar. The maximum likelihood halo fraction is $f_{b} $ = 20\\% (8\\% - 50\\%, 95\\% confidence interval) implying a total mass in MACHOs of M$_{\\rm MACHO}$ = $9_{-3}^{+4} \\times 10^{10}$ \\msolar~ ($r <$ 50 kpc). The EROS survey reports a result consistent with the central MACHO value, but also consistent with no Galactic microlensing halo (Afonso et al.~2003). In addition to searching for events with well-sampled, long-duration microlensing lightcurves, the MACHO \\& EROS groups conducted an analysis (``spike analysis'') in which they searched for very short timescale brightenings in order to place limits on low mass MACHOs. Their conclusion is that Milky Way halo dark matter cannot be comprised of objects in the mass range $2.4 \\times 10^{-7}$ \\msolar $< m <$ $5.2 \\times 10^{-4} M_\\odot$ (Alcock et al.~1996).  The M31 halo microlensing fraction $f_b$ we find is consistent with that seen by MACHO towards the LMC. Our central value is higher (by $0.7\\sigma$ using only our error, or $\\approx 0.4\\sigma$ combining both surveys' errors). We cannot argue for any inconsistency between the two surveys' results.  The positive, marginally significant halo signal we report is due in large part to the asymmetric distribution of events across the face of M31, slightly favoring the far side as would be expected from a microlensing dark matter halo. Our result is approximately as consistent with no halo as that reported by MACHO, however. On the basis of further M31 microlensing evidence, however, we tend to accept the positive halo indication. The VATT/Columbia survey serves as the pilot study for a larger survey, MEGA (``Microlensing Exploration of the Galaxy and Andromeda:'' Crotts et al.~2001), and the first results from this effort (de Jong et al.~2003) also shows a marginally significant farside surplus asymmetry. While this other work does not estimate $f_b$, such an asymmetry is a nearly unique marker of a dark matter halo of microlensing objects. Later data from this survey will encompass almost an order of magnitude more observations, so indicates that the result will become more clear in the near future.  Perhaps the most surprising conclusion drawn from the MACHO data is that the lenses lie in a mass range occupied by stellar objects and are well above the hydrogen burning limit ($\\sim$0.065 \\msolar: Chabrier \\& Baraffe [2000]) which defines the brown dwarf/stellar boundary. Creating plausible models which can explain these results without violating other astrophysical constraints has proven to be quite a challenge for theorists. The only candidates in this mass range to be excluded by direct observation are low-mass stars (Boeshaar, Tyson \\& Bernstein 1994; Graff \\& Freese 1996) and the most viable baryonic candidates are a population of old white dwarfs (WD) which have been previously overlooked because their colors are bluer than expected due to an increase in atmospheric H$_{2}$ opacity (Hansen 1999). Various surveys (Ibata et al.~1999, 2000; Oppenheimer et al.~2001a, b) have now discovered a significant number of such high-proper motion, faint WDs whose kinematics appear to be consistent with those of Galactic halo objects. Alternatively, there are those who would argue that these events are not associated with a dark matter halo but are, rather, explained as lensing by an intervening population of stars along the line-of-sight to the Magellanic Clouds (see Zaritsky \\& Lin 1997, Zhao 1998) or as self-lensing by stars within the LMC/SMC (Sahu 1994). As an alternative, we point out that primordial black holes with masses of $\\sim$1 \\msolar\\, can be formed from density perturbations created during the quark-hadron phase transition in the early universe (Crawford \\& Schramm 1982, Jedamzik 1997). However, although initially baryonic these black holes are classified as non-baryonic dark matter candidates because they form before the epoch of BBN and, therefore, are not subject to its constraints.  Comparing the current survey to MACHO, their constraints on $m$ is somewhat more restrictive. They argue for masses above the 0.07$M_\\odot$ hydrogen core-burning threshold, at a level of certainty of about 90\\%. While in M31 we cannot rule out the possibility that this represents a population of red main-sequence dwarf stars, in our own Galaxy this is ruled out. While our current survey adds little new information as to the nature of these objects, it tends to confirm that the population is seen to exist now in two galaxies, and is likely to be a universal phenomenon. With galaxy halo dark matter accounting for nearly the same fraction of universal closure density as baryons (Rubin 1993), with the WMAP value now set at $\\Omega_B = 0.046$ (Spergel et al.~2002). We are speaking of a contribution of about 1\\%, well within the baryon budget.  An M31 microlensing result might be contaminated by foreground lenses, as has been suggested for the LMC. Since the timescales of Galactic and M31 halo events are similar, given the same lens mass, an M31 survey is susceptible to foreground Galactic halo lensing. Extrapolating previous results (Zhao 1998), an optical depth of $10^{-7}$ with a sheet of matter at 10 kpc from the source (or observer), implies a surface density $\\sim 15 M_\\odot/pc^2$. Our signal implying a microlensing halo corresponds to an optical depth roughly an order of magnitude larger than this. In the Milky way, just judiciously covering the MDM farside field with such a sheet would only require $3\\times 10^5M_\\odot$. In M31, this requires $\\sim 10^9 M_\\odot$. The latter would imply an unusually massive tidal stream, whereas such a possibility in our Galaxy is perhaps reasonable if one ignores the double coincidence of its appearance both in front of the MACHO and current M31 survey fields. For either galaxy, a spherical shell would need about $\\sim 10^{11}M_\\odot$ at 10 kpc to make the optical depth, with mass at other locations scaling as the radius: $M_{shell} \\approx 10^{11} M_\\odot ( D / 10~kpc )$. This approaches a simple re-creation of the original halo lensing mass problem.  From this survey we have insufficient data to state whether these lenses arise in a thick disk or a true halo. To do so will require more events, scattered over a larger portion of the face of M31. Fortunately, such a survey is practible (Baltz et al.~2003) and is currently underway (Crotts et al.~2001). While the results presented here are interesting in their implications for the universality of disk galaxy halo dark matter, we would prefer to have better sampled lightcurves, to cover more of the face of M31 in order to get better leverage on the spatial variation of the lensing population, and to better catalog populations of variable stars in the same field which might masquerade as microlensing events. We look forward to additional progress in using M31 to generalize and extend the insight which has been won from microelensing searches in nearby galaxies."
    },
    "0403/astro-ph0403412_arXiv.txt": {
        "abstract": "We present $\\sim$2$''$-4$''$ aperture synthesis observations of the circumstellar disk surrounding the nearby young star TW Hya in the CO J=2--1 and J=3--2 lines and associated dust continuum obtained with the partially completed Submillimeter Array.  The extent and peak flux of the 230 and 345 GHz dust emission follow closely the predictions of the irradiated accretion disk model of \\citet{calvet_d02}. The resolved molecular line emission extends to a radius of at least 200 AU, the full extent of the disk visible in scattered light, and shows a clear pattern of Keplerian rotation. Comparison of the images with 2D Monte Carlo models constrains the disk inclination angle to $7^{\\circ}\\pm 1^{\\circ}$. The CO emission is optically thick in both lines, and the kinetic temperature in the line formation region is $\\sim$20~K. Substantial CO depletion, by an order of magnitude or more from canonical dark cloud values, is required to explain the characteristics of the line emission. ",
        "introduction": "TW Hya is the closest known classical T Tauri star, and exhibits high X-ray flux, large lithium abundance, and evidence for an actively accreting disk with accretion rate estimated to be from 10$^{-9}$ to 10$^{-8}$ M$_{\\odot}$ yr$^{-1}$ based on the large H${\\alpha}$ equivalent width and the X-ray emission (\\citealp{kastner_h02} and refs therein). Despite its apparently advanced age, estimated to be 5 to 20 Myr, TW Hya is surrounded by a disk of mass $\\sim 5 \\times 10^{-3}$ M$_{\\odot}$ inferred from dust emission (\\citealp{wilner_h00}). The disk is viewed nearly face-on, and is visible in scattered light to a radius of at least $3\\farcs5$, or 200 AU at a distance of 56 pc (\\citealp{krist_s00, trilling_k01, weinberger_b02, calvet_d02}).  Several molecular species (CO, HCN, CN, HCO$^+$, DCO$^+$) have been detected in the TW Hya disk at millimeter and submillimeter wavelengths using single dish telescopes (\\citealp{kastner_z97, vanzadelhoff_v01, thi01, vandishoeck_t03}). The proximity, isolation, and rich chemistry of the TW Hya disk make this system an excellent target for the study of the physical and chemical structure of protoplanetary environments at high spatial resolution using interferometry.  Models may be directly compared to aperture synthesis images of molecular lines to verify the inferences from low resolution observations, and to further probe the kinematics, physical conditions, and chemistry. A first step in this direction was the imaging analysis of HCO$^+$ J=1--0 emission observed with the Australia Telescope Compact Array by ~\\citet{wilner_b03}, though the spatial and spectral resolution were insufficient to examine the kinematics in any detail.  The prospects are excellent for interferometric observations of protoplanetary disks at submillimeter wavelengths (\\citealp{blake02}). 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.} , now nearing completion on Mauna Kea, is an ideal instrument for making these observations (\\citealp{ho_m04}). In this {\\em Letter}, we present early SMA observations of the TW Hya disk in the CO J=3--2 and J=2--1 lines and dust continuum, with a maximum resolution of $2''$. The resulting images unambiguously reveal the rotation of the disk around the star and constrain the disk inclination and size. ",
        "conclusions": ""
    },
    "0403/astro-ph0403138_arXiv.txt": {
        "abstract": "Gamma-ray spectra from cosmic-ray proton and electron interactions with dense gas clouds have been calculated using a Monte Carlo event simulation code, GEANT4. Such clouds are postulated as a possible form of baryonic dark matter in the Universe. The simulation fully tracks the cascade and transport processes which are important in a dense medium, and the resulting gamma-ray spectra are computed as a function of cloud column-density. These calculations are used for predicting the Galactic diffuse gamma-ray spectrum which may be contributed by baryonic dark matter; the results are compared with data from the EGRET  instrument, and used to constrain the fraction of Galactic dark matter which may be in the form of dense gas clouds. In agreement with previous authors, we find useful constraints on the fraction of Galactic dark matter which may be in the form of low column-density clouds ($\\Sigma\\la 10\\,{\\rm g\\,cm^{-2}}$). However, this fraction rises steeply in the region $\\Sigma\\sim10^2\\,{\\rm g\\,cm^{-2}}$, and for $\\Sigma\\ga200\\,{\\rm g\\,cm^{-2}}$ we find that baryonic dark matter models are virtually unconstrained by the existing gamma-ray data. ",
        "introduction": "The nature of dark matter remains one of the outstanding questions of modern astrophysics. The success of the cold dark matter cosmological model (albeit with ``dark energy'' now required: $\\Lambda$CDM) argues strongly for a major component of the dark matter being in the form of an elementary particle. However, the inventory of baryons which we can observe locally falls far short of the total inferred from observations of the cosmic microwave background fluctuations \\citep{Fuk04}, leaving open the possibility that there may be a significant baryonic component of dark matter. Furthermore, although $\\Lambda$CDM is very successful in describing the growth of structure in the universe on large scales, we still lack a direct detection of any of the candidate dark matter particles. Lacking this decisive piece of observational evidence, some authors have proposed models which include a large component of baryonic dark matter. In particular there have been many papers dealing with the possibility that cold, self-gravitating molecular clouds constitute a major component of the dark matter \\citep{Pfe94,DeP95,Hen95,Ger96, Com97, Wal99, Sci00a, Sci00b}. A variety of different forms, including isolated, clustered, and fractal, have been considered for the clouds, but all proposals involve dense gas of high column-density, in contrast to the diffuse gas in the interstellar medium which is easily detected in emission and/or absorption.  One of the fundamental predictions of a model featuring dense gas clouds is the gamma-ray emission resulting from cosmic-ray interactions within the clouds \\citep{DeP95, Kal99, War99, Sci00a}. Because of the potentially large total mass of gas involved, this process may yield a diffuse flux in the Galactic plane comparable to the flux from known sources for photon energies around 1 GeV \\citep{Sci00a}. Considering the high quality data on diffuse emission acquired by the EGRET detector aboard the Compton Gamma Ray Observatory \\citep{Hun97},  it is worth considering this source of gamma-ray emission in detail as it is possible to use these data to constrain the dark matter models (see \\citet{Sal96}; \\citet{Gil94}).  Most previous investigations of this problem have neglected the self-shielding and cascade phenomena which can be important at high column densities \\citep{Kal99, Sci00a}, and have employed emissivities appropriate to the low-density limit. These effects alter the emergent gamma-ray spectrum, and we note that this could be relevant to the observed excess Galactic flux above 1 GeV \\citep{Hun97}. We have noted elsewhere \\citep{Wal03} that massive ($M\\ga10^6\\;{\\rm M_\\odot}$) aggregates of dense gas clouds could potentially account for many of the unidentified discrete sources detected by EGRET \\citep{Har99}.  Here we present detailed calculations of the gamma-ray spectra arising from cosmic-ray interactions with dense gas clouds. We have used a Monte Carlo simulation code, GEANT4, developed for simulating interaction events in detectors used in high-energy particle physics. Not surprisingly, we find that the predicted spectra differ substantially between high and low column-density clouds, and we discuss the interpretation of our results in the context of the observed Galactic gamma-ray emission. Our calculations are undertaken for cold, dense molecular gas in clouds of radius $R\\sim 10^{13}$cm, similar to those proposed by \\citet{Wal98} to explain the extreme scattering events \\citep{Fie94} during which compact extragalactic radio sources are magnified and demagnified as a plasma ``lens'' moves across the line of sight (see \\citet{McK01} for a criticism of this model). However, the results of our calculations depend primarily on the column-density of the individual clouds, $\\Sigma$, under consideration, and their fractional contribution to the Galaxy's dark matter halo, and our results can be taken as representative of other models which are characterised by similar values of these quantities. ",
        "conclusions": "\\subsection{Comparison with data} Several physical processes contribute to the observed diffuse gamma-ray intensity. In order of decreasing fractional contribution these are thought to be pion production and bremsstrahlung from cosmic-ray interactions with {\\it diffuse\\/} Galactic gas, inverse Compton emission from cosmic-ray electrons interacting with ambient photons, and an isotropic background, which is presumably extragalactic and due to many faint, discrete sources. The sum of these contributions offers a good model for the diffuse emission which is actually observed in the 100~MeV--1~GeV band \\citep{Blo89, Hun97}; however, the $E>\\,1$~GeV emission is poorly modeled \\citep{Hun97} with the prediction \\citep{Ber93} amounting to only $\\sim2/3$ of the observed intensity.  Some authors \\citep{Gil94, Sal96} have used the low-energy ($E<\\,1$~GeV) data to argue that the agreement between model and data allows no room for any significant unmodeled emission and that these data therefore place tight constraints on any baryonic contribution to the dark matter halo of the Galaxy. This line of argument is clearly suspect because of the failure of the same model to account for the high-energy ($E>\\,1$~GeV) data. However, even if we ignore the high-energy data the calculations presented in \\S2.3 demonstrate that the gamma-ray constraints on high column density gas clouds in the dark halo are much weaker than those on low column density gas because the emissivity of the latter is much greater; consequently, we revisit the published constraints in the following section (\\S3.2), returning to the question of the high-energy data in \\S3.3.  \\subsection{Published constraints applied to high column-density gas} \\citet{Gil94} argued that the known (diffuse) gas accounts for essentially all of the gamma-ray intensity in the $E>100$~MeV band observed by the COS B satellite \\citep{Blo89}, and he suggested that any contribution from baryonic clouds should amount to no more than $10^{-5}\\,{\\rm ph\\,cm^{-2}s^{-1}sr^{-1}}$ ($E>100$~MeV) at high Galactic latitudes. Toward the Galactic poles we found (\\S3) $Q\\simeq2.5\\times10^{-3}{\\rm g\\,cm^{-2}}$, implying that the average emissivity of the material in the dark halo should be ${\\cal E} < 5.0\\times10^{-2}\\,{\\rm ph\\,s^{-1}g^{-1}}$ for $E>100\\;{\\rm MeV}$. From Table 1 (or figures 6 and 7) we see that this requirement is not met for the models which are in the thin material limit but that any model with $\\Sigma\\ga100\\,{\\rm g\\,cm^{-2}}$ is acceptable. In other words the constraint imposed by \\citet{Gil94} permits the Galaxy's dark halo to be entirely baryonic, provided the individual components have a column-density $\\Sigma\\ga100\\,{\\rm g\\,cm^{-2}}$.  \\citet{Sal96} pointed out that obervations at low Galactic latitudes can provide tighter constraints on any baryonic dark halo, because of the higher intensity expected when looking through the cosmic-ray disk edge-on. This point is manifest in the large values of $Q(b=0)$, which we computed in \\S2.3 (see also figure 8). \\citet{Sal96} argued that a suitable constraint on any baryonic component of the dark halo is that it should not contribute more than the uniform background intensity component observed in any given field. The strongest constraint then comes from observations of the Ophiuchus region \\citep{Hun94}, for which the requirement imposed by \\citet{Sal96} is an intensity contribution from the dark matter halo of $I_D<2.4\\times10^{-6}{\\,\\rm ph\\,cm^{-2}s^{-1}sr^{-1}}$ in the band 300\\ MeV $\\le E \\le$ 500\\$ MeV.  For this line of sight our calculation yields $Q(0, b=15^\\circ)=1.66\\times10^{-2}{\\rm g\\,cm^{-2}}$, and in the thin material limit, for which we find ${\\cal E}\\simeq2.4\\times10^{-2}\\,{\\rm ph\\,s^{-1}g^{-1}}$ (300\\ MeV $\\le E \\le$ 500\\$ MeV; see Table 1), this corresponds to a predicted intensity of $I_D\\simeq3.2\\times10^{-5}{\\,\\rm ph\\,cm^{-2}s^{-1}sr^{-1}}$ in the 300--500~MeV band, implying that at most 8\\% of the dark halo can be resident in low column density clouds.\\footnote{This result is a factor of 2 larger than the corresponding model of \\citet{Sal96} (with $L=3$~kpc and a spherical halo), a difference which is accounted for by their choice of a smaller core radius for the dark halo. \\citet{Sal96} chose a core radius of 3.5~kpc, which yields $Q(0, b=15^\\circ)=3.0\\times10^{-2}{\\rm g\\,cm^{-2}}$, whereas we have employed a core radius of 6.2~kpc (see \\S2.3).} However, our calculations  extend to clouds of higher column-densities, and we find that for $\\Sigma\\ga500\\,{\\rm g\\,cm^{-2}}$ the limit relaxes to the point where all of the dark halo is permitted to be in the form of high column density clouds.  \\subsection{Revision of constraints based on the observed spectrum} \\subsubsection{Galactic diffuse gamma-ray model} The constraints discussed in the previous subsection make reference only to the low-energy ($E<\\,1$~GeV) gamma-ray data. As noted earlier in this section, the observed intensity of the inner Galactic disk at high photon energies ($>\\,1$~GeV) is substantially greater than expected \\citep{Hun97, Ber93}. Much effort has been expended on explaining this discrepancy, with most of the attention given to models in which the Galactic cosmic-ray electron and/or proton spectra differ from their locally measured values \\citep{Bus01, Aha00, Ber00, Str00, Mor97}. However, none of these models offers a satisfactory explanation of the observed mean gamma-ray spectrum of the Galactic disk, and consequently a successful match to the low energy data should not be taken to mean that the emission model is correct. In turn this suggests that the constraints formulated by \\citet{Gil94} and \\citet{Sal96}, on the basis of the low-energy data, may be too restrictive.  The question then arises as to what constraints the gamma-ray data do in fact place on unmodeled emission, given the current state of understanding of the observed emission.  The Galactic diffuse emission model used in \\citet{Hun97} is based on \\citet{Ber93}, and contains the following contributions: \\begin{equation} I = I_{\\rm HI}+I_{\\rm HII}+I_{\\rm H_2} + I_{\\rm IC}+I_{\\rm EG}, \\end{equation} where $I_{\\rm HI}$ is the gamma-ray intensity contributed by cosmic-ray interactions with diffuse atomic hydrogen, and similarly for the ionised and molecular components of the interstellar medium (HII and ${\\rm H_2}$, respectively). The emissivity for these components \\citep{Ber93} is, of course, computed in the thin material limit, and each component therefore has the same spectral shape, differing only in normalisation. Here $I_{\\rm IC}$ is the gamma-ray flux by inverse Compton emission (cosmic-ray electrons up scattering low-energy photons), and $I_{EG}$ is the isotropic extragalactic background flux.  Although the atomic and ionised hydrogen components can be observed directly via their line emission, and are thus well constrained, this is not true for the molecular component. The molecular hydrogen column is assumed to be proportional to the CO emission line strength, as measured by the Columbia CO survey \\citep{Dam87}, for example, but the constant of proportionality (usually denoted by $X$) is unknown and one is forced to determine its value by fitting to the gamma-ray data.  Although the uncertainty in the best-fit determination of $\\langle X\\rangle$ (averaged over the whole sky) is small, the systematic uncertainties are acknowledged to be much larger, ``at least 10\\%--15\\%''  \\citep{Hun97}.  In turn, this estimate of the uncertainty is small in comparison with the differences among the various values of $X$ which have been reported in the literature (see the review by \\citet{Blo89}) and the likely range of variation in $X$ within the Galaxy \\citep{Hun97}. Although these are substantial uncertainties, molecular hydrogen contributes only a fraction of the total observed gamma-ray intensity roughly 20\\% of the local surface density of gas in the Galactic disk is thought to be in molecular form \\citep{Wou90} so the implied fractional uncertainty in the total Galactic emission is perhaps as small as 3\\%. Moreover, uncertainty in $X$ affects only the normalisation and not the spectral shape of the predicted emission from diffuse molecular hydrogen, so the observed high-energy excess cannot be explained in this way even if $X$ could assume an arbitrarily large value.  The other main contributions to uncertainty in the diffuse model are associated with (1) unmodeled spatial variations in the cosmic-ray spectral energy density and (2) unmodeled spatial variations in the low-energy photon spectral energy density; the former affects all of the Galactic contributions to $I$, whereas the latter affects only the inverse Compton component. These uncertainties affect both the normalisation and the spectral shape of the predicted gamma-ray emission; however, the uncertainties are difficult to quantify.  \\subsubsection{Constraints based on the gamma-ray data} For our purposes it is not actually necessary to quantify the uncertainties on the model input parameters; it suffices to use the discrepancy between model and data as a measure of the uncertainty in our understanding of the observed emission. In turn this measure determines the constraints which we can apply to any putative unmodeled emission, such as the contribution from dense gas which we are concerned with here. At photon energies $E>1\\;{\\rm GeV}$ the fractional discrepancy is roughly 60\\% \\citep{Hun97}, in the sense that the observed emission is 1.6 times larger than the model, and we henceforth adopt $0.6/1.6\\simeq40$\\% of the total observed intensity as our estimate of the unmodeled emission. Although this estimate is derived from data at high energies, the effects of the various contributing processes are all very widely spread, and {\\em the estimate therefore applies independent of photon energy.\\/} The constraints appropriate to high/low Galactic latitudes can now be re-evaluated.  At high Galactic latitudes the observed intensity is $I\\simeq1.5\\times10^{-5}\\,{\\rm ph\\,cm^{-2}s^{-1}sr^{-1}}$ for $E\\ge100\\;{\\rm MeV}$ \\citep{Kni96}, implying that any unmodeled emission should be $\\la6\\times10^{-6}\\,{\\rm ph\\,cm^{-2}s^{-1}sr^{-1}}$ in this band. This result is actually slightly stricter than the criterion used by \\citet{Gil94} and thus leads us to tighten our high-latitude constraints, relative to those quoted in \\S3.2: the observed high-latitude gamma-ray intensity constrains the amount of low column-density gas to $\\la20$\\% of the total density of the Galactic dark halo, with this fraction rising to 100\\% for gas clouds of column density $\\Sigma\\ga200\\;{\\rm g\\,cm^{-2}}$.  At low Galactic latitudes we can make use of the mean intensity of the inner Galactic disk, which has been accurately determined by \\citet{Hun97}. For example at 1~GeV the mean intensity ($|l|\\le60^\\circ$, $|b|\\le10^\\circ$) is $\\langle I\\rangle\\simeq 3\\times10^{-8}\\,{\\rm ph\\,cm^{-2}s^{-1}sr^{-1}MeV^{-1}}$, and our calculation of $\\langle Q\\rangle$ for this region yields (\\S2.3) $3.28\\times10^{-2}\\,{\\rm g\\,cm^{-2}}$, implying that the emissivity of the Galactic dark halo material must be, on average, ${\\cal E}\\le4.6\\times10^{-6}\\,{\\rm ph\\,s^{-1}g^{-1}MeV^{-1}}$. By comparison, the actual emissivity of low column-density gas is computed to be (\\S2.3, table 2) ${\\cal E}(1\\;{\\rm GeV})\\simeq1.4\\times10^{-5}\\,{\\rm ph\\,s^{-1}g^{-1}MeV^{-1}}$, implying that $\\la30$\\% of the Galaxy's dark halo may be comprised of low column density gas. For higher column densities the emissivity falls, and table 2 shows that for $\\Sigma=100\\,{\\rm g\\,cm^{-2}}$ the emissivity is only $5.5\\times10^{-6}\\,{\\rm ph\\,s^{-1}g^{-1}MeV^{-1}}$. The gamma-ray data on the inner Galactic disk thus indicate all of the Galaxy's dark halo to be made of dense clouds of column-density $\\Sigma\\ga100\\,{\\rm g\\,cm^{-2}}$.  \\subsection{Comment on the gas content of the galactic disk}  The constraint we have just given is based on the mean spectrum of the inner Galactic disk, in contrast to those given by \\citet{Sal96} who employed limits based on the angular structure of the observed gamma-ray intensity. Specifically, \\citet{Sal96} required that the putative contribution of emission from a baryonic component of the Galaxy's dark halo be less than that of the isotropic component of the observed intensity; this procedure seems to us to be less reliable than the procedure we have employed, for two reasons. First, even if the dark halo were spherically symmetric the emission from any baryonic component would not be, both because our point of observation is quite distant from the centre of the Galaxy and because the resulting gamma-ray emission is strongly dependent on the Galactic cosmic-ray distribution, which in turn is strongly concentrated in the disk of the Galaxy. Indeed the cosmic-ray distribution appears to correlate with the distribution of interstellar matter \\citep{Hun97}, thus complicating the interpretation of the observed gamma-ray intensity distribution. In particular this coupling leads to gamma-ray emission from a baryonic dark halo being correlated with the diffuse gas column density, even if the dense gas is uncorrelated with the diffuse gas. Second, on any given line of sight, such as the Ophiuchus field considered by \\citet{Sal96}, a highly structured dark matter halo might exhibit, by chance, a low dark matter column density. The chances of this are good if, as in the case of \\citet{Sal96}, the field is specifically chosen to have a low ``background'' intensity."
    },
    "0403/astro-ph0403128.txt": {
        "abstract": "{intracluster medium --- active galaxies --- jets --- cooling flows --- feedback --- entropy} Jets and winds are significant channels for energy loss from accreting black holes. These outflows mechanically heat their surroundings, through shocks as well as gentler forms of heating. We discuss recent efforts to understand the nature and distribution of mechanical heating by central AGNs in clusters of galaxies, using numerical simulations and analytic models. Specifically, we will discuss whether the relatively gentle  `effervescent heating' mechanism can compensate for radiative losses in the central regions of clusters, and account for the excess entropy observed at larger radii. ",
        "introduction": "Observational and theoretical lines of evidence suggest that accreting black holes in active galaxies release large amounts of kinetic energy to their environments. Observationally, a minority of active galaxies release the bulk of their power in the form of radio jets (Rees \\textit{et al.} 1982; Begelman \\textit{et al.} 1984). But more numerous radio quiet AGN may also generate much of their power in the form of fast winds, as indicated by recent spectral analyses of broad absorption line (BAL) QSOs in the optical/UV (Arav \\textit{et al.} 2001; de Kool \\textit{et al.} 2001) and X-ray (Reeves \\textit{et al.} 2003) bands. These studies suggest that the kinetic energy in BAL outflows --- which are thought to occur in most radio quiet QSOs --- is larger than previously thought and can approach the radiation output.  Theoretical arguments also point to the importance of the kinetic energy channel.  Numerical simulations suggest that accretion disks, which transfer angular momentum and dissipate binding energy via magnetorotational instability, may inevitably produce magnetically active coronae (Miller \\& Stone 2000).  These are likely to generate outflows that are further boosted by centrifugal force (Blandford \\& Payne 1982).  Unless radiation removes at least 2/3 of the liberated binding energy, very general theoretical arguments indicate that rotating accretion flows {\\it must} lose mass.  The physical reason is that viscous stresses transport energy outward, in addition to angular momentum.  If radiation does not remove most of this energy, then a substantial portion of the gas in the flow will gain enough energy to become unbound (Narayan \\& Yi 1995; Blandford \\& Begelman 1999, 2004). Blandford (this volume) discusses this effect and its consequences in more detail. While it may sometimes be possible to tune the system so that the gas circulates without escaping, any excess dissipation (i.e., increase of entropy) near a free surface of the flow will lead to outflow.  There are several possible sources of such dissipation, including magnetic reconnection, shocks, radiative transport, and the magnetocentrifugal couple mentioned above. If radiative losses are very inefficient, outflows can remove all but a small fraction of the matter supplied at large radii.  While it is probably safe to say that some substantial release of kinetic energy always accompanies accretion, we are not yet able to predict its magnitude, nor how it compares with radiative losses. During the growth of a supermassive black hole to mass $M_{BH}$, the integrated kinetic energy output could be comparable to the total release of binding energy, $\\sim 0.1 M_{BH} c^2$.  It is unlikely to be much smaller than a tenth that amount. Even with a kinetic energy output of $\\sim 0.01 M_{BH} c^2$, the effects on the environment can be dramatic.  At a kinetic energy conversion efficiency of $\\varepsilon_{KE} c^2$ per unit of accreted mass, an accreting black hole liberates $10^{19} (\\varepsilon_{KE}/0.01)$ ergs per gram.  In a galactic bulge with a velocity dispersion of $200 \\sigma_{200}$ km s$^{-1}$, the accretion of one gram liberates enough energy to accelerate $2 \\times 10^4 (\\varepsilon_{KE}/0.01) \\sigma_{200}^2 $ grams to escape speed --- provided that most of the energy goes into acceleration.  Since the typical ratio of black-hole mass to galactic bulge mass is $> 10^{-3}$ (H\\\"aring \\& Rix 2004), feedback from a supermassive black hole growing toward its final mass could easily exceed the binding energy of its host galaxy's bulge.  Such feedback has been invoked in various models to explain the $M_{BH}-M_{bulge}$ and $M_{BH}-\\sigma$ correlations (e.g., Silk \\& Rees 1998; Blandford 1999; Fabian 1999).  The effects of injecting a certain amount of kinetic energy into the environment of a supermassive black hole depend not only on the amount of energy injected, but also on the temporal and spatial characteristics of the injection process and the structure of the ambient medium. Explosive injection of a large amount of energy in a short time --- which might lead to intense, centrally concentrated heating at a shock --- will have very different consequences from gradual, intermittent, or spatially distributed injection, which would allow the surrounding medium to adjust.  Whether mechanical heating is partially offset by radiative cooling is an important factor, as is the presence of absence of small-scale density inhomogeneities.  The speed of a shock or sound wave propagating through a medium with a  `cloudy' thermal phase structure will be highest in the phase with the lowest density (the intercloud medium).  Dense regions will be overrun and left behind by the front, as first pointed out by McKee \\& Ostriker (1977) in connection with supernova blast waves propagating into the interstellar medium. Consequently, most of the energy goes into the gas which has the lowest density (and is the hottest) to begin with.  The global geometric structure of the ambient gas is important as well.   Since a wind or hot bubble emanating from an AGN will tend to follow the path of least resistance, a disk-like structure can lead to a  `blowout' of hot gas along the axis, leaving the disk intact.  Anisotropic injection of the kinetic energy, e.g., in a pair of jets or equatorial, can likewise affect certain regions while sparing others.  Given these complications, it is especially desirable to find a  `laboratory' where one can study the details of mechanical energy injections by active galaxies in a specific context.  Clusters of galaxies can serve this role well.  In the remainder of this paper we will discuss how recent X-ray observations of the intracluster medium (ICM) suggest a particular mode of heating by AGNs, which appears to be susceptible to theoretical modeling.  While mechanical heating is likely to be the only important form of AGN feedback into the hot, highly ionized atmospheres of galaxy clusters, we stress that other forms of energy injection, particularly radiative heating, must occur as well.  These effects are can be particularly important on smaller scales (i.e., the interstellar medium of the host galaxy) and in less highly ionized environments; they are discussed by Ostriker (this volume). ",
        "conclusions": "There is growing evidence, from observations as well as theory, that active galaxies can provide widespread mechanical heating of their environments.  In clusters, this heating apparently occurs in a rather gentle fashion, and is driven by buoyancy.  Simulations of buoyant plumes show significant and fast lateral spreading, generation of sound waves, cool rims of entrained gas surrounding the hot bubbles and a mismatch between X-ray and radio emission, resulting in `ghost cavities'.  The evenness and spatial distribution of heating may be adequate to balance cooling globally, prevent cooling catastrophes, and thus quench so-called `cooling flows'.  Energy spreading to larger radii during cluster assembly might be able to account for the observed entropy excesses in present-day clusters.  While the initial results are promising, numerical simulations have a long way to go before they can adequately the represent the physics of AGN heating.  Three-dimensional simulations are already being done, but their resolution needs to be increased in order to study the effects of mixing and thermal instability.  Magnetic fields, which have played little role in models to date, may have important effects on the dynamics, transport properties (viscosity and thermal conduction), and radio emissivity of clusters.  At the microphysical level, we need to understand why the hot (relativistic?) gas injected by active galaxies appears to mix relatively little with the ICM.  Indeed, we cannot exclude the possibility that some of the injected `cosmic rays' do stream through the thermal background.  If they couple effectively to the ICM via hydromagnetic waves, they will heat the gas in much the same fashion as expanding bubbles, as they traverse the pressure gradient (Loewenstein \\textit{et al.} 1991).  Finally, we need to better understand what sets the efficiency of kinetic energy output from black hole accretion flows, the speed and degree of collimation of the output (winds vs. narrow jets), and feedback effects that couple the evolution of the ICM to the growth rate of the black hole.  Whether (and how) such feedback fixes the $M_{BH}- M_{bulge}$ correlation by regulating the black hole mass, the galaxy mass, or both, remains to be seen."
    },
    "0403/astro-ph0403406_arXiv.txt": {
        "abstract": "We apply the intrinsically symmetrical, decelerating relativistic jet model developed by Laing \\& Bridle for 3C\\,31 to deep, full-synthesis 8.4-GHz VLA imaging of the two low-luminosity radio galaxies B2\\,0326+39 and B2\\,1553+24.  After some modifications to the functional forms used to describe the geometry, velocity, emissivity and magnetic-field structure, these models can accurately fit our data in both total intensity and linear polarization.  We conclude that the jets in B2\\,0326+39 and B2\\,1553+24 are at angles of $64^\\circ\\pm5^\\circ$ and $7.7^\\circ\\pm1.3^\\circ$ to the line of sight, respectively.  In both objects, we find that the jets decelerate from 0.7 -- 0.8$c$ to $<$0.2$c$ over a distance of approximately 10\\,kpc, although in B2\\,1553+24 this transition occurs much further from the nucleus than in B2\\,0326+39 or 3C\\,31. The longitudinal emissivity profiles can be divided into sections, each fit accurately by a power law; the indices of these power laws decrease with distance from the nucleus. B2\\,0326+39 also requires a discontinuity in emissivity to in order to fit a region with several bright knots of emission.  In B2\\,1553+24, the sudden brightening of the jet can be explained by a combination of rapid expansion of the jet and a continuous variation of emissivity. The magnetic fields in both objects are dominated by the longitudinal component in the high-velocity regions close to the nucleus and by the toroidal component further out, but B2\\,0326+39 also has a significant radial component at large distances, whereas B2\\,1553+24 does not. Simple adiabatic models fail to fit the emissivity variations in the regions of high velocity but provide good descriptions of the emissivity after the jets have decelerated. Given the small angle to the line of sight inferred for B2\\,1553+24, there should be a significant population of similar sources at less extreme orientations. Such objects should have long ($\\ga$200\\,kpc), straight, faint jets and we show that their true sizes are likely to have been underestimated in existing images. ",
        "introduction": "\\label{intro}  Evidence that the jets in low-luminosity, FR\\,I \\citep{FR74} radio galaxies are initially relativistic and decelerate on kpc scales has mounted in recent years. FR\\,I sources are thought to be the side-on counterparts of BL Lac objects, in which relativistic motion on parsec scales is well-established \\citep{UP95}.  Superluminal motions have been seen on milliarcsecond scales in several FR\\,I jets \\citep{Giovannini01} and on arcsecond scales in M\\,87 \\citep{Biretta95}. In FR\\,I sources, the lobe containing the main (brighter) jet is less depolarized than the counter-jet lobe \\citep{Morganti97b}. This can be explained if the main jet points toward the observer, suggesting that the brightness asymmetry is caused by Doppler beaming \\citep{Laing88}. The asymmetry decreases with distance from the nucleus \\citep{Laing99}, implying that the jets must decelerate.  Self-consistent models of the deceleration of relativistic jets by injection of matter lost from stars or entrained from the surrounding galactic atmosphere have been calculated by \\citet{Bicknell94}, \\citet{Komissarov94} and \\citet{Bowman96}.  \\citet[][hereafter LB]{LB} fit VLA images of the radio jets in the nearby FR\\,I radio galaxy 3C\\,31 using a sophisticated model to reproduce the observed total and polarized emission within 30\\,arcsec of the nucleus, where the jets are straight. They parameterized the three-dimensional distributions of velocity, emissivity and magnetic-field structure, calculated the brightness at each point within the jets in Stokes $I$, $Q$ and $U$, accounting for the effects of relativistic aberration, and integrated along the line of sight to reproduce the expected distributions on the sky.  They concluded that the jets in 3C\\,31 could be accurately modelled as intrinsically identical, axisymmetric, antiparallel, decelerating, relativistic flows, with locally random but anisotropic magnetic fields. Optimization of the model parameters placed tight constraints on the geometry, velocity, emissivity and field structure. \\citet{LB2} used this velocity field, together with estimates of the external pressure and density from {\\sl Chandra} observations \\citep{Hardcastle02} in a conservation-law analysis based on that of \\citet{Bicknell94}. They showed that there are self-consistent solutions for jet deceleration by injection of thermal matter and derived the variations of pressure, density, Mach number and entrainment rate along the jets. Finally, \\citet{Laing04} developed models of adiabatic, relativistic jets with velocity shear and applied them to 3C\\,31. They demonstrated that such models provide a reasonable description of the emissivity and magnetic-field variations at large distances from the nucleus but fail closer in, and inferred that significant reacceleration of relativistic particles is required precisely where X-ray synchrotron emission is observed \\citep{Hardcastle02}.  In the present paper we apply a modified version of LB's model to the jets of two FR\\,I radio galaxies: B2\\,0326+39 and B2\\,1553+24. Our principal aim is identical to that of LB: to estimate the distributions of velocity, emissivity and magnetic-field structure without introducing preconceptions about the (poorly understood) internal physics. By studying different objects we hope to be able to improve the range of physical scales we are able to probe, to identify which intrinsic features are common to all FR\\,I jets and which vary from object to object and to assess the dependence of the jet structure on power and environmental conditions.  Section \\ref{obs} presents our new, deep, full-synthesis VLA images.  In Section \\ref{model}, we outline the model, emphasizing the parameterizations of the geometry, velocity, emissivity and magnetic field which differ from those used by LB. Section \\ref{results} compares our best-fitting models with the observed data in a variety of ways to show the features that we are able to reproduce as well at those we cannot. In Section \\ref{physical}, we present the velocity, emissivity and field structures of the best-fitting models. Section \\ref{discuss} investigates whether the magnetic-field structure and emissivity are consistent with the assumptions of flux freezing and pure adiabatic energy loss and examines the idea that the jets reaccelerate. We then consider the appearance of the model for B2\\,1553+24 at large angles to the line of sight and the implications for the detectability of side-on counterparts. Finally, we briefly compare the models for 3C\\,31, B2\\,0326+39 and B2\\,1553+24. Section \\ref{ssfw} summarizes our conclusions and outlines possible avenues for further work.  We adopt a Hubble constant, $H_0$ = 70\\,$\\rm{km\\,s^{-1}\\,Mpc^{-1}}$ throughout and define spectral index $\\alpha$ in the sense $S(\\nu) \\propto \\nu^{-\\alpha}$. We use the notation $p = (Q^2+U^2)^{1/2}/I$ for the degree of linear polarization. ",
        "conclusions": "\\label{discuss} \\subsection{Adiabatic models} \\label{adiabatic}  The simplest physical picture of the evolution of the emissivity along a jet assumes that the radiating particles only lose energy adiabatically (synchrotron and inverse-Compton losses being negligible by comparison), that there are no dissipative processes such as particle acceleration or field-line reconnection, and that the magnetic field is convected passively with the (laminar) flow. Analytical equations for the emissivity in this case were first derived by \\citet{Burch79}; these were generalized by \\citet{Baum97} to describe a relativistic, decelerating flow with purely perpendicular or parallel field.  We use the latter expressions, modified to calculate the magnetic-field evolution self-consistently as in \\citet{Laing04}, to estimate the emissivity profiles expected for our models of jet shape and velocity.  We refer to the fits derived in the previous sections and by LB as {\\em free models} in order to distinguish them from {\\em adiabatic models}, following the terminology of \\citet{Laing04}.  \\subsubsection{Magnetic-field profiles} \\label{freeze}  \\citet{Baum97} showed that the variations of the magnetic field components in the quasi-one-dimensional approximation are: \\begin{eqnarray*} B_r &\\propto& (x\\beta\\Gamma)^{-1} \\\\ B_t &\\propto& (x\\beta\\Gamma)^{-1} \\\\ B_l &\\propto& x^{-2}              \\\\ \\end{eqnarray*} in the absence of shear ($x$ is again the jet radius). Given the field components at one point in the jet, we can then predict their evolution from the profiles of radius and velocity given in the previous section.  The predicted evolution of the field components in B2\\,0326+39 and B2\\,1553+24 is shown by the dashed lines in Figs~\\ref{fig:0326bprofiles} and \\ref{fig:1553bprofiles}, respectively. The initial conditions have been chosen to match the free models at the distances from the nucleus given in the figure captions (note that the field is forced to be longitudinal at the nucleus).  The qualitative evolution of the field components is consistent with the free models for both sources: a decrease in the longitudinal component is accompanied by an increase in the toroidal component. Rapid flaring and deceleration act in the same sense, leading to a large decrease in the relative fraction of longitudinal field, as observed. There are significant quantitative discrepancies, however, particularly in the regions where the jets decelerate. In both sources, the transition from longitudinal to transverse field is predicted to be less abrupt than is seen in the free models.  This discrepancy is qualitatively consistent with the effects of velocity shear due to a transverse velocity gradient, which will act to increase the longitudinal component and to delay the onset of the longitudinal to transverse transition.  This cannot be the whole story, however, as there are also large differences between the variations of the toroidal and radial components. The former increases, as expected, but the latter stays roughly constant: in the quasi-one-dimensional approximation, their ratio should remain constant. Similar problems occur in 3C\\,31 (LB; \\citealt{Laing04}).  \\citet{KO89,KO90} showed that the variation of the rms toroidal and longitudinal field components with distance could be very different from that predicted by the simple adiabatic approximation if the flow is turbulent, as is widely believed (e.g.\\ \\citealt{Bic84,DeY96}). Given that the gross departures from the flux-freezing predictions occur where the jets are fast, or are decelerating rapidly, it seems likely that turbulence affects both the magnitude and the configuration of the field in these regions.  \\subsubsection{Emissivity profiles} \\label{ademiss}  The emissivity function $\\epsilon$ can be written in terms of the magnetic field $B$, both as defined in Section~\\ref{Emissivity}, as: \\begin{eqnarray*} \\epsilon \\propto (x^2\\beta\\Gamma)^{-(1+2\\alpha/3)}B^{1+\\alpha} \\end{eqnarray*} \\citep{Baum97,Laing04}.  $B$ can be expressed in terms of the parallel-field fraction $f = \\langle B_l^2 \\rangle ^{1/2}/B$ and the radius $\\bar{x}$, velocity $\\bar{\\beta}$ and Lorentz factor $\\bar{\\Gamma}$ at some starting location using equation 8 of \\citet{Laing04}: \\begin{eqnarray*} B \\propto \\left[ f^2\\left(\\frac{\\bar{x}}{x}\\right)^4 + (1 - f^2)\\left(\\frac{\\bar{\\Gamma}\\bar{\\beta}\\bar{x}}{\\Gamma\\beta x}\\right)^2\\right ]^{1/2} \\end{eqnarray*} We can therefore predict the emissivity using our fitted jet width and velocity together with an estimate of the parallel-field fraction $f$. The resulting emissivity profiles for B2\\,0326+39 and B2\\,1553+24 are shown in Figs~\\ref{fig:0326profiles}(c) and \\ref{fig:1553profiles}(d), respectively. The solid lines show the emissivity profiles from our free model fits and the dashed lines the self-consistent adiabatic profiles, normalized to match at large $z$.  In both objects the adiabatic models agree poorly with the free models where the jet velocities are high. The emissivity falls off much too rapidly in the adiabatic models, just as in 3C\\,31 (LB). This is not a surprise, for the following reasons: \\begin{enumerate} \\item Velocity shear (required by the free model) has been neglected. \\item The magnetic-field evolution is more complicated than expected from simple flux-freezing in an axisymmetric laminar-flow model, even if shear is included (Section~\\ref{freeze}) and turbulence may dominate the field evolution \\citep{KO89,KO90}. \\item Where the jets are fast, we see complex, small-scale, non-axisymmetric structures, indicating that the flow is not laminar. \\item In B2\\,1553+24, there is optical synchrotron emission from the base of the main jet \\citep{Parma03}, implying continuing particle acceleration. \\end{enumerate}  Further from the nucleus, where the velocity has a low and roughly constant value, the adiabatic profile matches the free model very well in both objects. In B2\\,0326+39 this is within a region 4 to 12\\,kpc (Fig.~\\ref{fig:0326profiles}).  In B2\\,1553+24, the region of low, constant velocity extends from 30 to 200\\,kpc and the adiabatic profile agrees very well with the free model over approximately 1.5 orders of magnitude in emissivity (Fig.~\\ref{fig:1553profiles}).  This suggests that the outer jets in this source are modelled surprisingly well as constant-velocity, perpendicular-field, adiabatically-evolving flows.  \\subsection{Sidedness profiles and reacceleration} \\label{acceleration}  Our models, motivated by the gross features of the observed sidedness ratios, assume monotonic deceleration. There are theoretical reasons to expect jets to be reaccelerated by the pressure gradient of the external medium at large distances from the nucleus provided that the mass injection rate is not too large, for example if stellar mass loss dominates the mass injection \\citep{Komissarov94,Bowman96}.  The most obvious effect of reacceleration is a small increase of sidedness ratio with distance from the nucleus.  No such increase is obvious from the ridge-line profiles or images of sidedness ratio (Figs~\\ref{fig:0326ls} and \\ref{fig:1553ls}). We have therefore averaged the sidedness ratios in rings of constant distance from the nucleus in order to improve the signal-to-noise ratio (Fig.~\\ref{avside}).  This reveals considerable sidedness structure in the low-velocity regions. Both sources show sidedness minima, at $\\approx$14\\,kpc from the nucleus in B2\\,0326+29 and at $\\approx$8\\,kpc in B2\\,1553+24. The minimum for B2\\,0326+39 occurs roughly where \\citet{Worrall00} inferred that the synchrotron minimum pressure in the jets \\citep{Bridle91} becomes less than the pressure of the external medium.  Fig.~\\ref{avside} shows the average observed sidednesses as solid lines and our best-fitting models as dashed lines. The dotted lines show the profiles for models with the outer velocity parameters, $\\beta_0$, modified to fit the minimum and maximum sections of the observed profiles in the outer region. The velocity ranges are $\\beta_0 =$ 0.03 -- 0.25 for B2\\,0326+39 and 0.13 -- 0.19 for B2\\,1553+24.  [Note that the model sidedness profiles increase slightly with distance from the nucleus because progressively larger areas of jet edge, which have lower sidedness ratios than the centres, fall below the intensity blanking threshold].  If the changes in sidedness result from acceleration, then they should be associated with variations in the degree of polarization.  In B2\\,0326+39, the field structure can be roughly approximated by two-dimensional field sheets with equal radial and toroidal components (Section~\\ref{0326magnetic}). The expected changes in polarization are then straightforward to calculate \\citep{Laing80}.  For an acceleration from $\\beta = 0.03$ to $\\beta = 0.25$ between 16 and 18\\,arcsec from the nucleus, as implied by a naive interpretation of Fig.~\\ref{avside}(a), $p$ should vary from 0.49 to 0.64 in the main jet and from 0.45 to 0.38 in the counter-jet. These changes are not observed (Figs~\\ref{fig:0326lp}e and \\ref{fig:0326avlp}). The velocity increase is far larger than is expected from the effects of any pressure gradient in the external medium and we also note that the fit of an accelerating adiabatic model to the emissivity profile would be significantly worse than that shown in Fig.~\\ref{fig:0326profiles}.  In B2\\,1553+24, it is more difficult to exclude acceleration as the cause of the increase in sidedness ratio between 8 and 14\\,arcsec from the nucleus (Fig.~\\ref{avside}b): the predicted variations in degree of polarization are below the fluctuation level in Fig.~\\ref{fig:1553lp}(e) and we cannot average across the jet to reduce the noise; the velocity increase required ($\\beta_0 =$ 0.13 -- 0.19) is physically more reasonable than in B2\\,0326+39 and the adiabatic model fit to the emissivity is comparable in quality to that in Fig.~\\ref{fig:1553profiles}.  We conclude that there is no compelling evidence in favour of reacceleration of the jets in either source, and that it is most unlikely to be the sole cause of the increase in sidedness ratio with distance from the nucleus in B2\\,0326+39. Nevertheless, it is expected theoretically, there are hints that it might occur, and it is consistent with the total intensity and polarization data in B2\\,1553+24.  \\begin{figure*} \\includegraphics[width=17cm]{fig20.eps} \\caption{Sidedness profiles of: (a) B2\\,0326+39 at 0.5 arcsec resolution and (b) B2\\,1553+24 at 0.75 arcsec resolution, averaging over all points at the same distance from the nucleus. Solid line: data; dashed line: best fit model; dotted lines: models with $\\beta_0$ modified to fit the minimum and maximum sidednesses in the lower-velocity region, as described in the text. \\label{avside}} \\end{figure*}  \\subsection{The appearance of B2\\,1553+24 at other angles to the line of sight} \\label{1553incl}  We infer that the angle to the line of sight for B2\\,1553+24 is $\\approx$8$^\\circ$. Given that the B2 sample from which it is drawn is selected at the low frequency of 408\\,MHz, the usual assumption is that the emission is isotropic and therefore that the distribution of jet orientations for its members should be random. The source should therefore have many ($\\sim$100) counterparts of comparable total luminosity and larger $\\theta$, but the entire sample of B2 radio galaxies with jets defined in Tables\\,1 and 2 of \\citet{Parma87} has 43 members. A further concern is that the counterparts must be very large. The length of the main jet is at least 60\\,arcsec (Fig.~\\ref{fig:1553.montage}a), corresponding to 388\\,arcsec (340\\,kpc) at $\\theta = 60^\\circ$, the median angle to the line of sight. This is comparable in size with the longest jet in the sample, in NGC\\,315 \\citep{Willis81}, and far larger than the median ($\\approx$30\\,kpc; \\citealt{Parma87}). There is cause for suspicion unless: (a) the side-on counterparts of B2\\,1553+24 are not members of the B2 sample or (b) sources in that sample have linear sizes far larger than previously realised.  In order to understand potential selection effects, we have computed the appearance of the model brightness distribution for B2\\,1553+24 at various angles to the line of sight. We show the results for the median angle to the line of sight for an isotropic sample ($\\theta = 60^\\circ$) and for a source in the plane of the sky ($\\theta = 90^\\circ$) in Fig.~\\ref{fig:1553los}. The area chosen for the plots corresponds to the modelled region for the $60^\\circ$ case.  As well as being very long, the jets appear narrow, straight and (except for a small region around the nucleus) extremely faint.  The total flux density from the inner 200\\,kpc of the model jets (measured along the jets and excluding the core) is constrained to be 44\\,mJy for the $\\theta = 7.7^\\circ$ model but is 27\\,mJy for $\\theta = 60^\\circ$. There are three important selection effects: \\begin{enumerate} \\item The extended flux of the jets shows significant Doppler boosting at small $\\theta$. Given that there is little lobe emission in B2\\,1553+24 \\citep{deRuiter93,NVSS}, the jets probably dominate the flux at low radio frequencies, so end-on sources of this type will be selected preferentially in the B2 survey.  The integral source count $N(>S) \\propto S^{-3/2}$ for the luminosities and redshifts in question, so the size of the parent sample is effectively increased by a factor of $(44/27)^{3/2} \\approx 2$ if isotropic lobe emission is ignored. \\item Except for a small region around the nucleus, the model jets are faint. Further than 50\\,arcsec from the nucleus, the surface brightness is $\\la$0.16\\,mJy/beam area at 8.4\\,GHz with a beam of 2.75\\,arcsec FWHM. This scales to 0.47\\,mJy/beam area at 1.4\\,GHz for a spectral index $\\alpha = 0.6$, comparable with the detection threshold for typical VLA images of sources in the sample presented by \\citet{Fanti86,Fanti87}. In any case, emission on scales $\\ga$120\\,arcsec would not have been imaged reliably in existing VLA observations of the B2 sample (C configuration at 1.4\\,GHz; \\citealt{Fanti87}).  The combination of these two effects makes it extremely unlikely that the outer jets would have been detected. \\item Sources of large angular size with flux densities close to the limit of the catalogue could have been missed by the original B2 survey, which measured peak rather than integrated flux density. We have estimated the magnitude of this effect by convolving the model brightness distribution for $\\theta = 60^\\circ$ with the beam of the B2 survey (3 $\\times$ 10\\,arcmin$^2$). If the short axis of the beam is parallel to the jet axis, the ratio of peak/total flux is 0.73; if the long axis of the beam is aligned, the ratio becomes 0.95. A few sources could therefore have been missed (primarily those orientated E-W). \\end{enumerate} We conclude that the side-on counterparts of B2\\,1553+24 could have escaped identification, either because they were missed in the original survey or because their angular sizes have been greatly underestimated. Morganti \\& Parma (private communication) and Ledlow \\& Owen (2004; in preparation) have made more sensitive observations of radio galaxies in the B2 sample using the WSRT and the VLA in D configuration and have shown that a significant fraction of them have much longer radio jets than have previously been reported, extending many 100's of kpc or even further. We suggest that these form part of the missing population.  \\begin{figure*} \\includegraphics[height=17cm,angle=-90]{fig21.eps} \\caption{The brightness distribution of the model for the jets in B2\\,1553+24 at angles to the line of sight of: (a) $\\theta = 60^\\circ$ and (b) $\\theta = 90^\\circ$. The plots cover 394\\,arcsec (346\\,kpc) in projection on the sky. This corresponds to a distance from the nucleus along the jet axis of 200\\,kpc in both directions for $\\theta = 60^\\circ$, i.e.\\ to the size of the modelled region. The contour levels are 1, 2, 4, 8, 16, 32, 64, 128, 256, 512 $\\times$ 20$\\mu$Jy/beam and the model images have been convolved with a 2.75-arcsec Gaussian beam.\\label{fig:1553los}} \\end{figure*}   \\subsection{Similarities and differences between B2\\,0326+39, B2\\,1553+24 and 3C\\,31} \\label{compare}  In later papers, we will present a re-analysis of 3C\\,31 along with models of two further objects, all with the same functional descriptions of velocity, emissivity and magnetic field. This will allow a full comparison between them; here we make a few general observations about the similarities and differences between B2\\,0326+39, B2\\,1553+24 and 3C\\,31.  Notable similarities are: \\begin{enumerate} \\item All three objects can be modelled successfully by a decelerating jet model. \\item The jets are all initially well collimated, then flare before recollimating and expanding conically. \\item Close to the nucleus the jet velocity on-axis is consistent with a value of $\\beta \\approx 0.8$ in all cases. \\item Deceleration occurs over a region of $\\approx$10\\,kpc to an outer velocity of $0.0 \\la \\beta \\la 0.2$. \\item The ratio of edge to on-axis velocity is consistent with a constant value of $\\approx$0.7 everywhere, although in some places this is poorly constrained. \\item The exponent of the emissivity index is similar in the outermost modelled regions of all the jets. \\item At the jet edge the emissivity falls to $\\approx$0.25 of its on-axis value. \\item The longitudinal field component dominates in the high-velocity regions close to the nucleus and the toroidal component in the outer parts, qualitatively but not quantitatively as expected from flux freezing. \\item The variation of the radial field component is more complex than predicted by simple flux-freezing models. This effect cannot simply be due to velocity shear in a laminar, axisymmetric flow. \\item Close to the nucleus (in the high-velocity and deceleration regions), the variation of emissivity with distance from the nucleus is far less rapid than predicted by the adiabatic approximation. \\item In the low-velocity outer regions, the emissivity variation is reasonably well fitted by a quasi-one-dimensional adiabatic model. \\end{enumerate}  The three objects were expected to have a wide range of angles to the line of sight; this is confirmed by our modelling. In addition, there are some clear differences: \\begin{enumerate} \\item The conical outer region of 3C\\,31 has a large opening angle and is centred on the nucleus.  The outer regions of the other two objects are closer to cylindrical, with very small opening angles and vertices far behind the nucleus. [Note that 3C\\,31 does recollimate outside the modelled region (LB)]. \\item The velocity in the outer region of 3C\\,31 was modelled by LB as decreasing monotonically with distance from the nucleus (as required by the variation of sidedness ratio). In the B2 sources, it has an approximately constant asymptotic value. \\item The close connection between the jet geometry and the forms of the velocity and emissivity profiles inferred for 3C\\,31 is not general. \\item The velocity in B2\\,1553+24 remains at $\\beta \\approx 0.8$ until $\\approx$20\\,kpc from the nucleus (cf.\\ $\\approx$2\\,kpc in the other two objects). \\item There is no need for a discontinuity in emissivity in B2\\,1553+24: the brightening of the jet is consistent with expansion at constant emissivity. \\item Although the biggest single field component at large distances is always toroidal, there are significant differences in the details of the field-component evolution: 3C\\,31 has a mixture of longitudinal and toroidal components; B2\\,0326+39 has toroidal and radial components and B2\\,1553+24 has almost pure toroidal field. \\item In 3C\\,31, the radial field component in the flaring region increases towards the edge of the jet. There is no evidence for this effect in the B2 sources, although it cannot be excluded in B2\\,0326+39, where the signal-to-noise ratio is low. \\item The adiabatic approximation describes the emissivity evolution in B2,1553+24 over a much larger fraction of the jets than in the other two sources. This may be because it applies most accurately in the low-velocity outer region, which is relatively longer in this source. \\end{enumerate}"
    },
    "0403/astro-ph0403630_arXiv.txt": {
        "abstract": "A Monte-Carlo calculation of the atmospheric neutrino fluxes~\\cite{BGS,AGLS} has been extended to take account of the three-dimensional (3D) nature of the problem, including the bending of secondary particles in the geomagnetic field.  Emphasis has been placed on minimizing the approximations when introducing the 3D considerations.  In this paper, we describe the techniques used and quantify the effects of the small approximations which remain. We compare 3D and 1D calculations using the same physics input in order to evaluate the conditions under which the 3D calculation is required and when the considerably simpler 1D calculation is adequate. We find that the 1D and 3D results are essentially identical for $E_\\nu>5$~GeV except for small effects in the azimuthal distributions due to bending of the secondary muon by the geomagnetic field during their propagation in the atmosphere. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403540_arXiv.txt": {
        "abstract": "{ Possible forsterite nanoparticle model of the Extended Red Emission(ERE) is proposed on the basis of photo-luminescence of forsterite after gamma-ray and neutron irradiation. Forsterite exhibits interesting thermoluminescence spectrum similar to ERE of Red Rectangle after irradiation in low temperature. It is shown that the forsterite after thermoluminescence is over exhibits photo-luminescence(PL) when Ultraviolet ray is irradiated. The structure of PL spectrum is almost similar to that of thermoluminescence. In order to explain small variations of the peak position of wavelength of ERE spectrum, possible nanoparticle model of forsterite is investigated. Our model is consistent to the ISO observation data in near and middle infrared region, which suggest the existence of forsterite. ",
        "introduction": "Extended red emission (ERE) is a broad emission band with a peak wavelength between 600 and 850 nm, and with a width between 60 and 120 nm seen in many dusty astrophysical objects such as reflection nebulae, planetary nebulae, HII regions, halos of galaxies, and even in the Diffuse Interstellar Medium. The observation of ERE in the Diffuse Interstellar Medium shows that ERE is a general phenomenon. Though the carrier for ERE is not yet clear, the proposed carriers are hydrogenate amorphous carbon (HAC), quenched carbonaceous composite (QCC), $\\rm C_{60}$, carbon nanoparticles, polycyclic aromatic hydrocarbons (PAHs),  and silicon nanoparticles, and most of them appear to be unable to explain the observed ERE spectra (see Ledoux et al. 2001; Witt et al. 1998, for a summary).  We have suggested that thermoluminescence spectra of forsterite after $\\gamma$-ray irradiation are very similar to ERE of Red Rectangle (Koike K. et al., 2002, here after KK), and have discussed such possibility that thermoluminescence is related to the changes of property of interstellar and circumstellar matter by various irradiation in that space. It is, however, not so plausible that the irradiation energy is sufficient to explain such emissions.  Recently, we have found that the forsterite after thermoluminescence is over exhibits photo-luminescence(PL) when ultraviolet ray (UV) is irradiated. This fact seems to suggest a possible realistic mechanism of ERE. This paper is concerning to this problem.  \\begin{figure*} \\includegraphics[width=17cm,height=8cm]{Forste_UV-F0.eps} \\caption{Photoluminescence spectra of forsterite $(\\rm Mg_2SiO_4)$ single crystals under UV irradiation. The sample is firstly irradiated by $(\\gamma)$-ray in liquid nitrogen, and warmed till about 500K. After thermoluminescence has been over, it is irradiated by UV of a mercury lamp at room temperature.} \\label{Forste_UV_R} \\end{figure*}  Interstellar and circumstellar matter is irradiated by high energy electromagnetic and cosmic ray particles such as $\\gamma$ rays, neutrons, protons and heavy-ions etc. These irradiation will cause some changes on properties such as optical ones of these materials. Especially, it is known that extremely large fluxes of neutrons and $\\gamma$-rays have been emitted during super-nova explosions. Furthermore, interstellar and circumstellar space is typically at extremely low temperature and is always irradiated by electromagnetic radiation and by cosmic ray particles for long time-scale. The effect of this radiation will accumulate in the low temperature environment. In the circumstellar region of both young and evolved stars as well as in the solar system, forsterite and enstatite have been found (Waters et al., 1998ab ; Malfait et al., 1998; Wooden et al., 1999) by the Infrared Space Observatory (ISO) (Kessler et al., 1996). Carbonates such as dolomite $\\rm CaMg(CO_{3})_{2}$, breunnerite $\\rm Mg(Fe,Mn)(CO_{3})_{2}$, calcite $\\rm CaCO_3$, and $\\rm Mg$,$\\rm Ca$-bearing siderite $\\rm FeCO_3$ were found in CI chondrite (Endress, Zinner and Bischoff, 1996). Among these carbonates, Ca-bearing minerals such as dolomite and calcite were also detected in dust shells around evolved stars by ISO (Kemper et al., 2002). Especially, it should be noted that the broad emission feature responsible for extended red emission (ERE) appears at about the 600 -- 900 nm region in many reflection nebulae, and among reflection nebulae the Red Rectangle nebula shows stronger intensity by one order(Witt and Boroson, 1990). In the Red Rectangle nebula, both PAH- and crystalline silicates (forsterite and enstatite)-features were observed by ISO (Waters et al., 1998b). We have suggested, on the basis of the measurement of thermoluminescence for irradiated silicates and carbonates, our thermoluminescence spectra of forsterite at 645--655 nm is very similar to the ERE of the Red Rectangle (KK, 2002).  In this paper, we will investigate the changes of properties of silicates by irradiation in the context of astrophysics. Especially, it is emphasized that the forsterite after thermoluminescence is over exhibits photo-luminescence(PL) when ultraviolet ray (UV) is irradiated. The structure of PL spectrum is almost similar to that of thermoluminescence. Further, possible forsterite nanoparticle model of ERE is also discussed on the basis of this fact.   \\begin{figure} \\centering \\includegraphics[width=7.5cm]{Forste_UV-Red-F-2.eps} \\caption{Comparison with Observation Data of Red Rectangle Spectrum (From Witt A.N. and Boroson T.A. 1990) and Photoluminescence spectra of forsterite under UV irradiation (Fig.1) } \\label{Forste_UV-Red} \\end{figure} ",
        "conclusions": "It should be noted that interstellar and circumstellar space is typically at extremely low temperature and is always irradiated by both electromagnetic radiation and by cosmic ray particles over cosmological timescales. Furthermore, it is well known that extremely large fluxes of neutrons and gamma-rays are emitted during supernova explosions. We have ever suggested that thermoluminescence spectra of forsterite after $\\gamma$-ray irradiation are very similar to ERE of Red Rectangle. However, the energy efficiency of thermoluminescence seems to be not sufficient to explain ERE. In this paper, we have discussed on the fact that the photo-luminescence of irradiated forsterite exhibits similar spectrum to that of thermoluminescence. It should also be remembered that forsterite and enstatite have been found by many ISO observations in many oxygen-rich young and evolved stars.   Is the irradiated forsterite possible to be really one of the carrier of ERE? It is possible so far as its temperature is below about 1000K, because it is known in many insulators the effects of irradiation are almost maintained. However, the irradiation effects on insulators are known to fade out by annealing it to high temperature above this degree. In such case, can forsterite be a carrier of ERE?  In semiconductor silicon, it is well known that the existence of a small amount of impurity caused by a certain kind of minor elements such as As, P, B, Ga or In etc cause new energy levels and forms new type of semiconductors. Though little is known about forsterite, it seems to be possible that a certain kind of impurities cause some semiconductor-like structure of forsterite. Investigation of such possibility is further problem. In this context, it should be noted that the melting point of silicon is $1410^{\\circ}C$, while that of forsterite is very high, and is $1890^{\\circ}C$. It is known that the existence of forsterite is observed in addition to ISO spectrum, in meteorites, while silicon itself is not found in both of them. Thus, together with the investigation of formation condition of silicates in the previous section, the silicon nanoparticle model seems to be highly hypothetical one in the present stage.  We can get easily high-quality crystalline silicon because it is the most fundamental semiconductor, where the solid state properties is very well studied together with the problem of impurity. As for the forsterite, it is very difficult to get the high-quality one. Then, we have used the synthesized one by CZ method in laboratory. For the impurity component in our forsterite sample, Takei and Kobayashi (\\cite{takei}) reported  previously that spectrographic analysis shows that the forsterite sample is pure. However, neutron activation analysis reveals that it contains an infinitesimal quantity of Ir at about 16--18 wt ppm. This level of Ir impurity is below the detection limit of spectrographic analysis. We have also measured our sample using radio activation analysis, and confirmed that our sample is almost pure; that is, other elements except for $\\rm Mg, Si, O$ and an infinitesimal quantity of Ir are not detected.  In order to confirm the reappearance of thermoluminescence of forsterite in other sample, we have get forsterite powder for china and porcelain from a pottery \"Marusu\" (Japan), where details of the method of creation of it is not open. We have investigated thermoluminescence of the \"Marusu\" forsterite, and get almost the same result as our synthesized one."
    },
    "0403/astro-ph0403589_arXiv.txt": {
        "abstract": "We have modeled the unusual orbital light curve of V1494 Aquilae (Nova Aquilae 1999 No.2) and found that such an unusual orbital light curve can be reproduced when there exist two-armed, spiral shocks on the accretion disk.  V1494 Aql is a fast classical nova and found to be an eclipsing system with the orbital period of 0.1346138 days in the late phase of the nova outburst.  Its orbital light curve shows a small bump at orbital phase 0.2, a small dip at 0.3, sometimes a small bump at 0.4, and a large bump at $0.6-0.7$ outside eclipse.  Such a double- or triple-wave pattern outside eclipse has never been observed even though overall patterns look like some supersoft X-ray sources or eclipsing polars. We have calculated orbital light curves including the irradiation effects of the accretion disk and the companion by the hot white dwarf. These unusual patterns can be reproduced when we assume two-armed spiral shocks on the accretion disk.  Especially, triple-wave patterns are naturally obtained.  This result strongly suggests the existence of two-armed spiral shocks on the accretion disk in the late phase of the nova outburst. ",
        "introduction": "Angular momentum transport plays an essential role in accretion disks of cataclysmic variables. Two mechanisms have been proposed so far: one is the turbulent viscosity as adopted in the standard accretion disk model of \\citet{sha73} and the other is the direct dissipation by tidal spiral shocks on the accretion disk as first demonstrated by \\citet*{saw86}.  The turbulent viscosity is a local physical process while the tidal spiral shocks have a global structure on the accretion disk. Therefore, we have a chance to observe a global spiral shock structure when they play an essential role in the angular momentum transport of the accretion disk. Such an observational evidence first came from the Doppler maps of the dwarf nova IP Peg outburst \\citep*[e.g.,][]{ste97}.  We have long believed that tidal spiral structures can be detected even in orbital light curves of cataclysmic variables if the spiral patterns are prominent. At last we find such an evidence of spiral patterns on the accretion disk from the orbital light curves of V1494~Aquilae.  V1494~Aquilae (Nova Aquilae 1999 No.2) was discovered by A. Pereira on 1999 December 1.785 UT at $m_V \\sim 6.0$ \\citep{per99}.  It reached the visual maximum brightness of $m_V \\sim 4.0$ on Dec 3.4 UT ($t_0=$JD 2,451,515.9$\\pm 0.1$) and decayed at the rate of $t_2= 6.6 \\pm 0.5$ days and $t_3= 16 \\pm 0.5$ days \\citep{kis00}.  Early phase spectra were taken by \\citet{fuj99} and \\citet{aya99}, which show P-Cyg profiles of hydrogen Balmer lines with a blueshifted component of 1020 km~s$^{-1}$ and 1200 km~s$^{-1}$, respectively. Thus, V1494~Aql has been classified as a fast nova. \\citet{kat04t} have summarized overall development of the V1494~Aql light curve until autumn of 2003 (for about 3 years).  Detailed spectral features have been reported by \\citet{iij03}.  A short-term periodic modulation was first reported by \\citet{nov00}, who found 0.03 mag variations with a period of 0.0627 days from their 2000 June observation. \\citet{ret00} analyzed 31 nights of CCD photometry during 2000 June-August and suggested a periodicity of 0.13467(2) days.  The full amplitude of the variation in the $R$-band increased from 0.03 mag in June to 0.07 mag in August.  Their folded light curve shows a double-wave structure with a shallower dip at phase 0.5, with about half the amplitude of the main periodicity. \\citet{bos01} reported, based on unfiltered and $R$-band CCD photometry obtained on 12 nights during 2001 June-July, a robust change in the shape and amplitude of the 0.13467 days period. It had an eclipse shape with depth about 0.5 mag. A second shallow eclipse (about 0.1 mag deep) at phase 0.5 can be seen. \\citet{bar03} refined, based on $V$-band photometry during June and September of 2002, the orbital period of 0.1346141(5) days. They pointed out that the orbital light curve is quite unusual in the sense that it is not similar to ordinary cataclysmic variables with a hot spot on the accretion disk. \\citet*{pav03} made a multicolor photometry and concluded that the eclipse depth is deeper in longer wave lengths (deeper in $I$-band than in $V$-band).  \\citet{pav03} further discussed that such an eclipsing characteristic (deeper in longer wave lengths) cannot be explained by the accretion disk or the irradiation effect of the companion because these two light sources have the opposite nature (deeper in short wave lengths). Thus, they suggested, as a model of V1494~Aql orbital light curve, a self eclipsing accretion column in a magnetic polar.  \\citet{kat04t} provided a time-development of the orbital light curves among 2001 November-December, 2002 August, and 2003 June. It sometimes shows a triple-wave structure as well as a double-wave structure.  Kato et al. suggested that some structures (probably on the accretion disk) fixed in the binary rotational frame are responsible for the orbital light curve because of the stability of out-of-eclipse light curve patterns.  Very recently, \\citet*{kis04} pointed out that a close companion to V1494~Aql (located $1\\farcs4$ southwest) is brighter than V1494~Aql itself in the very late phase of the outburst.  They corrected the brightness of the orbital light curve and found that the depth of the eclipse is about twice deeper in magnitudes than before correction.  The depth of the primary eclipse of V1494~Aql has become deeper as the nova decayed.  Such a feature was also observed in the recurrent nova CI~Aql 2000 outburst \\citep[e.g.,][]{mat01}, in which the irradiation of the accretion disk plays an essential role \\citep{hac03ka}.  Therefore, we expect the accretion disk in V1494~Aql is also responsible for the unusual wavy structure of the orbital light curve. In this Letter, we model the orbital light curve. Almost the same orbital light curve model as in the supersoft X-ray sources is adopted \\citep*[e.g.,][]{sch97, hac03kb, hac03kc}, which includes the irradiation effects of the accretion disk and of the companion. In addition, we assume two-armed spiral structures on the accretion disk to reproduce a triple-wave structure outside eclipse. In \\S 2, our numerical model for V1494~Aql is briefly introduced and summarized.  The numerical results are given in \\S 3.  Discussion follows in \\S4.            \\placefigure{v1494_figure_late_i785} \\placetable{system_parameters_V1494_Aql} ",
        "conclusions": "The orbital light curve varies from night to night as shown in Figure 3 of \\citet{kat04t}.  Such a scatter is also clearly shown in the present Figures 2 and 3 for Kiss et al.'s (2004) data. To reproduce this variation, we have changed parameter $\\xi_1$ and $\\xi_2$ independently (see Fig. \\ref{vmag_wd1000m03_quiescence_kiss_2}). The variation appears in orbital phase of $0.1-0.3$ and the upper and lower bounds for these variations can be well reproduced by the models of $\\xi_1=0.26$, $\\xi_2=0.44$ for the upper bound, and $\\xi_1=0.44$, $\\xi_2=0.24$ for the lower bound.  Here, $\\xi_1$ and $\\xi_2$ represent the height of the spiral arm in the rear and front side on the disk of Figure \\ref{v1494_figure_late_i785}, respectively. The other parameters are the same as those in Figures \\ref{v1494_figure_late_i785} and \\ref{vmag_wd1000m03_quiescence_kiss}.  It is remarkable that the peak at orbital phase $\\sim 0.65$ is rather stable and never disappears. This is consistent with the robustness of the period for the out-of-eclipse shape as emphasized by \\citet{kat04t}.  \\citet{pav03} discussed the main light source for the wavy structures of the out-of-eclipse orbital light curve. They argued, as a light source, (1) accretion disk, (2) ellipsoidal shape of the companion, (3) reflection (irradiation) effect of the companion by the hot white dwarf, (4) intermediate polar activity, and (5) polar activity self-eclipsed by the accretion column.  They rejected possibilities (1)-(4) and suggested a viable model of polar activity. Their reasons rejecting possibilities (2) and (4) are reasonable but not for possibilities (1) and (3) as discussed below.  \\citet{pav03} rejected the possibility of accretion disk by pointing out two reasons: (i) the maximum duration of the eclipse of the disk limited by the Roche lobe size and it cannot exceed a quarter of the orbital phase while the overall duration of the eclipse is as large as 0.45 of the orbital period.  (ii) The peak of light emission of the accretion disk lies in the blue region of the spectrum, so that the amplitude of the eclipse should increase as the wavelength becomes shorter (``as the wave length is increased'' in their text is mistaken). However, they observed the opposite trend: the amplitude in the red region is much more deeper than in the visual region.  As for the first conjecture (i), based on the accretion disk model together with the irradiation effects of the disk and the companion, we have already constructed the orbital light curves that match well the observational data. Second, we should be careful with the above statement (ii) because the depth of the primary eclipse varies from night to night as seen in Figure 3 of \\citet{kat04t} and its depth changes as large as $\\Delta R_c \\sim 0.4$ between different periods of the observations as shown in Figures 2 and 3.  Pavlenko et al.'s (2003) $R$ and $I$ light curves are not simultaneous ones but taken in different periods ($R$ is earlier than $I$).  The $V$ light curve is reconstructed from Barsukova \\& Goranskii's (2003) data.  They also rejected  possibility (3) based on the same statement as (ii) discussed above.  However, their multi-color results may simply suggest the fact that the eclipse depth varies from period (night) to period (night). Therefore, we cannot conclude statement (ii) only from the different depths at the different periods.  For the assumed WD luminosity of $L_{\\rm WD}= 3,000~L_\\sun$, the distance modulus is obtained to be $(m-M)_R = 12.05$ and the corresponding distance is estimated to be $d \\sim 1.4$ kpc (see Table \\ref{system_parameters_V1494_Aql}). However, this does not mean the real distance to V1494~Aql because we do not know the real luminosity of $L_{\\rm WD}$ at the time of the two observations. For instance, if we adopt $L_{\\rm WD}= 750~L_\\sun$, the shape of the orbital light curve hardly changes but the distance modulus becomes $(m-M)_R = 11.45$ and $d \\sim 1$ kpc."
    },
    "0403/astro-ph0403010_arXiv.txt": {
        "abstract": "Refraction of wave propagation in a corotating pulsar magnetospheric plasma is considered as a possible interpretation for observed asymmetric pulse profiles with multiple components. The pulsar radio emission produced inside the magnetosphere propagates outward through the rotating magnetosphere, subject to refraction by the intervening plasma that is spatially inhomogeneous. Both effects of a relativistic distribution of the plasma and rotation on wave propagation are considered.  It is shown that refraction coupled with rotation can produce asymmetric conal structures of the profile. The differential refraction due to the rotation can cause the conal structures to skew toward the rotation direction and lead to asymmetry in relative intensities between the leading and trailing components. Both of these features are potentially observable. ",
        "introduction": "Pulsar radio emission is thought to originate in the region deep inside the pulsar magnetosphere populated with a relativistic electron-positron pair plasma (e.g. Blaskiewicz, Cordes, \\& Wasserman 1991; Melrose 2000). The radio waves propagate through the magnetosphere, subject to reflection and refraction (e.g. Barnard \\& Arons 1986; Petrova 2000; Fussell, Luo \\& Melrose 2003, and hereafter FLM) or absorption (Blandford \\& Scharlemann 1976; Lyubarskii \\& Petrova 1998; Luo \\& Melrose 2001; FLM). Since these propagation effects can give rise to observable features in the pulse profile and polarization, study of them can provide insight to the radiation processes inside the magnetosphere.  In the polar cap model,  relativistic electron positron plasmas are produced in the cascade above the polar cap (PC) as a result of rotation-induced particle acceleration (Sturrock 1971; Ruderman \\& Sutherland 1975; Arons \\& Scharlemann 1979; Daugherty \\& Harding 1982;  Zhang \\& Harding 2000; Hibschman \\& Arons 2001). Primary electrons (positrons) are accelerated  to ultra-high energies by a rotation-induced parallel electric field, emitting high energy photons through curvature radiation or inverse Compton scattering. High energy photons decay into electron/positron pairs in the strong pulsar magnetic field, forming an outflowing relativistic pair plasma, referred to as the pulsar plasma, which has a very broad distribution in parallel momenta. Although the mechanism for the radio emission is not well understood, it is generally believed that the emission is produced from collective plasma processes in the pulsar plasma, through either maser emission or plasma instabilities (e.g. Melrose 2000). Regardless of the specific emission mechanism, the radiation must be produced directly on or converted to modes that can escape the magnetosphere to reach the observer.  In the pulsar plasma, there are two escape modes, the X mode with polarization perpendicular to the local plane of the magnetic field and wave vector, and the LO mode with polarization in this plane (e.g. Kennett, Melrose \\& Luo 2000, hereafter KML, and references therein). Observations of jumps between two orthogonal polarizations in the position angle of the radiation support the hypothesis that the radio emission propagates through the magnetosphere in these two orthogonal modes (Stinebring et al. 1984; McKinnon \\& Stinebring 2000). Therefore we only consider propagation of these two modes.  A pulsar plasma is spatially inhomogeneous due to the radial and transverse dependence of the density in the open field line region. Wave propagation in such an inhomogeneous medium can be affected by refraction and reflection, which ultimately determines the emergent pulse profile. Because the dipole field, $B$, decreases in the radial direction as $B\\propto1/R^3$, where $R$ is the radial distance to the star's center, the density of the plasma that streams along the field lines must have the same radial dependence, $N\\propto 1/R^3$, corresponding to the longitudinal characteristic length scale for the density variation, $L\\sim R$, considerably larger than the star's radius. Variation in the plasma density in the direction perpendicular to the field lines is mainly due to the electron/positron pair cascade being nonuniform across the PC, which gives rise to a much smaller (than $R$) transverse length scale of order the radius of the open field line cone. Although the inhomogeneities can be on a still much smaller scale, possibly as a result of nonstationary acceleration or highly localized pair cascades, we only discuss the case that both radial and transverse scales are much larger than the relevant wavelength and that the refraction is the dominant effect that changes the ray path. Wave propagation can be then treated in the geometric optics approximation in which two orthogonal modes propagate independently of each other (Barnard \\& Arons 1986).  The effect of refraction on wave propagation in the pulsar magnetosphere was discussed in the cold plasma approximation by several authors (e.g. Barnard \\& Arons 1986; Petrova 2000; Weltevrede et al. 2003). Numerical study of electron/positron pair cascades above the PC shows that the distribution of a pulsar plasma has a relativistic spread in the plasma rest frame (e.g. Daugherty \\& Harding 1982; Zhang \\& Harding 2000; Hibschman \\& Arons 2001; Arendt \\& Eilek 2002). The relativistic distribution strongly modifies the wave properties (KML) and therefore a self-consistent model for wave propagation needs to include the relativistic effect of the plasma. Past work on refraction of wave propagation generally ignores rotation and the ray path initially in the field line direction would remain two dimensional, confined to the plane of the magnetic field lines. The resultant pulse profile has the same symmetry as the intensity distribution at the emission origin (e.g. Petrova 2000). Observations often show pulse profiles with asymmetric multiple components, i.e. there is asymmetry in intensity as well as location in pulse longitude of conal components (Lyne \\& Manchester 1988; Gangadhara \\& Gupta 1998; Gupta \\& Gangadhara 2003). It is believed that distortion of the profile is due to aberration (Blaskiewicz, Cordes \\& Wasserman 1991; Hirano \\& Gwinn 2001; Gupta \\& Gangadhara 2003) or absorption (Luo \\& Melrose 2001; FLM). Wave propagation in rotating magnetospheres was recently discussed in FLM with emphasis on the cyclotron absorption. Cyclotron absorption including rotation can lead to differential absorption producing asymmetric pulse profiles. In their discussion, FLM assumed the resonance region to be at a substantial fraction of the light cylinder (the radius at which the corotation speed equals the light speed, $c$) where refraction can be ignored.  The purpose of this paper is to explore the effect of refraction on wave propagation in the relativistic pulsar plasma that corotates with the star and the implication of such effect for interpretation of pulse profiles. We consider effects of both a relativistic distribution of the plasma and rotation. Due to rotation rays that originate from the leading and trailing components are subject to different refraction. It is suggested that such differential refraction can significantly distort the pulse profile. Strong refraction occurs in the region with a large density gradient, which is referred to as the refraction region and is located well below the cyclotron resonance region (FLM). To concentrate on the effect of refraction, one assumes the strong magnetic field approximation, in which the X mode is not affected by the medium and propagates approximately in a straight line through the magnetosphere. The LO mode is strongly refracted  and is discussed here. Following a similar procedure to that described in FLM, the ray path in the rotating medium is evaluated numerically in the geometric optics formalism. The emergent pulse profile is obtained at radial distances beyond which the refraction is no longer effective.  In Sec. 2, the wave dispersion of relativistic pulsar plasma is summarized with emphasis on the LO mode in the strong magnetic field approximation; The formalism of ray tracing in the rotating magnetosphere is described in Sec. 3. The ray path is obtained by numerically solving the ray equations including the relativistic distribution and rotation. The result is applied to interpretation of asymmetry of pulse profiles (Sec. 4). ",
        "conclusions": "The effect of refraction on wave propagation in rotating magnetospheres is considered including both effects of relativistic distribution and rotation. Since the relativistic plasma dispersion function is peaked at around $z\\approx1$ and remains small otherwise, inclusion of the intrinsic relativistic effect tends to suppress of refraction at frequencies $\\omega\\geq\\omega_c$. Refraction is sensitive to the plasma density profile in the transverse direction and causes rays to focus and bifurcate, which is qualitatively similar to the result from the nonrotation case (e.g. Petrova 2000). The focusing effect tends to produce a core component even when the emission region has a conal geometry. The bifurcation leads to split of the emission cone into two or more nested cones, giving rise to profiles with multiple components.  One of the distinct features of the model discussed here is the prediction of asymmetry in the pulse profile. The conal components are skewed toward the rotation direction, which is qualitatively in agreement with recent observations (e.g. Gangadhara \\& Gupta 2001). The predicted profiles also show asymmetry in relative intensities between the leading and trailing components, which are common in observed pulse profiles. It is suggested that similar asymmetry seen in observations can be produced by combination of aberration and the differential refraction due to rotation. The latter is the dominant cause for the asymmetry if the radiation is produced at low altitudes.  Two other modes that are not discussed here include the X mode and low-frequency Alfv\\'en mode. In general, pulsar plasma is gyrotropic due to that the electron and positron distributions are not identical. Although the X mode is not purely transverse, with a small longitudinal component of along the magnetic field, one expects refraction of the X mode to be weaker for the LO mode. Inclusion of the X mode would lead to pulse profiles of the two modes that are displaced with one mode dominating any particular part of the pulse longitude. The observational implication of such a displaced profile needs to be explored. The low-frequency Alfv\\'en mode can be refracted in the low-frequency regime. However, the wave becomes subluminal as it propagates to underdense regions and is subject to strong damping.  Pair production above the PC can be oscillatory over the characteristic time scale about $\\Delta t_0=h_0/c$, where $h_0$ is the typical length of polar gap acceleration. The pulsar plasma formed from the cascade may consist of many outflowing clouds (e.g. Usov 1987; Asseo \\& Melikidze 1998). For $h_0\\sim R_*$, we have $\\Delta t_0\\sim 10^{-4}\\,\\rm s$. Since this time can be shorter than the ray propagation time, further work on the wave propagation including the time-dependence of the medium is needed."
    },
    "0403/astro-ph0403226_arXiv.txt": {
        "abstract": "Present in over $45\\%$ of local spirals, boxy and peanut-shaped bulges are generally interpreted as edge-on bars and may represent a key phase in bar evolution. Aiming to test such claims, the kinematic properties of self-consistent 3D N-body simulations of bar-unstable disks are studied. Using Gauss-Hermite polynomials to describe the major-axis stellar kinematics, a number of characteristic bar signatures are identified in edge-on disks: 1) a major-axis light profile with a quasi-exponential central peak and a plateau at moderate radii (Freeman Type~II profile); 2) a ``double-hump'' rotation curve; 3) a sometime flat central velocity dispersion peak with a plateau at moderate radii and occasional local central minimum and secondary peak; 4) an $h_3-V$ correlation over the projected bar length. All those kinematic features are spatially correlated and can easily be understood from the orbital structure of barred disks. They thus provide a reliable and easy-to-use tool to identify edge-on bars. Interestingly, they are all produced without dissipation and are increasingly realized to be common in spirals, lending support to bar-driven evolution scenarios for bulge formation. So called ``figure-of-eight'' position-velocity diagrams are never observed, as expected for realistic orbital configurations. Although not uniquely related to triaxiality, line-of-sight velocity distributions with a high velocity tail (i.e.\\ an $h_3-V$ correlation) appear as particularly promising tracers of bars. The stellar kinematic features identified grow in strength as the bar evolves and vary little for small inclination variations. Many can be used to trace the bar length. Comparisons with observations are encouraging and support the view that boxy and peanut-shaped bulges are simply thick bars viewed edge-on. ",
        "introduction": "} Three-dimensional (3D) N-body simulations of bar-unstable disks have consistently shown that, soon after a bar forms, it buckles and settles with an increased thickness and velocity dispersion, appearing boxy or peanut-shaped (B/PS) when viewed edge-on \\citep[e.g.][]{cs81,cdfp90,rsjk91}. This is particularly important in view of the large number of vertically extended structures commonly refered to as B/PS bulges in edge-on spirals. The recent survey of \\citet*{ldp00a} reveals that at least 45\\% of bulges across all morphological types are B/PS \\citep[but see also the earlier studies of][]{j86,sa87,s87}, suggesting that a significant fraction of bulges may form through dynamical instabilities in disks rather than dissipational collapse \\citep*[e.g.][]{els62} or accretion of smaller systems \\citep[e.g.][]{sz78}.  Verifying such claims is hard, as reliably identifying bars in edge-on systems is difficult. A plateau in the major-axis light profile of an edge-on spiral has long been claimed to indicate the presence of a bar \\citep*[e.g.][]{cc87,hw89,ldp00b}, but axisymmetric features could equally give rise to it and end-on bars would likely remain undetected. A kinematic identification is clearly called for, analogous to the use of longitude-velocity diagrams in the Galaxy (\\citealt{p75,ml86,bgsbu91}; etc). This was first proposed and used in the context of external bulges by \\citet{km95}, while \\citeauthor*{ba99} (\\citeyear{ba99}, hereafter \\citeauthor{ba99}) and \\citeauthor*{ab99} (\\citeyear{ab99}, hereafter \\citeauthor{ab99}) refined those kinematic bar diagnostics using, respectively, periodic orbit calculations and hydrodynamical simulations. These were then applied to relatively large samples of galaxies by \\citet{mk99} and \\citet{bf99}, who successfully showed a close relationship between B/PS structures and bars. Using the ionized-gas kinematics, these works were however only able to probe the galactic potentials in the equatorial plane and were restricted to intermediate and late-type spirals.  Using N-body simulations, we develop in this paper fully self-consistent stellar kinematic bar diagnostics for edge-on systems. Those diagnostics are most easily applicable to gas-poor early-type spirals (where the bulges are large and contain little dust, perfect for photometric studies), and they allow to probe galactic potentials out of the disk plane (to test the barred nature of B/PS structures at large galactic heights). The diagnostics rely exclusively on readily observable quantities and are thus not meant as a study of barred N-body models. They have already been successfully applied by \\citet{cb04} to the complete sample of \\citet{bf99}, and they are being applied at different heights to a subset of the galaxies by Zamojski et al.\\ (2005, in preparation). \\citet{baabdf05} and \\citet*{aab05} present $K$-band observations of the same sample, quantifying the B/PS structure. The goal of all those studies is to study the vertical structure of bars, and ultimately to clarify their relationship to bulges in the context of bar-driven evolution scenarios \\citep[e.g.][]{pn90,fb93,fb95,a03a}.  We review existing bar diagnostics for edge-on disks in \\S~\\ref{sec:past_diagn} and describe the N-body simulations used throughout this paper in \\S~\\ref{sec:nbody}. The basic stellar kinematic bar diagnostics are presented in \\S~\\ref{sec:bar_diagn} along with their viewing angle dependence. \\S~\\ref{sec:time_evol} and \\S~\\ref{sec:incl_dep} describe, respectively, the time evolution of the diagnostics (as the bar evolves) and their inclination dependence (for galaxies close to but not perfectly edge-on). We discuss the uniqueness of the bar diagnostics and compare them to existing observations in \\S~\\ref{sec:discussion}, and our results are summarized in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "} To constrain bar-driven evolution scenarios for the formation of bulges, and more specifically to provide tools to probe the nature of boxy and peanut-shaped (B/PS) bulges, we have used self-consistent 3D N-body simulations of bar-unstable disks to study the kinematic signatures of edge-on bars. Quantifying the major-axis stellar kinematics with Gauss-Hermite polynomials, a number of features can be used as bar diagnostics: 1) a steep quasi-exponential central light profile with a shoulder or plateau at moderate radii; 2) a double-hump rotation curve, possibly showing a local maximum and minimum at moderate radii; 3) a sometime broad central velocity dispersion peak with a shoulder or plateau (and possibly a secondary maximum) at moderate radii; 4) line-of-sight velocity distributions (LOSVDs) with a high-velocity tail (i.e.\\ an $h_3-V$ correlation) over the projected bar length. A local central $\\sigma$ minimum can also be present for strong bars seen approximately end-on. The positions and lengths of those features are further correlated, providing an easy and reliable tool to identify bars in edge-on disks. We note that despite previous claims, a ``figure-of-eight'' is never seen in stellar position-velocity diagrams (PVDs), as expected from realistic models and orbital configurations.  Interestingly, while some of the kinematic bar features can individually be created by axisymmetric density distributions, the $h_3-V$ correlation appears to be a particularly reliable tracer of triaxiality (although it is not uniquely related to it). A number of those kinematic bar signatures can also be used as proxies for the bar length, and thus as an indirect measure of the bar pattern speed. The sharpness of the kinematic features generally decreases for weaker bar and/or increasing viewing angle (from end-on to side-on), introducing some degeneracy between the two but ensuring complementary with morphological bar signatures (strongest for side-on bars). All characteristic features of the kinematic profiles are established the moment the bar forms and grow stronger with time, as the bar lengthens and strengthens. They vary little for small inclination variations but noticeable differences appear for $i\\lesssim80\\degr$.  Existing data and the recent work of \\citet{cb04} show that most B/PS bulges indeed show the kinematic signatures identified here, lending support to evolution scenarios where those bulges are formed through vertical disk instabilities. Detailed work on the morphological and photometric properties of both simulated and observed B/PS bulges are ongoing \\citep{baabdf05,aab05} and should provide more tests of the consistency of these scenarios. Preliminary results support our current conclusions.  The current simulations also systematically produce Freeman Type~II (truncated) surface brightness profiles with approximately exponential peaks, as is observed in most bulges (at least late-types). The local central $\\sigma$ minimum present in the strongest bars (and produced without dissipation) is also increasingly thought to be common in spirals."
    },
    "0403/astro-ph0403156_arXiv.txt": {
        "abstract": "{We present infrared H- and K-band spectra of a companion candidate $3^{\\prime \\prime}$ north of the young star GSC~08047-00232, a probable member of the nearby young Horologium association. From previously obtained JHK-band colors and the magnitude difference between primary and companion candidate, the latter could well be substellar (Neuh\\\"auser et al. 2003) with the spectral type being roughly M7-L9 from the JHK colors (Chauvin et al. 2003). With the H- and K-band spectra now obtained with ISAAC at the VLT, the spectral type of the companion candidate is found to be M6-9.5. Assuming the same age and distance as for the primary star ($\\sim 35$ Myrs, 50 to 85 pc), this yields a mass of $\\sim 25$ Jupiter masses for the companion, hence indeed substellar. After TWA-5 B and HR 7329 B, this is the third brown dwarf companion around a nearby ($\\le 100$~pc) young ($\\le 100$~Myrs) star. A total of three confirmed brown dwarf companions (any mass, separation $\\ge 50$ AU) around 79 stars surveyed in three young nearby associations corresponds to a frequency of $6 \\pm 4~\\%$ (with a correction for missing companions which are almost on the same line-of-sight as the primary star instead of being separated well), consistent with the expectation, if binaries have the same mass function as field stars. Hence, it seems that there is no brown dwarf desert at wide separations. ",
        "introduction": "Brown dwarfs (BDs) are objects with mass below the hydrogen burning mass limit, i.e. below $\\sim 0.078$~M$_{\\odot}$ (e.g. Burrows et al. 1997). It is not yet clear, how BDs form. According to one of the suggested scenarios, accreting companions get ejected as BD embryos by encounters inside a forming multiple system (e.g. Reipurth \\& Clarke 2001). In such a case, BDs may be different from stars regarding multiplicity, kinematics, and disk properties, so that such a scenario can be tested observationaly, e.g. by comparing the secular evolution of star-star binaries with star-BD binaries: If single BDs became single by the ejection from a multiple system, then the fraction of star-BD binaries should decrease faster with time than star-star binaries.  A possible secular evolution in stellar binaries has been studied by Bouvier et al. (2001) and Patience et al. (2002) with as yet inconclusive results. In order to compare such results with star-BD binary fractions, one would first need significant statistics. Hence, several groups search for BDs as companions to stars.  As far as young stars ($\\le 100$ Myrs) in nearby associations ($\\le 100$ pc) are concerned, only two brown dwarfs have been confirmed so far as companions by both common proper motion as well as spectroscopy, namely TWA-5~B (Lowrance et al. 1999, Neuh\\\"auser et al. 2000) in the TW Hya association (TWA, Webb et al. 1999) and HR~7329~B (Lowrance et al. 2000, Guenther et al. 2001) in the Horologium association (HorA, Torres et al. 2000) and/or $\\beta$ Pic moving group (Zuckerman \\& Webb 2000).  In Sect. 2, we present previous and new imaging observation and investigate the astrometry, i.e. whether the pair is co-moving. Our spectroscopic observations, the data reduction, and the final H- and K-band spectra are shown in Sect. 3. We obtain a mass estimate for the companion and interprete the results in Sect. 4. A discussion in terms of BD companion frequency is given in the last section. ",
        "conclusions": ""
    },
    "0403/gr-qc0403111_arXiv.txt": {
        "abstract": "We discuss the Sagnac effect in standard Minkowski coordinates and with an alternative synchronization convention. We find that both approaches lead to the same result without any contradictions. When applying standard coordinates to the complete rim of the rotating disk, a time-lag has to be taken into account which accounts for the global anisotropy. We propose a closed Minkowski space-time as an exact equivalent to the rim of a disk, both in the rotating and non-rotating case. In this way the Sagnac effect can be explained as being purely topological, neglecting the radial acceleration altogether. This proves that the rim of the disk can be treated as an inertial system.  In the same context we discuss the twin paradox and find that the standard scenario is equivalent to an unaccelerated version in a closed space-time. The closed topology leads to preferred frame effects which can be detected only globally.  The relation of synchronization conventions to the measurement of lengths is discussed in the context of Ehrenfest's paradox. This leads to a confirmation of the classical arguments by Ehrenfest and Einstein. ",
        "introduction": "The interpretation of the Sagnac effect is a longstanding problem in the theory of special relativity. Although different approaches of explanation agree on the observable effects, which are in turn consistent with experiments, the interpretation in the context of special relativity is still a matter of debate, as the continuously high publication rate on the subject shows. The problem is closely related with time measurements in moving reference frames and especially with the synchronization of clocks at different positions in rotating systems.  This paper is not intended to be a review of previous work on the subject but instead tries to derive the theory and conventions from first principles in so far as it seems appropriate in the given context. It is meant to be more or less self-contained, so that some overlap with previous articles cannot be avoided. We do not claim a full axiomatic foundation of special relativity, like the one presented in the book of \\citetp{reichenbach24},\\nocite{reichenbach69} of course. For an overview of the Sagnac effect and closely related subjects, including experiments, theoretical analysis and philosophical interpretation, we refer to the articles of \\citetp{post67}, \\citetp{hasselbach93}, \\citet{stedman97} and especially the extensive discussion of synchronization issues by \\citet{anderson98}\\phantom{.} These reviews also include extended lists of references for further reading. Very recently, \\citet{rrf04} compiled a book about ``Relativity in Rotating Frames'' containing a number of articles about the Sagnac effect, its interpretation and related issues. The article by \\citet{rizzi04} themselves comprises another review of the subject.  The outline of this paper is as follows. First we will briefly describe the problem the Sagnac effect poses for special relativity. After this introduction we will describe the principles of relativity. The Sagnac effect will be described using the standard conventions, finding that this is possible without contradictions. This proves that special relativity can be used also on the circumference of a rotating disk. However, we will see that the standard coordinates are not valid globally in this case, which leads us to the discussion of a more general approach. We will derive general coordinates and transformations which are still equivalent and compatible with standard relativity although they use a different convention. We will argue that the synchronization can be chosen arbitrarily and show that in certain situations the choice of a non-standard convention can be more convenient, although not necessary.  We will discuss why the standard coordinates are not valid globally in the situation of a rotating disk, and show that the problem can be formulated without explicitly considering the acceleration. Instead, the closed topology of the rim of a rotating disk has to be taken into account. Locally (which will be defined), the rim is indistinguishable from Minkowski space. Only if our measurement process somehow encloses the complete circle, do we notice the rotation and are able to detect any non-trivial effects. This in a natural way leads us to the central point of this paper, which is the topological interpretation. We will construct a (spatially) closed Minkowski space-time which will be revealed to be exactly equivalent to the rim of a rotating disk. With this model at hand, we can easily avoid some of the difficulties present in previous discussions of the Sagnac effect. This proves that acceleration is not the main reason for the problems of interpretation. Special relativity is valid, and the rim of the rotating disk can even be seen as an inertial frame, as long as the radial direction is not probed by the experiment.  After this discussion, we will turn to the twin paradox which is shown to be closely related to the Sagnac effect, especially when it is discussed on the rotating disk or equivalently in closed space-time. As before, we will propose a scenario which avoids any acceleration effects while it still leads to exactly the same physical effects as the standard situation. In this way the twin paradox is interpreted in terms of space-time topology as well.  Finally, we will discuss measurements of lengths, which are related with the synchronization problem. We will find that, although the synchronization of clocks is a matter of convention, spatial lengths of objects at rest have a meaning as invariant intervals in space-time only if these are measured along lines of standard simultaneity. The measurement of lengths of moving objects, on the other hand, is more a matter of discussion. The result will generally depend on the details of the experiment which implicitly applies a synchronization convention in order to define the lengths. This will lead to our interpretation of the Ehrenfest paradox which agrees with classical approaches but disagrees with several recent publications on the subject.  In all our presentation we will avoid any sophisticated mathematical methods in order not to hide the physical content behind the formalism. We put the emphasis on a self-consistent and logical derivation of the arguments, without hiding any implicit assumptions. Although we tried to include references to most of the relevant publications, we do not claim completeness in the discussion of previous work. ",
        "conclusions": "All the discussions in this paper have in common the importance of clear definitions of the concepts used. Once \\emph{time} and \\emph{length/space} are defined unambiguously, the measurements in all the (real and thought) experiments can be translated relatively easily into their corresponding space-time concepts. The rim of a rotating disk can be treated as an inertial system without any contradiction, as long as the radial coordinate is not probed. Unless the disk is surrounded, the situation is exactly equivalent to that of linear motion as discussed in standard relativity textbooks. When the global structure, i.e.\\ the topology, of the rotating circle becomes relevant, one has to ensure that the locally valid coordinates and space-time descriptions really match or ``mesh'' globally. On the rotating disk, we found that local Minkowski coordinates can not be extended globally,  which gives the explanation for the Sagnac effect. In order to overcome any lack of confidence in our notion that the acceleration is not important for the discussion of effects on the rim of the rotating disk, we showed the equivalence with the situation of a topologically closed Minkowski space-time, which can physically not be distinguished from the rim of a rotating disk. We have to keep in mind that the standard notion of special relativity using Lorentz transformations leads to coordinates which are valid locally. In standard open space-time situations, these coordinates can be extended to form globally valid systems. Periodic boundary conditions or closed space-time topology, on the other hand, will restrict the allowed range of these coordinates. Additional discontinuities or time-lags can appear when matching the local coordinates globally. This leads to preferred frame effects which are purely global and of topological nature. Although most fundamental physics is usually formulated in terms of partial differential equations, we have to keep in mind that such a local description can naturally not explain the world completely. Boundary conditions and thus the topology of space-time can play an equally important role.    \\newcommand{\\mnras}{Mon.\\ Not.\\ R.\\ Astron.\\ Soc.} \\newcommand{\\pz}{Phys.\\ Z.} \\newcommand{\\fp}{Found.\\ Phys.} \\newcommand{\\fpl}{Found.\\ Phys.\\ Lett.} \\newcommand{\\physrep}{Phys.\\ Rep.} \\newcommand{\\grg}{Gen.\\ Relativ.\\ Gravit.} \\newcommand{\\ajp}{Am.\\ J.\\ Phys.} \\newcommand{\\ejp}{Europ.\\ J.\\ Phys.} \\newcommand{\\cqg}{Class.\\ Quant.\\ Gravit.} \\newcommand{\\grqc}{ArXiv Gen.\\ Relat.\\ Quant.\\ Cosmol.\\ e-print} \\newcommand{\\physics}{ArXiv Phys.\\ e-print} \\newcommand{\\ncim}{Nuovo Cimento} \\newcommand{\\rpp}{Rep.\\ Prog.\\ Phys.} \\newcommand{\\ppsa}{Proc.\\ Phys.\\ Soc.\\ London, Sect.~A} \\newcommand{\\comptrend}{C.\\ R.\\ Acad.\\ Sci.} \\newcommand{\\jdp}{J.\\ Phys.\\ (Paris)} \\newcommand{\\annphys}{Ann.\\ Phys.\\ (Leipzig)}  \\let\\cite=\\ocite"
    },
    "0403/astro-ph0403464.txt": {
        "abstract": "{One of the goals of the ground-based support program for the \\corot\\ and \\romer\\ satellite missions is to characterize suitable target stars for the part of the missions dedicated to asteroseismology. We present the detailed abundance analysis of nine of the potential \\corot\\ main targets using the semi-automatic software \\vwa. For two additional \\corot\\ targets we could not perform the analysis due to the high rotational velocity of these stars. For five stars with low rotational velocity we have also performed abundance analysis by a classical equivalent width method in order to test the reliability of the \\vwa\\ software. The agreement between the different methods is good. We find that it is necessary to measure abundances extracted from each line relative to the abundances found from a spectrum of the Sun in order to remove systematic errors. We have constrained the global atmospheric parameters \\teff, \\logg, and [Fe/H] to within $70-100$~K, $0.1-0.2$~dex, and 0.1~dex for five stars which are slow rotators ($\\vsini < 15$~\\kms). For most of the stars stars we find good agreement with the parameters found from line depth ratios, \\halpha\\ lines, Str\\\"omgren indices, previous spectroscopic studies, and also \\logg\\ determined from the \\hipparcos\\ parallaxes. For the fast rotators ($\\vsini > 60$~\\kms) it is not possible to constrain the atmospheric parameters. ",
        "introduction": "\\corot\\ (COnvection, ROtation, and planetary Transits) is a small space mission, dedicated to asteroseismology and the search for exo-planets (Baglin \\etal~2001). Among the targets of the asteroseismology part of the mission, a few bright stars will be monitored continuously over a period of 150 days. These bright targets will be chosen from a list of a dozen F \\& G-type stars, located in the continuous viewing zone of the instrument. The final choice of targets needs to be made early in the project, as it will impact on some technical aspects of the mission. A precise and reliable knowledge of the candidate targets is required, in order to optimize this final choice.  Among the information which needs to be gathered on the candidate targets, fundamental parameters like effective temperature, surface gravity, and metallicity will play a major role in the selection of targets for \\corot. Projected rotation velocities, as well as detailed abundances of the main chemical elements will also be taken into account in the selection process. {\\ bf Thus, the aim of this study is to obtain improved values for the fundamental parameters and abundances of individual elements for the \\corot\\ main targets.}  This information on the targets will subsequently be used for the selected stars, in conjunction with asteroseismological data obtained by \\corot, to provide additional constraints for the modelling of the interior and the atmospheres of these stars.  % Above: Contrib from CATALA % Below: Hans' org. text  In \\sect~\\ref{observations} we summarize the spectroscopic observations and data reduction, in \\sect~\\ref{sec:fundamental} we discuss the determination of the fundamental parameters from spectroscopy, photometry and \\hipparcos\\ parallaxes and we summarize previous spectroscopic studies of the target stars. In \\sect~\\ref{sec:methods} we describe the three different methods we have used for abundance analysis. In \\sect~\\ref{constrain} we discuss how we constrain the fundamental atmospheric parameters using abundance results for a grid of models. In \\sect~\\ref{sec:abundances} we discuss the abundances we have determined. Lastly, we give our conclusions in \\sect~\\ref{conclusions}.    \\begin{table} \\centering \\caption[]{Log of the observations for the spectra of the proposed \\corot\\ targets we have analysed. The signal-to-noise ratio in the last column is calculated around 6\\,500 \\AA\\ in bins of 3~\\kms. \\label{tab:log}} \\begin{tabular}{r|rc|r|r} \\hline HD      & \\multicolumn{1}{c}{Date}           & UT start      & $t_{\\rm exp}$   & S/N  \\\\ \\hline 43318  & 19-Jan-98      & 22:38         &  900            & 120 \\\\ 43587  & 14-Jan-98      & 22:32         & 1800            & 250 \\\\ 45067  & 15-Jan-98      & 22:20         & 1200            & 260 \\\\ 49434  & 17-Jan-98      & 22:33         & 600             & 160 \\\\ 49933  & 21-Jan-98      & 22:17         & 1200            & 210 \\\\ 55057  & 15-Jan-98      & 23:10         & 900             & 270 \\\\ 57006  & 10-Dec-00      & 00:41         & 1800            & 250 \\\\ 171834  & 05-Sep-98      & 19:00         & 1800            & 300 \\\\ 184663  & 18-Jun-00      & 01:30         & 1000            & 170 \\\\ 46304  & 17-Jan-98      & 22:05         & 600             & 170 \\\\ 174866  & 04-Jul-01      & 00:39         & 1500            & 190 \\\\ \\hline \\end{tabular} \\end{table}    \\begin{table} \\centering %\\footnotesize \\caption{Str\\\"omgren photometric indices of the \\corot\\ main targets taken from Hauck \\& Mermilliod (1998). HD~46304 and HD~174866 are shown separately: abundance analysis has not been made for these two stars due to their high \\vsini.\\label{tab:stromgren_corot}}  \\begin{tabular}{r|ccccc} \\hline \\multicolumn{1}{c|}{HD} & $V$  &  $b-y$ & $m_1$ &   $c_1$ & ${\\rm H}_\\beta$ \\\\ \\hline 43318  & 5.65 & 0.322 & 0.154 &   0.446 &   2.644 \\\\ 43587  & 5.70 & 0.384 & 0.187 &   0.349 &   2.601 \\\\ 45067  & 5.87 & 0.361 & 0.168 &   0.396 &   2.611 \\\\ 49434  & 5.74 & 0.178 & 0.178 &   0.717 &   2.755 \\\\ 49933  & 5.76 & 0.270 & 0.127 &   0.460 &   2.662 \\\\ 55057  & 5.45 & 0.185 & 0.184 &   0.876 &   2.757 \\\\ 57006  & 5.91 & 0.340 & 0.168 &   0.472 &   2.625 \\\\ 171384  & 5.45 & 0.254 & 0.145 &   0.560 &   2.682 \\\\ 184663  & 6.37 & 0.275 & 0.149 &   0.476 &   2.665 \\\\ \\hline 46304  & 5.60 & 0.158 & 0.175 &   0.816 &   2.767 \\\\ 174866  & 6.33 & 0.122 & 0.178 &   0.960 &   2.822 \\\\ \\hline \\end{tabular} \\end{table} ",
        "conclusions": "}  \\begin{itemize} \\item  We have performed a detailed abundance analysis of nine potential \\corot\\ main targets. The accuracy of the abundances of the main elements is of the order 0.10~dex when including possible errors on microturbulence and inadequacies of the applied 1D LTE atmospheric models.  \\item We have compared three different methods for the analysis which show very good agreement. The discrepancies are due to the different parameters we have used for the stellar atmosphere models.  \\item We have found no evidence for chemically peculiar stars.  \\item We have constrained \\teff, \\logg, and metallicity to within $70-100$~K, $0.1-0.2$~dex, and 0.1~dex for the five slow rotators in our sample. For the four moderate rotators we cannot constrain the fundamental parameters very well, \\ie\\ \\teff, \\logg, and [Fe/H] to within $200$~K, $0.5$~dex, and 0.15~dex.  \\item For most of the stars our results for the fundamental parameters agree with the initial estimates from Str\\\"omgren photometry, line depth ratios, and \\halpha\\ lines, and the \\hipp\\ parallaxes. For HD~43318 we have found a somewhat lower \\teff\\ and \\logg. For HD~49933 we find a \\teff\\ about 200~K hotter than previous studies. For the fast rotators HD~171834 and HD~184663 we find a quite high \\logg\\ value compared to other methods, but the uncertainty on our estimate is large (0.5 dex).  \\item  For the two \\corot\\ targets HD~46304 and HD~174866 abundance analyses were not possible due to the very high \\vsini.  \\end{itemize}  Suggestions for future studies of the \\corot\\ targets:  \\begin{itemize}  \\item From the comparison of two independent analyses (method A and C; \\cf~\\ref{sec_comp_meth}) we have found evidence that a careful (re-)normalization of the spectra may be very important but this must be investigated further.  \\item To probe the interior of the stars with the asteroseismic data from \\corot\\ we need to know the abundances of elements which affect the evolution of the stars, namely C, N, and O. Thus, spectra that cover the infrared regions with good C, N, and O lines must be obtained.  % \\item To observe spectra with  higher resolution could improve %       the abundance estimates and hence the ``fine-tuning'' of the atmospheric %       parameters. % This will give more precise EQW -- but this is not the %                   % main uncertainty here !! Models mv. are far worse error %                   % sources.  \\item Recent 3D atmospheric models should also be used in the analysis to explore the importance of the models.  % We have shown that the \\vwa\\ software % can potentially be used for larger samples of stars.  %   We note that a future paper in this series %   will present similar results for the proposed \\romer\\ targets.  \\end{itemize}  We finally note that the next paper in this series will give the results of the abundance analysis for some recently proposed \\corot\\ primary targets, the possible secondary \\corot\\ targets, as well as the proposed targets for the \\romer\\ mission.  %=================================================================="
    },
    "0403/astro-ph0403360_arXiv.txt": {
        "abstract": "We have identified two quiescent GRBs (bursts having two or more widely-separated distinct emission episodes) in which the post-quiescent emission exhibits distinctly different characteristics than the pre-quiescent emission. In these two cases (BATSE GRBs 960530 and 980125), the second emission episode has a longer lag, a smoother morphology, and softer spectral evolution than the first episode. Although the pre-quiescent emission satisfies the standard internal shock paradigm, we demonstrate that the post-quiescent emission is more consistent with external shocks. We infer that some observed soft, faint, long-lag GRBs are external shocks in which the internal shock signature is not observed. We further note that the peak luminosity ratio between quiescent episodes is not in agreement with the ratio predicated by the lag vs. peak luminosity relationship. We briefly discuss these observations in terms of current collapsar jet models. ",
        "introduction": "There is strong evidence that the majority of gamma-ray burst (GRB) pulses result from internal shocks in relativistic winds \\citep{sar97a,kss97,dm98,rrf00,np02}. However, a number of bursts exhibit a soft component indicative of external shocks that could be interpreted as onset of afterglow \\citep{gib99,bur99,gib02}. This emission begins preferentially towards the end of the burst or even after the GRB has ended (as suggested by the BeppoSAX WFC and NFI observations of serveral x-ray afterglows \\citep{cos00}), but can also appear at a time early enough to overlap the short-timescale emission (as observed in GRB 980923 \\citep{gib99}). In addition, extended soft $\\gamma-$ray emission has been observed by co-adding the fluxes of many BATSE bursts; this might also be an indicator of the same external shock phenomenon \\citep{con02}.    ``Quiescent'' GRBs may provide new clues to the origins and physics of GRB shocks. Quiescent bursts are GRBs that release their gamma-ray energy in more than one distinct emission episode \\citep{rm01,rmr01}; e.g. they have at least one extended period during which emission is absent. A rigorous definition of quiescent time can be problematic due to the diverse nature or complexity of GRB temporal structures.    \\citet{rm01} introduced a quiescent GRB definition for a sample of bright BATSE bursts that used a sliding temporal window with a width of $0.05 \\times {\\rm T90}$ on summed, background-subtracted, BATSE 4-channel 64-ms data. These authors chose GRBs which had at least one time period for which the count rate dropped below $2\\sigma$ of the summed background rate. This criterion was applied to only the brightest, longest bursts in the BATSE 4B Catalog \\citep{pac99}.    It is possible that this definition is too permissive, since it might allow for inclusion of some GRBs that do not have distinct emission episodes. For example, some bright bursts that do not satisfy this criterion (because the first episode's decay and the second episode's rise overlap) might be classified as quiescent if they are observed at lower signal-to-noise, as might some bursts with faint, long, underlying pulses. Because the existing quiescent definition might be too permissive, we have tentatively further required quiescent GRBs to be those in which the emission drops to the background for an extended period of time that is greater than or equal to the duration of the previous emission episode. This ensures that the quiescent time is greater than the interpulse duration.                            In their study of very bright quiescent GRBs, \\citet{rm01} operated under the assumption that the post-quiescent episodes are due to internal shocks. We offer the hypothesis that quiescent GRBs potentially provide a laboratory in which {\\it external} shock signatures indicative of the onset of afterglow may be isolated from the rest of the burst. We have thus far found two GRBS that support this hypothesis. ",
        "conclusions": "At least two quiescent GRBs exhibit multiple lags. The short-lag component is consistent with internal shocks while an external shock explanation is preferred for the smooth long-lag component. The long-lag component is broad and faint, and is possibly present but unrecognizable in many bursts. The apodization of \\citep{nor00} implicitly recognizes its existence in order to remove it. It is possible that many or all of the long-lag bursts in the sample of \\citet{nor02} exhibit only the smooth and broad long-lag component.  Because the short and long lags found in these quiescent GRBs are so different, we have assumed that there are two distinct GRB pulse types. This hypothesis needs to be verified. The best way to approach this problem is via statistical and/or data mining studies of GRB pulse characteristics. It is important to determine the existence of one continuous pulse distribution or poossibly distinctly different distributions. Detailed analyses of BATSE and $Swift$ GRB data are planned to corroborate this work.  The lag vs. peak luminosity relation apparently requires re-calibration. It is not clear that long-lag pulses obey the same relation as short-lag pulses, or if they even obey a lag vs. peak luminosity relation. Since there is a correlation between lag and variability, the variability vs. peak luminosity relation \\citep{rei01} might require similar re-calibration.  The simplest explanation for the large number of long-lag GRBs found in the BATSE database is that they are GRBs viewed from off the jet axis. The tightly-beamed, variable, short-lag internal shocks would not be seen from large viewing angles, whereas the more widely-beamed, smooth, long-lag external shocks would be. However, it is difficult to reconcile this explanation at the present time with some of the other similar phenomena associated with GRBs (e.g. x-ray rich GRBs, Gamma-Ray Flashes, and Orphan Afterglows)."
    },
    "0403/astro-ph0403683_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{intro}  We present here a brief review of energy and power spectra, radiative processes giving rise to them, spectral states, and some aspects of variability of black-hole binaries. Our emphasis is on interpretation of the X-ray and soft \\g-ray (hereafter X\\g) observations in terms of simple physical models and unifying various aspects of the wealth of the present observational data. We also discuss millisecond flares recently discovered from Cyg X-1\\cite{gz03}, where we present some important new results. We refer the reader to Ref.~\\citen{mr04} for an exhaustive review of observational aspects of black-hole binaries, and, e.g., to Refs.\\ \\citen{d02,z99,b04} and \\citen{z00} for some other relevant reviews.  In our presentation, we use the so-called $\\nu F_\\nu$ representation for both energy and power spectra. The motivation for this choice is provided by a trivial mathematic consideration. Namely, any differential dependence, e.g., flux per unit energy, ${\\rm d}F/{\\rm d}E$, equivalently expressed as $F(E)$ or $F_E$, can be plotted as $F(E)$ vs $E$. However, when we plot $F$ vs.\\ $\\log E$, it should be $F(\\log E)={\\rm d}F/{\\rm d}\\log E$, i.e., $F(E) ({\\rm d}\\,\\log E/{\\rm d}E)^{-1}\\propto E F(E)$. This gives us $F$ per a logarithmic interval of $E$. Graphically, $EF(E){\\rm d}\\log E$ represents the area under the curve, $F$, in a plot vs $\\log E$. Then, a peak in this representation shows us at which photon energy range most of the power is radiated, or which frequency range corresponds to most of variability of a source. An added benefit of this representation is that plots vs.\\ photon energy or wavelength become identical, as $\\lambda F(\\lambda) = E F(E)$. We also plot the energy spectra with the same length per decade on each axis, in order to enable direct comparison of the shape spectra on different figures.  On the other hand, it is common in X-ray astronomy to plot photon flux vs.\\ energy. However, this makes it incompatible with most of other branches of astronomy (e.g., those from the radio to the UV), where energy flux is used. Also, energy is a much more fundamental quantity in physics than photon number, with many physical laws concerning the former, not the latter (as photon number is not conserved in most of physical processes). Indeed, all quantities of radiative transfer, e.g., flux, specific and average intensity, radiation pressure, are defined in terms of energy.  A practical consequence of plotting photon number per unit energy on a logarithmic plot is that such spectra look usually extremely steep, covering many orders of magnitude on the vertical axis, appear very similar to each other, and make it very difficult to actually see any spectral features. This is due to the artificial steepening of such a spectrum by $E^{-2}$, described above.  It is also very common in X\\g\\ astronomy to plot instrumental counts instead of photons or energy. The motivation for it is that high-energy instruments have rather limited energy resolution, which makes the process of deconvolution not unique. However, this purist point of view would actually require plotting only instrumental counts, as plotting together predictions of a model and/or residuals is already model-dependent. In fact, residuals given as, e.g., ratios or contributions to $\\chi^2$, are {\\it identical\\/} for either counts vs.\\ model or a deconvolved spectrum vs.\\ model. However, the latter has the enormous advantages of presenting a physical spectrum, allowing comparison with other spectra, and showing features originating in the source. The count-rate plot, though model-independent, is by its nature instrument-dependent, showing mostly features of the instrumental response, not related to the physics of the source. ",
        "conclusions": "We have reviewed spectral states of black hole binaries and their respective dominant radiative processes. In the hard state, the dominant process is thermal Comptonization of soft photons, most likely blackbody ones  from an outer accretion disk. The thermal character of the distribution of the scattering electrons is probably related to the most likely source geometry, consisting of an inner hot accretion flow\\cite{a95,ny95,z98}, see Fig.\\ \\ref{geo}a. The electrons in the flow are powered by energy transfer from the hot ions, which receive most of the available gravitational energy. The stochastic nature of the energy transfer, probably via Coulomb interactions, results in the electron distribution being nearly thermal.  In addition, the cold medium gives rise to reprocessing, including Compton reflection and fluorescence. Relativistic broadening observed in some cases indicates that, in addition to an outer disk, there is some cold medium close to the black hole, e.g, from collapse of the innermost hot flow\\cite{y01}. The variability pattern observed on long timescales indicates that the dominant driver of the variability is changing irradiation of the hot flow by soft seed photons from the outer cold disk\\cite{z02,z03}.  \\begin{figure}[t!] \\centerline{\\includegraphics[width=11cm]{hard.ps}} \\vskip 0.7cm \\centerline{\\includegraphics[width=11cm]{soft.ps}} \\caption{(a) A schematic representation of the likely geometry in the hard state, consisting of a hot inner accretion flow surrounded by optically-thick accretion disk. The hot flow consitutes the base of the jet (with the counter-jet omitted from the figure for clarity). The disk is truncated far away from the minimum stable orbit, but it overlaps with the hot flow. The soft photons emitted by the disk are Compton upscattered in the hot flow, and emission from the hot flow is partly Compton-reflected from the disk. (b) The likely geometry in the soft state consisting of flares/active regions above an optically-thick accretion disk extending close to the minimum stable orbit. The soft photons emitted by the disk are Compton upscattered in the flares, and emission from the flares is partly Compton-reflected from the disk\\cite{z02}. } \\label{geo} \\end{figure}  In the soft states, Compton scattering is by a hybrid, thermal/nonthermal electron distribution. The scattering plasma forms, most likely, active regions above the surface of an accretion disk, which emit the energetically dominant blackbody emission, see Fig.\\ \\ref{geo}(b). The active regions are powered by energy transfer from inside the disk, most likely via magnetic buoyancy and subsequent magnetic field annihilation, leading then to acceleration of particles. This is inherently a nonthermal process, which explains the observed nonthermal photon spectra, showing no high-energy cutoff up to high energies. The ratio of the power in the active regions to the disk power increases in the sequence of the soft states: ultrasoft, high, intermediate/very high. Thus, the last state is characterized by a powerful corona surrounding most of the disk, in a modification of the picture of Fig.\\ \\ref{geo}(b).  In a range of $L/\\ledd$, either the hard/hot and soft/cold accretion flows are possible. This leads to hysteresis in the long-term lightcurve of black-hole binaries, and the corresponding long-term limit cycle.  A very interesting new phenomenon in black-hole binaries is that of millisecond flares, extreme events occurring on timescales comparable to $GM/c^3$. In particular, the strongest flare discovered from Cyg X-1 appears to be the most extreme event detected from black-hole accretion yet."
    },
    "0403/astro-ph0403611_arXiv.txt": {
        "abstract": "Using the deepest and finest resolution images of the Universe acquired with the Hubble Space Telescope and a similar image taken 7 years later for the Great Observatories Origins Deep Survey, we have derived proper motions for the point sources in the Hubble Deep Field--North. Two faint blue objects, HDF2234 and HDF3072, are found to display significant proper motion, 10.0 $\\pm$ 2.5 and 15.5 $\\pm$ 3.8 mas yr$^{-1}$. Photometric distances and tangential velocities for these stars are consistent with disk white dwarfs located at $\\sim$ 500 pc. The faint blue objects analyzed by Ibata et al. (1999) and Mendez \\& Minniti (2000) do not show any significant proper motion; they are not halo white dwarfs and they do not contribute to the Galactic dark matter. These objects are likely to be distant AGN. ",
        "introduction": "Major observational campaigns have searched for dark matter in the form of massive compact halo objects (MACHOs) using microlensing events (e.g. Alcock et al 1997; Afonso et al. 2003; Udalski et al. 1992). The detection of 13--17 microlensing events toward the Large Magellanic Cloud during 6 years by the MACHO collaboration implies that a significant fraction (20$\\%$) of the halo of the Galaxy may be in the form of compact halo objects (Alcock et al. 2000). The time scale of these lensing events eliminates the possibility of MACHOs having substellar masses. The MACHO collaboration finds a most probable mass of 0.5 \\msun\\, which supports the idea of a massive halo comprised of baryonic matter in the form of low luminosity white dwarfs (Kawaler 1996). Recent observations by the EROS group provide further evidence that less than 25\\% of a standard dark matter halo can be composed of objects with a mass between 2 $\\times$ 10$^{-7}$ \\msun\\ and 1 \\msun\\ (Afonso et al. 2003).  Halo white dwarf stars are expected to have large proper motions as a result of their high velocities relative to the Sun. HST proper motion studies of the Globular Cluster NGC 6397 showed that most of the required dark matter in the solar vicinity can be accounted for by a population of old white dwarfs representing the thick disk and halo of the Galaxy (Mendez 2002). Claims by Oppenheimer et al. (2001) and Ibata et al. (2000) that they had found a significant population of halo white dwarfs from kinematic surveys are tantalizing. Their discoveries seemed to be consistent with earlier findings of an old population of white dwarfs in the Hubble Deep Field (Mendez \\& Minniti 2000). However, further analysis by several groups showed that the sample of Oppenheimer et al. (2001) could also be interpreted as the tail of a kinematically warmer white dwarf component, better explained by the thick disk population of the Galaxy (Reid et al. 2001; Reyle et al. 2001; Mendez 2002; Bergeron 2003).  The Hubble Deep-Field (HDF) provides a unique window on the Universe (Williams et al. 1996; Flynn et al. 1996). The extreme depth of the HDF provides an unprecedented advantage to find faint stellar objects as well as to study very distant galaxies. The advantage of going deep is that it allows us to search for faint stellar components of the Galaxy in the regions of the color--magnitude diagram that are devoid of any contamination by standard Galactic stars. The lack of ordinary disk stars is due to the finiteness of the Galaxy (Flynn et al. 1996). Mendez \\& Minnitti (2000) claimed that the faint blue objects found in the HDF--North and HDF--South are Galactic stars based on the observed number of blue sources and extragalactic sources in the two fields. Independent proper motion measurements for five of these faint blue sources by Ibata et al. (1999) suggested that they are cool halo white dwarfs which could account for the entire missing mass in the solar neighborhood. Third epoch data on these five objects, however, did not show any significant proper motion (R. Ibata, private communication; Richer 2001).  We use the original Hubble Deep Field -- North data and images of the same field taken 7 years later for the Great Observatories Origin Deep Survey (GOODS) to measure proper motions of the point sources analyzed by Ibata et al. (1999) and Mendez \\& Minniti (2000). ",
        "conclusions": "The nature of the faint blue objects in the Hubble Deep Field may be crucial to understanding the contribution of low luminosity halo white dwarfs to micro-lensing events and the dark matter content of the Galaxy. Apparent proper motions for 5 faint blue objects (Ibata et al. 1999) was enough to explain the entire missing mass in the halo of the Milky Way. Mendez \\& Minniti (2000) claimed that the faint blue objects are white dwarf stars located at heliocentric distances of up to 2 kpc and belong to the Galactic halo. They found a local halo white dwarf mass density of 4.64 $\\times$ 10$^{-3}$ \\msun\\ pc$^{-3}$, which would account for about 30--50$\\%$ of the dark matter in the Galaxy.  With the advantage of a 7--year baseline, we are able to place better limits on the proper motion measurements of the faint blue objects. Using the proper motion information, we also derived distances and tangential velocities for these objects. Figure 7 shows the tangential velocities and distances for objects brighter than $V\\approx27$ assuming that they are main sequence stars or DA white dwarfs. All of the main sequence stars exhibit halo kinematics and distances, whereas all of the likely white dwarfs exhibit disk kinematics and distances.  Following Gilmore, King, \\& van der Kruit (1989; see also von Hippel \\& Bothun 1990) we use the analytical form of the density profile for the thin disk and thick disk \\begin{equation} \\frac{\\nu_{0}(z)}{\\nu_{0}(0)}= {0.96}\\ {e^{-z/250 pc}} + {0.04}\\ {e^{-z/1000 pc}} \\end{equation} with a local normalization of 0.11 \\msun\\ pc$^{-3}$ (Pham 1997). We use the form \\begin{equation} \\nu_{halo}(r) \\propto \\frac{exp[-7.669\\ (R/R_{e})^{(1/4)}]}{(R/R_{e})^{(7/8)}} \\end{equation} for the halo (Young 1976), where $R$ is the distance from the Galactic center, and R$_{e}$ is the scale factor. $R$ is related to the distance $r$ from the observer to a star by \\begin{equation} R^{2} = R_{0}^{2} + r^{2} - 2 r R_{0}\\ {\\rm cos}b\\ {\\rm cos}l \\end{equation} with $R_{0}$ the solar Galactocentric distance, and {\\it b} and {\\it l} the Galactic coordinates for the HDF--North. We use R$_{0}=$7.8 kpc (Gilmore, King, \\& van der Kruit 1989), R$_{e}=$2.7 kpc (de Vaucouleurs \\& Pence 1978) and a local normalization for the halo of (1/800) $\\times$ 0.11 \\msun\\ pc$^{-3}$ (Chen et al. 2001; Gilmore, King, \\& van der Kruit 1989). Using equations 2, 3, and 4, we calculated the expected number of stars in the HDF. We expect to find 2 thin disk, 3 thick disk, and 11 halo objects in the HDF--North.  We have also used Reid \\& Majewski (1993) star count models to predict the number of stars in the HDF--North. We found that 2 thin disk, 4 thick disk, and 14 halo objects are expected in the HDF--North. Both simple analytical models and more sophisticated star count models, when extrapolated to the photometric depth of the HDF, predict similar number of stars (16--20) in the HDF--North.  There are 14 stars brighter than $V=27$ and 17 objects fainter than $V=27$ classified as stars by SExtractor. The observed number of stars and the predictions of star count models are in good agreement for $V\\la27$ (see also Mendez et al. 1996 and Mendez \\& Minniti 2000). On the other hand, there seems to be an excess of point sources in the Hubble Deep Field -- North for $V\\ga27$. Unfortunately, SExtractor classification cannot be trusted at these magnitudes. Furthermore, we did not detect significant proper motion for all but two of these objects. The two faint, possibly moving objects, HDF774 and HDF1816, may be halo white dwarfs. One of the problems with any analysis using these objects is that the observations are beyond the completeness limit, and any calculation based on them is subject to a significant completeness correction. The rest of the objects fainter than $V=27$ are probably extragalactic objects (see section 3.2).  The five faint blue objects analyzed by Mendez \\& Minniti (2000) do not exhibit any significant proper motion; they are not halo white dwarfs. These objects do not account for the MACHO optical depth and are not the source of the Galactic dark matter. Their stellar nature is not confirmed either. The colors of HDF161 are consistent with our QSO simulations. The faint blue objects may be distant AGN.  Holberg et al. (2002) used a local sample of white dwarfs complete out to 13 pc, and found the local mass density of white dwarf stars to be 3.4 $\\pm$0.5 $\\times$ 10$^{-3}$ \\msun\\ pc$^{-3}$. Using this normalization factor in equations 2 and 3, we estimate the expected number of white dwarfs in the Hubble Deep Field. We expect to find 0.05 disk white dwarfs, 0.09 thick disk white dwarfs, and 0.33 halo white dwarfs in the Hubble Deep Field North. We have also used Reid \\& Majewski (1993) star count models to predict the number of white dwarfs in the HDF. The results are roughly consistent: 0.10 disk, 0.25 thick disk, and 0.5 halo white dwarfs are expected.  We have discovered two likely white dwarfs, HDF2234 and HDF3072, brighter than $V=27$ in the HDF--North. They are located at distances of $\\sim$ 500 pc and have tangential velocities $\\sim$30 km s$^{-1}$. Their kinematic properties are consistent with being thin disk or thick disk objects (see Table 3 and Figure 7). The expected number of thin disk + thick disk white dwarfs is found to be 0.14 -- 0.35. We have found 6 to 14 times more disk white dwarfs in the HDF--N than expected from the models. Assuming Poisson statistics, the probability of finding two white dwarfs is 1\\% if the expected number of white dwarfs is 0.14 and 4\\% if the expected number of white dwarfs is 0.35. The number of disk + thick disk white dwarfs may be substantially underestimated. Due to small number statistics, however, this statement is only a 2--3 $\\sigma$ result and it heavily depends on the fact that HDF2234 and HDF3072 are white dwarfs. Follow--up spectroscopy of these two objects is needed to confirm this result.  Mendez \\& Minniti (2000) have found 22 Galactic stars and 10 faint blue objects in the Hubble Deep Field -- South. A natural test to check the space density of disk and halo white dwarfs would be to obtain second epoch observations of the HDF--South to find high proper motion objects. Also, the HST/ACS Ultra--Deep Field observations of the Chandra Deep Field -- South will be useful to search for faint blue objects at fainter magnitudes and to improve the morphological classification of these objects at brighter magnitudes. The Ultra--Deep Field will be $\\sim$ 1.5 mag deeper than the HDF and HDF--South (Beckwith et al. 2003)."
    },
    "0403/astro-ph0403427_arXiv.txt": {
        "abstract": "The photon emissivity from the bremsstrahlung process $e^-e^-\\rightarrow e^-e^-\\gamma$ occuring in the electrosphere at the bare surface of a strange quark star is calculated. For surface temperatures $T< 10^{9}$K, the photon flux exceeds that of $e^+e^-$ pairs that are produced via the Schwinger mechanism in the presence of a strong electric field that binds electrons to the surface of the quark star. The average energy of photons emitted from the bremsstrahlung process can be 0.5 MeV or more, which is larger than that in $e^+e^-$ pair annihilation. The observation of this distinctive photon spectrum would constitute an unmistakable signature of a strange quark star and shed light on color superconductivity at stellar densities.  \\bigskip \\noindent PACS:  26.60.+c, 95.30.Cq, 97.60.Jd ",
        "introduction": "\\label{sec_intro}   Recently, it has been pointed out that thermal emission from the bare surface of a strange quark star, due to both photons and $e^+e^-$ pair production, can produce luminosities well above the Eddington limit ($\\sim 10^{38}~{\\rm erg~sec}^{-1}$) for extended periods of time, from about a day to decades, depending on the superconducting phase of quark matter~\\cite{Page02}. The spectrum of emitted photons is significantly different from that of a normal cooling neutron star ($30 <E/{\\rm keV} < 500$ instead of $0.1 < E/{\\rm keV} <2.5$).  This distinctive spectrum and temperature evolution, if observed, would constitute an unmistakable detection of a strange quark star and shed light on color superconductivity at stellar densities. The predicted characteristics are well within the capabilities of the INTEGRAL satellite~\\cite{Integral} launched toward the end of 2002. \\vskip 0.2cm A strange quark matter (SQM) star  differs from a normal neutron star (NS) in several important ways, which include  \\vskip 0.2cm  \\noindent (1) A SQM star is self-bound~\\cite{witten84,Chin79,Farhi84,Haensel86,AO88} while a NS requires gravitational forces for its binding.  As a consequence, the pressure vanishes at vanishing baryon density in a NS with a surface of normal matter.  The interior of the NS, however, may contain any or a combination of non-nucleonic matter such as hyperons, pion or kaon condensates, or quark matter~\\cite{LP2003,Alford2001}. In contrast, a self-bound SQM star is made entirely of quark matter up to its bare surface and the pressure vanishes at a finite but supra-nuclear baryon density. In the context of the MIT bag model with first order corrections due to gluon exchange, the baryon density $n_B$ at which the pressure vanishes is~\\cite{PBP90} \\beq \\label{nzero} n_B(P=0)=\\biggl(\\frac{4B}{3\\pi^{2/3}}\\biggr)^{3/4} \\biggl(1-\\frac{2\\alpha_c}{\\pi}\\biggr)^{1/4}, \\eeq where $B$ is the bag constant and $\\alpha_c=g_c^2/(4\\pi)$ is the quark-gluon coupling constant. For typical values of $B$ and $\\alpha_c$~\\cite{Farhi84}, the baryon density at vanishing pressure is about 2 to 3$n_0$, where $n_0=0.16$~fm$^{-3}$ is the saturation density of nuclear matter. The density in Eq.~(\\ref{nzero}) is not significantly affected by the finite strange quark mass~\\cite{PBP90} or by pairing gaps in the quark phase~\\cite{Reddy2002}.  \\vskip 0.1cm  \\noindent (2) The mass versus radius relations of NS and SQM stars differ significantly, although in the range of masses ($1<M/M_{\\odot}<2$) observed to date \\cite{TC99}, the calculated radii are similar ($R\\sim 10$ to 15 km). This makes it difficult to distinguish these two classes on the basis of their gross physical properties alone.  \\vskip 0.2cm  The light curves of NS and SQM stars, determined by surface photon emission, can however be very different under certain conditions. Central to the discussion of photon emission from a SQM star is the physical size and composition of its surface layer. If quark matter is in the Color-Flavor-Locked (CFL) superconducting phase (in which all three flavors take part in pairing) up to the surface, no electrons are required or admitted~\\cite{RW03,AR02,SRP03}.  Surface photon emission from the CFL phase has been found to saturate the blackbody limit at early times~\\cite{RVO}. In contrast, charge neutrality cannot be satisfied by three flavors of quarks alone in the two-flavor superconducting phase (2SC). In this case, an electron concentration $n_e/n_B$ of about $10^{-4}$ to $10^{-3}$ is required to achieve charge neutrality.  Normally, the bare surface of a quark star is an inefficient emitter of thermal radiation below temperatures $T\\sim 10^{11}$K because the plasma frequency of quark matter is high, of order 20 MeV~\\cite{Alcock86}.  However, in the 2SC phase, strong electric fields can arise because electrons in the surface layer are bound by electrostatic interaction to quark matter. This layer of electrons, which we will call the ``electrosphere'', is typically $10^{3}$ fm thick~\\cite{Alcock86} and can have radial electric fields whose magnitude ($\\sim 5\\times 10^{17}~{\\rm V~cm}^{-1}$) exceeds the critical value required for electron-positron pair production from the QED vacuum.  The critical value for the electric field was estimated by Schwinger to be \\cite{Schwinger} \\beq E_{\\rm cr}=\\frac{m_e^2c^3}{e\\hbar}\\simeq 1.3\\times 10^{16}~{\\rm V~cm}^{-1} \\quad, \\eeq where $m_e$ and $e$ are the electron's mass and charge, respectively.  For a homogeneous electric field $E \\gg E_{\\rm cr}$, the vacuum pair-creation rate per unit volume is given by \\cite{Schwinger} \\beq R_{\\pm}=\\frac{m_e^4c^5}{24\\pi\\hbar^4}\\biggl(\\frac{E}{E_{\\rm cr}}\\biggr)^2\\simeq 1.7\\times 10^{50}\\biggl(\\frac{E}{E_{\\rm cr}}\\biggr)^2~{\\rm cm}^{-3}~{\\rm s}^{-1} ~. \\eeq Usov~\\cite{Usov98} first exploited the possibility of thermal emission of hard photons from the annihilation of $e^+e^-$ pairs created by super-critical electric fields in degenerate matter and demonstrated that the associated photon emissivity is large enough to be observable.  Subsequently, Page and Usov~\\cite{Page02} have shown that the bare surface of a quark star under super-critical electric fields leads to a light curve that is significantly different from those produced by either quark matter surrounded by a layer of normal matter, or by surfaces of compact objects made entirely of normal matter. Their principal conclusion was that hard photon emission from $e^+e^-$ annihilation can dominate significantly over the blackbody spectrum at lower luminosities, and that the variation in the photon spectrum with luminosity could be used as an observational signature of a young bare strange quark star. Furthermore, calculations of light curves from quark matter in its superconducting phases (CFL or 2SC) have shown that it may be possible to set bounds on quark pairing gaps provided they are in the range 0.5 MeV to a few MeV \\cite{Page02}.  \\vskip 0.2cm  Our objectives in this paper are to  \\noindent (1) point out that, depending on the temperature, the bremsstrahlung process $e^-e^- \\rightarrow e^-e^-\\gamma$ in the electrosphere can lead to photon luminosities that are significantly larger than the $e^+e^-$ pair production luminosities in the presence of strong electric fields, \\\\ \\noindent (2) present a calculation of the corresponding emissivity including minimal screening effects, and investigate the role of the Landau-Pomeranchuk-Migdal (LPM) effect, and \\\\ \\noindent (3) compare the photon emissivity from the $e^-e^- \\rightarrow e^-e^-\\gamma$ bremsstrahlung process to other processes occurring in the thin electrosphere at the surface of the star. Examples of such processes include $e^+e^-$ pair production which occurs due to the presence of strong electric fields, and equilibrium and non-thermal bremsstrahlung radiation from quark-quark collisions in the uppermost layer of quark matter.  \\vskip 0.2cm  \\noindent We wish to emphasize that our discussion applies only to the case in which a charge neutralizing surface electron layer is present. This can be realized in the 2SC, the gapless CFL (gCFL) or crystalline color superconducting phases~\\cite{AR02,KR}, but not in the pure CFL phase.  \\vskip 0.2cm   This paper is organized as follows. Sec. II is devoted to a qualitative discussion of photon propagation in the electrosphere. Here the physical characteristics of the electrosphere, the photon mean free path, the photon dispersion relation, and the role of the Landau-Pomeranchuk-Migdal (LPM) effect that suppresses the emission of low-energy photons are addressed. The calculation of the photon emissivity from the bremsstrahlung process in degenerate matter (with and without modifications from the LPM effect) in the electrosphere, and a characterization of the non-thermal nature of the emitted photon spectrum are detailed in Sec. III.  This section also contains numerical results and discussion of the bremsstrahlung radiation flux, including a comparison with the flux of the $e^+e^-$ pair annihilation process and quark-related processes. Sec. IV contains our conclusions and outlook. Appendix A outlines the evaluation of an integral required in Sec. III.  \\vskip 0.2cm ",
        "conclusions": "We have calculated the photon emissivity from the $e^-e^- \\rightarrow e^-e^-\\gamma$ bremsstrahlung process occurring in the electrosphere of a bare strange quark star. For degenerate electrons and for low energy photons, unscreened magnetic interactions provide the largest contribution to the emissivity, while the contribution from screened electric interactions is comparatively negligible.  For temperatures in the range $10^8$K$<T<10^9$K, the emissivity from the bremsstrahlung process surpasses that of the $e^+e^-$ pair production process occurring in the presence of a strong electric field in the electrosphere. Within the soft photon approximation, the mean energy of the bremsstrahlung photons escaping the electrosphere is $\\sim 0.5$ MeV. This estimate may be slightly altered if the entire spectrum is sampled. Under strongly degenerate conditions, the contribution of high energy photons to the emissivity is expected to be small in comparison to that from low energy photons.  Photons from the bremsstrahlung process are non-thermal for sufficiently thin electrospheres and moderate electron chemical potentials, because such photons do not undergo sufficient number of scatterings in the electrosphere.  Inclusion of the Landau-Pomeranchuk-Migdal effect, which arises due to multiple scatterings of the emitting electron, reveals that even in the presence of a large suppression from this effect the bremsstrahlung process emerges as the dominant source of radiation from the star's surface for temperatures in the range $10^8$K$<T<10^9$K.  We may regard the LPM-inclusive bremstrahlung rates as a conservative estimate, since we have approximated photon exchange by a static Coulomb field.  If retardation effects were included, the suppression will be smaller leading to somwewhat higher rates. Both the ``bare'' and LPM-corrected rates take into account the dielectric suppression effect of the electrosphere. The consequence is that the LPM suppression is two orders of magnitude less than that in the case the photon is emitted into vacuum. There is an additional mitigating effect of about 15\\% on the LPM suppression due to the strong external electric field, which can bend the electron trajectory causing the formation time of the emitted photon wave packet to decrease.  This implies that the emission angles can be increased beyond $\\theta_{max}=m_e/E$. However, this only changes the emissivity by a pre-factor $\\sim 4/3$, and to logarithmic accuracy our principal results are valid even in the presence of strong electric fields. In view of this, our LPM-corrected rates should be treated as a lower bound.  An intriguing result of our calculations is that for $T\\leq 10^{-3}~\\mu_e(0)$, bremsstrahlung emission becomes comparable to blackbody emission. This indicates that bremstrahlung absorption has to be taken into account in estimating the optical thickness of the electrosphere (see the discussion following Eq.~(\\ref{Thomson})). This will lead to better estimates of the emissivity applicable in the temperature window excluded by Eq.~(\\ref{numemiss}). The remaining task is to perform calculations of the thermal evolution of strange quark matter stars including the non-thermal bremsstrahlung and thermal pair annihilation processes in order to provide baseline calculations of the surface luminosity versus age. Results of such calculations will be reported separately."
    },
    "0403/astro-ph0403561_arXiv.txt": {
        "abstract": "{ We revisit the possibility that an excess in the CMB power spectrum at small angular scales (CBI, ACBAR) can be due to galaxy clusters (or compact sources in general). We perform a Gaussian analysis of ACBAR-like simulated data based on wavelets. We show how models with a significant excess should show a clear non-Gaussian signal in the wavelet space. In particular, a value of the normalization $\\sigma_8 = 1$ would imply a highly significant skewness and kurtosis in the wavelet coefficients at scales around 3 arcmin. Models with a more moderate excess also show a non-Gaussian signal in the simulated data. We conclude that current data (ACBAR) should show this signature if the excess is to be due to the SZ effect. Otherwise, the reason for that excess should be explained by some systematic effect. The significance of the non-Gaussian signal depends on the cluster model but it grows with the surveyed area. Non-Gaussianity test performed on incoming data sets should reveal the presence of a cluster population even for models with moderate-low $\\sigma_8$'s. ",
        "introduction": "In the last couple of years there have been several results suggesting that there is a significant excess in the power spectrum of CMB experiments at small scales ($\\theta < 5$ arcmin or $\\ell > 2000$) (CBI, Pearson et al. 2003; ACBAR, Kuo et al. 2002). The best candidate for this excess (if confirmed) is the SZ effect signal in galaxy clusters. We expect the excess in power due to galaxy clusters to be observed at some point at scales smaller than $\\approx$ few arcmin (or $\\ell > 3000$) although the exact value is model dependent. However, the amplitude of the excess claimed by recent experiments is larger than what it is expected for the current most fashionable models. In fact, if that excess is confirmed to be due to galaxy clusters, it would require high values for the normalization of the matter power spectrum ($\\sigma_8 \\approx 1$) (e.g Bond et al. 2002). These high values of $\\sigma _8$ would contradict the values derived directly from observations of galaxy clusters (e.g Efstathiou et al 2002). Another contribution to the excess in power could be due to unresolved non-subtracted point sources. It is expected to be more important for experiments at low frequencies in the microwave band (like CBI) than at frequencies around 150 GHz (like ACBAR). Since a strong contribution from radio sources is expected in the former case and in the later the extragalactic point source contribution is at the lowest level along the microwave band, a more significant residual after the subtraction/estimation analysis can be present at centimeter rather than millimeter wavelengths. Moreover, the effect of clustering of point sources should be considered in the estimation of the residual (Toffolatti et al. in preparation).\\\\ An interesting aspect of the excess in power is that no compact sources (clusters or point sources) are clearly seen directly in the data, thus suggesting that the origin  of this excess may be due to a different nature other than the compact sources. One has only to realize that the power spectrum is an averaged quantity over the surveyed area and that the compact (bright) sources are not distributed over the whole area. As a consequence, many pixels will enter in the average with negligible values. On the other hand,  the CMB does distribute over the whole surveyed area. This implies that if the power spectrum of compact sources dominates the power spectrum of the CMB at small scales, then those sources should be seen in the map at small scales. Then, where are they ? The debate is still open and it is not clear whether the excess is due to compact sources or is just a systematic effect.\\\\ In this paper we will study the SZ contribution to the excess found in the power spectrum and its cosmological implications. In particular we will analyze the Gaussian deviations introduced in the wavelet coefficients. Since the CMB has not shown (up to date) any departures from Gaussianity at least at those small scales (for larger scales see Vielva et al. 2003) therefore any non-Gaussian detection would point out the presence of compact sources.  The analysis presented in this paper will be useful to confirm or reject the hypothesis that the excess is due to galaxy clusters in current CMB data. The analysis would be also a desirable consistency check with the near future incoming data sets. \\\\ The reader will find very interesting for instance the pioneering work of Aghanim \\& Forni (1999) and the more recent of Rubi\\~{n}o-Martin \\& Sunyaev (2003) where similar questions are considered. ",
        "conclusions": "We have presented a small (but interesting) excercise showing how current data (ACBAR CMB5 field) is expected to have non-Gaussian signals if the claimed excess in the power spectrum is fully due to the SZ effect. With this excercise we want to send a brief but clear note to the community. Current data should be checked for non-Gaussian signals at small scales before any excess in the power spectrum is claimed as due to the SZ effect.  The potentiality of this approach is based on the fact that, up to date, no intrinsic deviations from Gaussianity have been found in the CMB at these small scales. Any deviations would indicate the presence of compact sources which are expected to be the dominant foreground at arcmin scales.\\\\ We have presented the result of the level of non-Gaussianity we should expect in ACBAR-like data if a significant excess in power is due to galaxy clusters. Given the uncertainties in the estimation of the power spectrum by ACBAR and CBI, we explore two models which are in the lower and upper limit of these uncertainties and we find that for the model in the upper limit we should expect a clear non-Gaussian signal at MHW scales around 3 arcmin (the Gaussian and A-model distributions in figures \\ref{fig_1} and \\ref{fig_2} overlap by less than 1 \\%). In the other case (lower limit B-model), we should expect a non-Gaussian signal in $\\approx 75 $\\% of the cases if we look at the kurtosis and in $\\approx 80 $\\% if we look at the skewness (both at $> 99$ \\% confidence limit). By performing this analysis on the ACBAR data (for example), one could confirm or reject several galaxy cluster models.  \\\\ Although we have based our simulations on galaxy clusters, a similar argument could be made if the excess is due to point sources. A way to distinguish between both would be the sign of the skewness of the MHW coefficients: a positive sign would point out to point sources whereas a negative one would imply the presence of galaxy clusters (see Rubi\\~{n}o-Martin \\& Sunyaev 2003). For experiments where the level of non-Gaussianity is expected to be small, a more efficient manner to discriminate between the CMB and galaxy clusters is by defining a Fisher discriminant which contains information about both the skewness and kurtosis and also at several scales (Mart\\'\\i nez-Gonz\\'alez et al. 2002). These techniques will prove to be very useful with the near future data.  We would like to conclude by being provocative by insisting that actual data should be checked for non-Gaussian signals before suggesting that the excess is due to the SZ effect. If the data is found to be compatible with Gaussianity, then the systematics may be the reason for that excess. For instance, the power spectrum may have been overestimated at small scales if the covariance matrix of the instrumental noise is slightly undetermined at those small scales. Although we are not the first suggesting that non-Gaussianity studies are useful to detect SZ signatures (See Aghanim \\& Forni 1999 or  Rubi\\~{n}o-Martin \\& Sunyaev 2003), we are the first in predicting that if a significant excess claimed by recent experiments is due to the SZ effect, in general, non-Gaussian signals should be found with a high significance in at least one of those data sets, the CMB5 field of ACBAR. The reader may consider our approach too simplistic but we think is important to highlight our main point with simple arguments, current data should be checked for non-Gaussian signals."
    },
    "0403/astro-ph0403082_arXiv.txt": {
        "abstract": "We discuss the propagation of electromagnetic waves on a rectangular lattice of polarizable point dipoles. For wavelengths long compared to the lattice spacing, we obtain the dispersion relation in terms of the lattice spacing and the dipole polarizabilities. We also obtain the dipole polarizabilities required for the lattice to have the same dispersion relation as a continuum medium of given refractive index $m$; our result differs from previous work by Draine \\& Goodman (1993). Our new prescription can be used to assign dipole polarizabilities when the discrete dipole approximation is used to study scattering by finite targets.  Results are shown for selected cases. ",
        "introduction": "The discrete dipole approximation (DDA) provides a flexible and general method to calculate scattering and absorption of light by objects of arbitrary geometry \\citep{Drai88,DrFl94}. The approximation consists of replacing the continuum target of interest by an array of polarizable points, which acquire oscillating electric dipole moments in response to the electric field due to the incident wave plus all of the other dipoles. Broadly speaking, there are two criteria determining the accuracy of the approximation: \\begin{itemize} \\item the interdipole separation $d$ should be small compared to the wavelength of the radiation in the material ($|m|kd < 1$, where $m$ is the complex refractive index, and $k=\\omega/ c$ is the wavenumber {\\it in vacuo}); \\item the interdipole separation $d$ should be small enough to resolve structural dimensions in the target. \\end{itemize} With modern workstations, it is now feasible to carry out DDA calculations on targets containing up to $\\sim10^6$ dipoles \\citep{DrFl94}. In addition to calculation of scattering and absorption cross sections, the DDA has recently been applied to computation of forces and torques on illuminated particles \\citep{DrWe96,DrWe97}  If the dipoles are located on a cubic lattice, then in the limit where the interdipole separation $kd\\rightarrow0$, the familiar Clausius-Mossotti relation (see, e.g, \\citet{Ja62}) can be used to determine the choice of dipole polarizabilities required so that the dipole array will approximate a continuum target with dielectric constant $\\epsilon$. \\citet{Drai88} showed how this estimate for the dipole polarizabilities should be modified to include radiative reaction corrections, and \\citet{DrGo93} derived the $O[(kd)^2]$ corrections required so that an infinite cubic lattice would have the same dispersion relation as a continuum of given dielectric constant.  Nearly all DDA calculations to date have assumed the dipoles to be located on a cubic lattice. If, instead, a rectangular lattice is used, it will still be possible to apply FFT techniques to the discrete dipole approximation, in essentially the same way as has been done for a cubic lattice \\citep{GoDF91}. However, the ability to use different lattice constants in different directions might be useful in representing certain target geometries.  The objective of the present report is to obtain the dispersion relation for propagation of electromagnetic waves on a rectangular lattice of polarizable points. This dispersion relation will be expanded in powers of $(kd)$, where $d$ is the characteristic interdipole separation. For a lattice of specified dipole polarizabilities, this will allow us to determine the dispersion relation for electromagnetic waves propagating on the lattice when $(kd)^2\\ll 1$. Alternatively, if we require the lattice to propagate waves with a particular dispersion relation, inversion of the lattice dispersion relation will provide a prescription for assigning dipole polarizabilities when seeking to approximate a continuum material with a rectangular lattice of polarizable points. ",
        "conclusions": "The principal results of this paper are as follows:  \\begin{enumerate} \\item We have derived the exact mode equation (\\ref{eq:mode_equation}) for propagation of electromagnetic waves through a rectangular lattice of polarizable points.  \\item If the electric polarizability tensor $\\alpha$ for the lattice points is specified, then we can solve the mode equation to obtain the dispersion relation for waves propagating on the lattice. Alternatively, if we wish the lattice to have the same dispersion relation as a continuum medium of complex refractive index $m$, then we can solve the mode equation for the appropriate polarizability tensor.  \\item In the case of a cubic lattice, our derived ``lattice dispersion relation'' polarizability tensor (\\ref{eq:alpha_cubic}) is diagonal but anisotropic, and differs from the isotropic polarizabilities obtained by DG93.  \\item The new polarizability prescription has been tested in DDA scattering calculations for spheres. The resulting cross sections have accuracies comparable to those obtained with the DG93 prescription. \\end{enumerate}"
    },
    "0403/astro-ph0403207_arXiv.txt": {
        "abstract": "{ We report on a search for low-energy neutrino (antineutrino) bursts in correlation with the 8 time coincident events observed by the gravitational waves detectors EXPLORER and NAUTILUS (GWD) during the year 2001.  The search, conducted with the LVD detector (INFN Gran Sasso National Laboratory, Italy), has considered several neutrino reactions, corresponding to different neutrino species, and a wide range of time intervals around the (GWD) observed events. No evidence for statistically significant correlated signals in LVD has been found.  Assuming two different origins for neutrino emission, the cooling of a neutron star from a core-collapse supernova or from coalescing neutron stars and the accretion of shocked matter, and taking into account neutrino oscillations, we derive limits to the total energy emitted in neutrinos and to the amount of accreting mass, respectively. ",
        "introduction": "The analysis of the data collected in coincidence by the gravitational wave bar detectors EXPLORER and NAUTILUS during the year 2001 (Astone et al., 2002) shows an excess (8 events against 2.6 expected from the background) when the two detectors are favorably oriented with respect to the Galactic Disc. Moreover, this result comes from the present day most sensitive experiments for the detection of gravitational wave bursts and a search for neutrino bursts in correlation with the 8 GWD events is, therefore, appropriate.  A few astrophysical transient sources are indeed expected to produce associated bursts of neutrinos and gravitational waves. It is well known that most of the energy ($99\\%$) released in the gravitational core collapse of a massive star is carried away by neutrino originated both from the matter accretion in the shock and from the cooling of the proto-neutron star (see for example Burrows et al., 1992). Depending on the collapse dynamics, some fraction of the total energy is emitted in GW ( Thorne 1988, Muller 1997 ), asymmetric supernovae in our Galaxy being the best candidate sources for GW bar detectors. Two coalescing neutron stars would  also constitute a source for both neutrinos and gravitational waves. From the point of view of GW emission, it is likely that the merging event would produce powerful gravitational wave bursts, and, even if the physics of the merger is not known, there are estimates that, for binary systems of large mass, coalescence waves are likely to be stronger than the inspiral ones. Some amount of the kinetic energy is converted in thermal energy so that the hot remnant would probably emits thermal neutrinos.\\\\ The search for a neutrino burst associated to the events detected by the GWD EXPLORER and NAUTILUS in 2001 has been performed with the LVD apparatus, operating in the INFN Gran Sasso National Laboratory (Italy) since 1992 with the main purpose of searching for neutrinos from gravitational stellar collapses within the whole Galaxy.  The paper is planned as follows: in Sect.2 we briefly describe the LVD detector, and we explain the selection of the data. In Sect.3 we present the results of the analysis: a time interval spanning from 12 $h$ preceding each of the 8 events up to 12 hours later, has been scanned, searching for any excess over the statistical fluctuation of the background. Further, a search for a $\\nu$ signal coincident in time with every event has been performed. We conclude in Sect.4, where we discuss the results of the search, taking into account $\\nu$ oscillations, and considering the two processes for $\\nu$ emission, i.e., cooling and accretion. Since we do not find any neutrino burst candidate associated with the 8 GWD mentioned events, in the following assumed scenarios we derive upper limits: \\begin{itemize} \\item on the neutrino flux, without reference to any particular source, \\item  on the total amount of energy emitted in neutrinos, in the cooling case, \\item on the accreting mass, in the accretion case. \\end{itemize} ",
        "conclusions": "We have conducted a search for low-energy antineutrino (neutrino) bursts with the LVD detector in coincidence with the 8 event excess found by the gravitational waves detectors EXPLORER and NAUTILUS during the year 2001.\\\\ We have found no evidence for any statistically relevant signal in LVD, in three different reaction channels (inverse beta decay, charged current and neutral current interactions with $^{12}C$) corresponding to different neutrino species, over a wide range of time durations, for any of the 8 events. Consequently, we have derived $90 \\%$ fluence upper limits to antineutrino and neutrino emission from an average GW event, at different energies in the range of sensitivity of the LVD detector.\\\\ We have then related the result of the search with two possible simplified models for neutrino emission, i.e., ``cooling'' and ``accretion'', deriving limits, on the one side, to the total energy emitted in neutrinos at the source, and, on the other, to the amount of accreting mass. Assuming a source distance $d=10$ kpc, possible candidates as new-born and colliding neutron stars have been excluded by this analysis. This makes even more challenging and interesting the search for a likely astrophysical source for the reported GWD events.\\\\[1cm] {\\it Acknowledgments}. The authors are grateful to the director and the staff of the National Gran Sasso Laboratory for their constant and valuable support. The authors thank precious comments by Francesco Vissani.\\\\[1cm]  {\\bf Appendix A. Limits calculation: inverse $\\beta$ decay interactions.}\\\\[0.2cm] The number of $\\bar{\\nu}_e$ interactions due to inverse beta decay in a detector is given by: \\begin{center} \\begin{equation} N_{ev}^{IBD}=M \\cdot N_{p} \\cdot \\int_{Q}^{\\infty} \\frac{dN_{\\bar{\\nu}_e}}{dE_{\\bar{\\nu}_e}} \\cdot \\sigma(E_{\\bar{\\nu}_e}) \\cdot \\varepsilon_n \\cdot \\varepsilon(E_d,E_{th}) dE_{\\bar{\\nu}_e} \\label{Nevcoolp} \\end{equation} \\end{center} where: $M$ is the detector active mass in ton; $N_p=9.36\\cdot10^{28}$ is the number of free protons per ton; $\\varepsilon_n$ is the neutron detection efficiency; $\\varepsilon(E_d,E_{th})$ is the $e^+$ detection efficiency; $E_d=E_{\\bar{\\nu}_e}-Q+2m_ec^2$, with $Q=M_n+m_e-M_p=1.8$ MeV is the positron detectable energy; $E_{th}$ is the detector energy threshold; $\\sigma(E_{\\bar{\\nu}_e})$ is the cross section (Vogel and Beacom, 1999; Strumia and Vissani, 2003); $\\frac{dN_{\\bar\\nu_e}}{dE_{\\bar\\nu_e}}$ is the antineutrino energy spectrum.\\\\ In the case of cooling process it is: \\begin{center} \\begin{equation} \\frac{dN_{\\bar\\nu_e}}{dE_{\\bar\\nu_e}}= \\frac{E_{B}}{4\\pi d^2} \\cdot \\frac{120}{7\\pi^4} \\cdot F_{\\bar\\nu_e} \\end{equation} \\end{center} where $d=10 Kpc$ is the assumed source distance;\\\\ $E_{B} = \\sum_{i}f_{\\nu_i}E_{B}$ is the total energy emitted in neutrinos;\\\\ $F_{\\nu}$ is the term accounting for different $\\nu$ oscillation scenarios\\footnote{NC data are not affected by oscillations. However, the limits from IBD data stay almost the same even assuming that MSW oscillations are completely absent. In fact, on accounting for vacuum oscillations we get $P_{\\bar{e}\\bar{e}}=1-sin^2 2\\theta_{12}/2 \\sim 0.6$ in all scenarios; this is practically the same value implied by MSW oscillations in the scenarios $NH$ and $IH_{non\\ ad.}$, $P_{\\bar{e}\\bar{e}}=cos^2\\theta_{12}\\sim 0.7$.}: \\begin{itemize} \\item for $NH_{ad}$; $NH_{non-ad}$; $IH_{non-ad}$: \\begin{center} \\begin{equation} F_{\\bar{\\nu}_e}= \\frac{f_{\\nu_e}}{T_{\\bar{\\nu}_e}^4}|U_{e1}|^2\\frac{E^2}{1+e^{E/ T_{\\bar{\\nu}_e}}} + \\frac{f_{\\nu_x}}{T_{\\bar{\\nu}_x}^4}|U_{e2}|^2\\frac{E^2}{1+e^{E/ T_{\\bar{\\nu}_x}}} \\label{osc1ae} \\end{equation} \\end{center} \\item for $IH_{ad}$: \\begin{center} \\begin{equation} F_{\\bar{\\nu}_e}= \\frac{f_{\\nu_e}}{T_{\\bar{\\nu}_e}^4}|U_{e3}|^2\\frac{E^2}{1+e^{E/ T_{\\bar{\\nu}_e}}} + \\frac{f_{\\nu_x}}{T_{\\bar{\\nu}_x}^4}(1-|U_{e3}|^2)\\frac{E^2}{1+e^{E/ T_{\\bar{\\nu}_x}}} \\label{osc2ae} \\end{equation} \\end{center} \\end{itemize} with: $|U_{e1}|^2 \\approx cos^2\\theta_{12}=0.67$, $|U_{e2}|^2 \\approx sin^2\\theta_{12}=0.33$, $|U_{e3}|^2 \\geq 10^{-4}$ for the adiabatic case and $|Ue3|^2 \\leq 10^{-6}$ for the non adiabatic one (M. Apollonio et al., 1999) (J.Bahcall and C.Pena-Garay, 2003).\\\\  In the case of mass accretion process it is (Loredo and Lamb, 2002): \\begin{center} \\begin{equation} \\frac{dN_{\\bar\\nu_e}}{dE_{\\bar\\nu_e}}= \\frac{1}{4\\pi d^2} \\cdot A_a \\cdot Y_n \\cdot M_{hot} \\cdot F'_{\\bar\\nu_e} \\end{equation} \\end{center} where:\\\\ $M_{hot}$ is the mass of hot emitting material;\\\\ $Y_n$ is the neutron fraction;\\\\ $A_a=\\frac{1+3g^2_A}{8}\\frac{\\sigma_{0}c}{m_n(m_ec^2)^2}\\frac{8\\pi}{(hc)^3}$, with $g_A=1.254$, $\\sigma_{0}=1.7 \\cdot 10^{44} cm^2$;\\\\ and with respect to $F'_{\\nu}$: \\begin{itemize} \\item for $NH_{ad}$; $NH_{non-ad}$; $IH_{non-ad}$: \\begin{equation} F'_{\\bar{\\nu}_e}= |U_{e1}|^2\\frac{E^4}{1+e^{E/ T_{\\bar{\\nu}_e}}} \\label{osc3ae} \\end{equation} \\item for $IH_{ad}$: \\begin{equation} F'_{\\bar{\\nu}_e}=|U_{e3}|^2\\frac{E^4}{1+e^{E/ T_{\\bar{\\nu}_e}}} \\simeq 0 \\label{osc4ae} \\end{equation} \\end{itemize} \\vspace{0.2cm}  {\\bf Appendix B. Limits calculation: neutral current interactions}\\\\[0.2cm] The number of interactions in the detector due to the neutral current is given by: \\begin{center} \\begin{equation} N_{ev}^{NC}=M \\cdot N_{C} \\cdot \\varepsilon_C \\cdot \\int_{15.1\\ {\\rm{MeV}}}^{\\infty}[\\frac{dN_{\\bar{\\nu}_i}}{dE_{\\bar{\\nu}_i}} \\sigma(E_{\\bar{\\nu}_i}) + \\frac{dN_{\\nu_i}}{dE_{\\nu_i}} \\sigma(E_{\\nu_i})] dE \\label{Nevcoolcn} \\end{equation} \\end{center} where: $N_{C}=4.24\\cdot10^{28}$ is the number of $^{12}C$ nuclei per ton; $\\varepsilon_C$ is the detector efficiency for $15.1$ MeV gamma; $\\sigma (E_{\\nu})$ is the cross section (M. Fukugita et al., 1988). \\\\ The neutrino energy spectrum in the case of cooling process is: \\begin{center} \\begin{equation} \\frac{dN_{\\bar\\nu_i}}{dE_{\\bar\\nu_i}}=\\frac{dN_{\\nu_i}}{dE_{\\nu_i}}= \\frac{E_{B}}{4\\pi d^2} \\cdot \\frac{120}{7\\pi^4} \\cdot F_{i} \\end{equation} \\end{center} with $F_i=   \\frac{f_{\\nu_e}}{T_{{\\nu}_e}^4}\\frac{E^2}{1+e^{E/ T_{{\\nu}_e}}} + 2 \\cdot \\frac{f_{\\nu_x}}{T_{{\\nu}_x}^4}\\frac{E^2}{1+e^{E/ T_{{\\nu}_x}}}$\\\\ while, for the mass accretion case, we considered all the events as $\\bar\\nu_e$ and we used: \\begin{center} \\begin{equation} \\frac{dN_{\\bar\\nu_i}}{dE_{\\bar\\nu_i}}=\\frac{dN_{\\bar\\nu_e}}{dE_{\\bar\\nu_e}}= \\frac{1}{4\\pi d^2} \\cdot A_a \\cdot Yn \\cdot M_{hot} \\cdot \\frac{E^4}{1+e^{E/ T_{\\bar \\nu_e}}} \\end{equation} \\end{center}"
    },
    "0403/astro-ph0403031_arXiv.txt": {
        "abstract": "In this work we investigate the evolution of matter density perturbations for two different quintessence models. One of them is based on the Einstein theory of gravity, while the other is based on the Brans-Dicke scalar tensor theory. We show that it is possible to constrain the parameter space of the models using the determinations for the growth rate of perturbations derived from data of the 2-degree Field Galaxy Redshift Survey. ",
        "introduction": "In the past few years, it has become apparent that the energy budget of our universe is dominated by an unknown component called \"dark energy\". The WMAP table of \"best\" cosmological parameters \\cite{lambda}, for instance, gives a $0.73\\pm0.04$ abundance for it, and a value of its equation of state $\\omega<-0.78$. In order to relieve some problems of the popular $\\Lambda CDM$ model (like the fine tuning issue), a dynamical $\\Lambda$-term has been proposed as representative of the dark energy. It's more popular version is a slowly rolling scalar field named quintessence. Many alternative cosmological models have been proposed, and indeed it is a challenge the work of ruling out all the \"incorrect ones\" on observational grounds. For instance many different potentials for these self interacting scalar fields (quintessence) have been proposed. However, it is obvious the importance of the observational exploration of the proposed cosmological models, and in this paper we give a further step in this direction. In \\cite{Isra1} two cosmological models are proposed. In both cases two fluids fill the universe: a background fluid of ordinary matter and a self-interacting scalar-field fluid accounting for the dark energy, both in Einstein's theory and in Brans-Dicke gravity. A linear relationship between the Hubble expansion parameter and the time derivative of the scalar field is used to derive exact cosmological attractor-like solutions. And a priori assumptions about the functional form of the potential or the scale factor behavior are not necessary. All these features render two very interesting models that motivated us to proceed with the observational check of them. In the original paper, the authors made a first check with some observational facts, related with Cosmic Microwave Background (CMB), nucleosynthesis, structure formation (restriction on quintessence density during galaxy formation epoch) and type Ia supernovae (SN Ia). Now we proceed the observational checking considering another aspect of structure formation: the galaxy motions and clustering, i. e., the evolution of density perturbations in the Universe. In the past few years, observations of the large scale structure of the Universe have improved greatly. The development of fiber-fed spectrographs that can simultaneously measure spectra of hundreds of galaxies has provided large redshift surveys such as the 2-degree Field Galaxy Redshift Survey (2dFGRS) and the Sloan Digital Sky Survey (SDSS). In particular, the Anglo-Autralian Telescope of the 2dFGRS has obtained the redshift of a quarter million galaxies. This collaboration has produced abundant data and technical papers \\cite{2dFGRS} about galaxy motions and clustering, and we will refer to some of this, in particular their velocity/density comparisons.  The paper is organised as follows: in section II we outline the main characteristics of the models, in section III the main aspects of velocity/density comparisons are exposed and the equation for the growth of perturbations is solved, in section IV the observational check is presented and interpreted, while in section V conclusions are drawn. Finally, references are supplied. ",
        "conclusions": "In the models presented here we found that the study of perturbation growth is a good tool to constraint the parameter space. Research on the origins and evolution of the large-scale of the universe is one of the hottest topics in cosmology. In this work, we have used the relation between the peculiar velocity field of the galaxies, the growth rate of perturbations and the density bias in galaxy formation to make another step in the observational check of two quintessence models, resulting in a further constrain on the parameter space of the models. We plan in the oncoming future proceed this works using another cosmological probes, like CMB, for instance."
    },
    "0403/astro-ph0403341_arXiv.txt": {
        "abstract": "Using the method of multiple scales, one can derive an analytic solution that describes the behaviour of weakly coupled, non-linear oscillations in  nearly Keplerian discs around neutron stars or  black holes close to the 3:2 orbital epicyclic resonance.  The solution obtained agrees with the previous numerical simulation. Such result may be relevant to the kilohertz quasi-periodic variability in X-ray fluxes observed in many Galactic black hole and neutron star sources. With a particular choice of  tunable parameters, the solution fits accurately the observational data for Sco X-1. ",
        "introduction": "Many Galactic black hole and neutron star sources in low  mass X-ray binaries show both chaotic and quasi-periodic variability in their observed X-ray fluxes. Some of the quasi-periodic oscillations (QPOs) are in the kHz range and often come in pairs $(\\nu_{\\rm upp}, \\nu_{\\rm down})$ of {\\it twin peaks} in the Fourier power spectra (e.g., van der Klis, 2000). In all four black hole sources with twin peak kHz QPOs pairs, $\\nu_{\\rm upp}/\\nu_{\\rm down} = 3/2\\,$ (McClintock \\& Remillard, 2003). For  neutron stars sources the ratio of twin-peak frequencies is mostly  close to $3/2$ too (Abramowicz, Bulik, Bursa, Klu\\'zniak, 2003). Based on these and other properties of QPOs Klu\\'zniak \\& Abramowicz (2000) concluded that twin peak kHz QPOs are due to a resonance in the accretion disc oscillation modes and they noticed that the specific $3:2$ ratio could be a consequence of strong gravity.  According to the standard Shakura-Sunyaev accretion disc model, matter spirals down into the central black hole along stream lines that are located almost on the equatorial plane $\\theta = \\theta_0 = \\pi/2$, and that locally differ only slightly from a family of concentric circles $r = r_0 = const$. The small deviations, $\\delta r = r - r_0$, $\\delta \\theta = \\theta - \\theta_0$ are governed, with accuracy to linear terms, by   \\begin{equation} \\label{Equation1} \\delta \\ddot r + \\omega_r^2 \\,\\delta r = \\delta a_r , ~~~~ \\delta \\ddot \\theta + \\omega_{\\theta}^2\\,\\delta \\theta = \\delta a_{\\theta} . \\end{equation}  \\noindent Here the dot denotes the time derivative. For purely Keplerian (free) motion $\\delta a_r = 0$, $\\delta a_{\\theta} = 0$ and the above equations describe two uncoupled harmonic oscillators with the eigenfrequencies $\\omega_{\\theta}$, $\\omega_r$.  \\noindent We shall start with an argument which appeals to physical intuition and which shows that the resonance now discussed is a very natural, indeed necessary, consequence of strong gravity. Assume that in thin discs, random fluctuations have $\\delta r \\gg \\delta \\theta.$\\footnote{In actual calculations this additional assumption is not made.}  Thus, $\\delta r \\delta \\theta$ is a first order term in $\\delta \\theta$ and should be included in the first order equation for vertical oscillations (\\ref{Equation1}). The equation now takes the form,   \\begin{equation} \\label{Equation5} \\delta \\ddot \\theta + \\omega_{\\theta}^2\\left [ 1 + h\\,\\delta r \\right ] \\delta \\theta = \\delta a_{\\theta}, \\end{equation}  \\noindent where $h$ is a known constant. The first order equation for $\\delta r$ has the solution $\\delta r = A_0 \\cos (\\omega_r \\,t)$. Inserting this in (\\ref{Equation5}) together with $\\delta a_{\\theta} = 0$, one arrives to the Mathieu equation ($A_0$ is absorbed in $h$),   \\begin{equation} \\delta \\ddot \\theta + \\omega_{\\theta}^2\\left [ 1 + h \\,\\cos (\\omega_r \\,t)\\right ] \\delta \\theta = 0, \\end{equation}  \\noindent that describes the {\\it parametric resonance}. From the theory of the Mathieu equation one knows that when   \\begin{equation} {\\omega_r \\over \\omega_{\\theta}} = {\\nu_r \\over \\nu_{\\theta}} = {2 \\over n}, ~~~~n =1,\\,2, \\,3 ..., \\end{equation}  \\noindent the parametric  resonance is excited (Landau \\& Lifshitz, 1976). The resonance is strongests for the smallest possible value of $n$. Because near black holes and neutron stars $\\nu_r < \\nu_{\\theta}$ (see Figure \\ref{fig1}), the smallest possible value for resonance is $n = 3$, which means that $2\\,\\nu_{\\theta} = 3\\,\\nu_r$. This is an example of a situation in which parametric resonance works in a thin accretion disc and it intuitively explains the observed 3:2 ratio (Klu\\'zniak and Abramowicz, 2002), because, obviously,  \\begin{equation} \\nu_{\\rm upp} = \\nu_{\\theta},~~~~\\nu_{\\rm down} = \\nu_r. \\end{equation}  \\noindent Of course, in real discs neither $\\delta r = A_0 \\cos (\\omega_r \\,t)$, nor $\\delta a_{\\theta} = 0$ exactly; moreover in realistic situations, for purely geodesic motion ($\\delta a_{\\theta}=\\delta a_r=0$) the system does not show increasing amplitudes (the higher terms prevent this from occurring).\\\\ However, one may expect that because these equations are approximately obeyed in thin discs, the parametric resonance will indeed be excited. Such a resonance  was found in numerical simulations of oscillations in a nearly Keplerian accretion disc by Abramowicz et al. (2003). \\begin{figure}[!tb] \\begin{center} \\FigureFile(0.48\\textwidth,0.48\\textwidth){freq.eps} \\end{center} \\caption{Epicyclic  frequencies in Schwarzschild's metric: meridional (solid line) and radial one (dashed line). The radius is in units of $r_G$. The vertical lines indicates the radii ($r_{3:2}$,$r_{5:3}$,$r_{4:2}$) where the ratio between the two frequencies is $3:2$ , $5:3$ and $4:2$ (from right to left).} \\label{fig1} \\end{figure} ",
        "conclusions": "The Taylor expansion of the relativistic geodesics to the third order leads to two coupled harmonic oscillators: the purely geodesics motion is stable, while when it is perturbed there are cases for which parametric resonance may occur.\\\\ Using the method of multiple scales one can highlight that, owing to the curvature (through the way it determines the effective potential), the first allowed resonance between the radial an the vertical epicyclic frequencies  is the $3:2$, in agreement with the numerical analysis.\\\\ Moreover the deviation of the slope from $3:2$ is easily explained as a property of such a non-linear resonance\\\\ The analysis of this toy-model reinforces the theory that indeed the observed pairs of QPOs may be due to parametric resonance, and finally to the strong gravitational field alone.  \\medskip The work reported here is a part of my 2003 laurea thesis {\\it Risonanza parametrica di dinamiche quasi geodetiche con campi gravitationali forti} at the Physics Department, University of Trieste. I thank my external thesis supervisor ({\\it correlatore}), Marek Abramowicz, for suggesting the subject to me, the  many discussions and his constant support . I also thank my local supervisor ({\\it relatore}), Fabio Benatti, for his advice and help and Wlodek Klu\\'zniak for the help in revising this paper. I was working on the problem at the University of Trieste and several other institutions: at SISSA (Trieste, Italy), Chalmers University (G\\\"oteborg, Sweden), UKAFF (Leicester University, England), and the Cargese 2003 Spring School on Black Holes (Cargese, France). I thank all these institutions for hospitality. This work have been supported by the  EU grant in Leicester,by MA's VR Swedish goverment grant and by the Polish KBN Grant 2P03D01424.   \\appendix"
    },
    "0403/astro-ph0403388_arXiv.txt": {
        "abstract": "We explore theoretical models of the ionization ratios of the Li-like absorbers \\ion N5, \\ion O6, and \\ion C4, in the Galactic halo.  These ions are believed to form in nonequilibrium processes such as shocks, evaporative interfaces, or rapidly cooling gas, all of which trace the dynamics of the interstellar medium.  As a useful new diagnostic, we focus on velocity-resolved signatures of several common physical structures: (1) a cooling Galactic fountain flow that rises, cools, and recombines as it returns to the disk; (2) shocks moving toward the observer; (3) a conductive interface with the observer located in the hotter gas. This last geometry occurs with the solar system inside a hot bubble, or when one looks out through the fragmenting top shell of our local bubble blown into the halo as part of the Galactic fountain.  In Paper~II, these models are compared to ionization-ratio data from FUSE and {\\it Hubble Space Telescope}. ",
        "introduction": "The nature and dynamics of the interstellar medium (ISM) of galaxies determines how the energy and matter released by stars are redistributed through the universe.  It is thus critical to understand the ISM, and in particular that of the Milky Way, which can be observed with greater sensitivity and resolution than other galaxies. The ISM in the disk of our galaxy consists of several phases, from dense and cold molecular gas to hot and fully ionized.  Stars heat and disperse the dense clouds in which they form, and that hot gas cools and recombines eventually to complete the cycle and form more stars. It is particularly interesting to consider the interface between this multiphase medium and intergalactic space, where hot gas is released from supernovae to several kiloparsecs in altitude, forming a hot diffuse medium first proposed by \\cite{spitzer56} as the Galactic corona.  The Galactic corona or halo is almost certainly a dynamic object.  It is difficult to construct static halo models, because models supported by thermal pressure are thermally unstable.  If conduction is sufficiently important to stabilize small-scale instabilities, then the entire (nearly isothermal) halo is unstable to collapse or expulsion as a wind \\citep[see][for stability arguments]{bregman80,field65}.  Cosmic-ray supported static halos have been proposed \\citep[e.g.,][]{boulares90}, but there are considerable uncertainties as to how cosmic rays are confined by the Galaxy (the Galactic magnetic field topology in particular), and these models do not explain the high and intermediate velocity clouds.  Considering these things, \\citet{shapfield76} first proposed a ``Galactic fountain'' of supernova-heated gas rising buoyantly above the disk until it cools and falls back to the disk.  Smooth Galactic fountain models have been constructed by several authors. In particular, \\citet{bregman80} discussed the issues of radial flow in the Galactic gravitational field for a supersonic hot flow which reaches several kiloparsecs height. More recent models \\citep{houck90,breit93} consist of transsonic flows, which require cooler initial temperatures and only rise to 1--2~kpc. These models can help to explain the population of intermediate velocity halo clouds, believed to exist at those lower heights and velocities compared to the more distant high velocity clouds \\citep[e.g.,][]{hivcdist}.  The cooling layer certainly fragments into small clouds by Rayleigh-Taylor instability even if no other inhomogeneities exist \\citep[e.g.,][]{berry98}.  Superbubbles, worms, and shell-like structures have been observed in \\citep[e.g.,][]{heiles84} the upper disk/low halo (few hundred parsecs altitude), but it is difficult to ascertain whether the hot gas is being expelled into the halo.  A related issue is the filling factor of hot gas in the disk, which determines the rate at which hot gas can escape into the halo. Estimates for that filling factor and for the rate of reheating by halo supernovae are based on the evolution of supernova remnants in smooth media, but our current understanding of the ISM is increasingly inhomogeneous and dynamic, so those arguments are probably of limited value \\citep{kahn98}.  Recently, the combination of good observations, sophisticated numerical models, and sufficient comprehension of the Galactic fountain may be beginning to allow identification of fountain-like rising structures \\citep{avillez01}.  The dynamics of interstellar gas in general and specifically in the Galactic halo may be best understood by studying the hot phase of the ISM (coronal gas at millions of degrees), and gas at temperatures $\\sim$10\\ts{5}~K intermediate to the hot phase and cooler phases (10\\ts{4}~K).  Of the different interstellar gas phases, the coronal gas is most directly linked to the main sources of energy in the ISM, supernovae and stellar winds.  Slightly cooler gas is most closely linked to transient and dynamical processes.  This gas is typically short-lived because the cooling time is short at at 10\\ts{5}~K. The lithium-like ions of common metals, in particular \\ion O6, \\ion N5, and \\ion C4, are sensitive tracers of interstellar gas at several times 10\\ts{5}~K.  The resonance absorption lines of these ions are observable with ultraviolet spectrographs on the {\\it Far Ultraviolet Spectroscopic Explorer} (FUSE), and the {\\it Space Telescope Imaging Spectrograph} (STIS) and {\\it Goddard High Resolution Spectrograph} (GHRS) aboard the {\\it Hubble Space Telescope} (HST).  Section \\ref{hotmodels} describes previous models of high-ion column densities.  Section \\ref{interpretation} describes new models of the dynamical signatures of the Li-like ions and their interpretation. We summarize in \\S~4. ",
        "conclusions": "We have considered the diagnostic power of velocity-resolved column density ratios in understanding the Galactic halo.  Column density ratios of Li-like ions in the Galaxy are useful to diagnose the physical formation mechanism of the gas and to study the interstellar gas cycle, and a survey of these ions can reveal general trends. In Paper II, we present a survey of sightlines observed with FUSE and HST, in which the distribution of N(\\N) and N(\\O) in the halo does not appear to favor a dominant physical production mechanism.  Here, we have presented models of interfaces and cooling nonequilibrium gas, focusing on the velocity-resolved N(\\N)/N(\\O) signatures. In particular, we consider conductive interfaces, radiative shocks and supernova remnants, and cooling gas in a Galactic fountain.  Typical ion column density ratios are summarized in Table~\\ref{summary}, along with a typical observed slope observed in FUSE and HST data. One important type of hot/cool gas interface for which the ionization-velocity signature has not been discussed is the turbulent mixing layer \\citep{mixing}. The structure of these Li-like ion producing structures is uncertain. It is likely that there is a turbulent cascade of eddies which would wash out any velocity-ionization signature, but there may possibly be large Kelvin-Helmholtz rolls with some coherent velocity signature.  \\begin{deluxetable}{lr} \\tabletypesize{\\footnotesize} \\tablecaption{\\label{summary} Summary of models} \\tablecomments{Models predict a slope in N(\\N)/O(\\O) with line-of-sight velocity for observers looking through the structure.  Galactic fountain model includes \\citet{bbc} cooling } \\tablehead{ \\colhead{Model} & \\colhead{Slope of Ratio log[N(\\N)/N(\\O)]} } \\startdata \\parbox{3in}{Radiative shock} & $\\sim$-0.0015 dex~(km~s\\ts{-1})\\ts{-1} \\\\ \\parbox{3in}{Mature SNR shell \\citep{shelton98}} & $\\gtrsim$-0.02 dex~(km~s\\ts{-1})\\ts{-1} \\\\ \\parbox{3in}{\\citet{houck90} fountain} & $\\sim$+0.002 dex~(km~s\\ts{-1})\\ts{-1} \\\\ \\tableline Indicative observed slope & -0.0032+/-0.0022(r)+/-0.0014(s) dex~(km~s\\ts{-1})\\ts{-1} \\enddata \\end{deluxetable}  The observable velocity-ionization trends are weak, because even very strong trends are washed out by the large thermal width of the gas at different parts of the flow. These trends could be further complicated by thermal instabilities, which likely occur in cooling gas such as in a fountain flow, and can also affect shocks and supernova shells.  Fragementation of the gas into parcels with complex density and velocity structure, as seen in the time-dependent shock model of \\citet{sutherland} could complicate the signatures described here.  Additional confusion can result when long sightlines pass through multiple structures.  As we show in Paper II, the dispersion of N(\\N)/N(\\O), both integrated and velocity-resolved, indicates that no single production scenario known to date can completely explain the Galactic halo.  To truly understand the physical production of Li-like ions in the halo, one needs to analyze gas in localized areas of physical space, rather than velocity space.  Absorption spectroscopy towards many halo stars with close angular separation and different distances could help to isolate gas at a specific altitude. Similarly, the gas above known superbubble shells or chimneys could be isolated.  These observations have a greater chance of distinguishing between models of hot gas production than observations along long lines of sight."
    },
    "0403/astro-ph0403177_arXiv.txt": {
        "abstract": "{ A sample of 37 Galactic, 10 LMC and 6 SMC cepheids is compiled for which individual metallicity estimates exist and $BVIK$ photometry in almost all cases. The Galactic cepheids all have an individual distance estimate available. For the MC objects different sources of photometry are combined to obtain improved periods and mean magnitudes. A multi-parameter Period-Luminosity relation is fitted to the data which also solves for the distance to the LMC and SMC.  When all three galaxies are considered, without metallicity effect, a significant quadratic term in $\\log P$ is found, as previously observed and also predicted in some theoretical calculations. For the present sample it is empirically determined that for $\\log P < 1.65$ linear $PL$-relations may be adopted, but this restricts the sample to only 4 LMC and 1 SMC cepheid.  Considering the Galactic sample a metallicity effect is found in the zero point in the $VIWK$ $PL$-relation ($-0.6 \\pm 0.4$ or $-0.8 \\pm 0.3$ mag/dex depending on the in- or exclusion of one object), in the sense that metal-rich cepheids are brighter. The small significance is mostly due to the fact that the Galactic sample spans a narrow metallicity range. The error is to a significant part due to the error in the metallicity determinations and not to the error in the fit. Including the 5 MC cepheids broadens the observed metallicity range and a metallity effect of about $-0.27 \\pm 0.08$ mag/dex in the zero point is found in $VIWK$, in agreement with some previous empirical estimates, but now derived using direct metallicity determinations for the cepheids themselves. ",
        "introduction": "The importance of the cepheid Period-Luminosity relation ($PL$-relation) has long been recognised and is the basis of the determination of the Hubble constant by Mould et al. (2000) and Freedman et al. (2001). The most important uncertainties in this derivation are the zero point of the $PL$-relation based on an adopted distance modulus (DM) to the LMC, and the adopted metallicity correction.  Mould et al. (2000) and Freedman et al. (2001) adopt corrections in the Wesenheit-index ($W$ = $V - 2.55 (V-I)$) of $-0.24 \\pm 0.16$ and $-0.2 \\pm 0.2$ mag/dex, respectively. This metallicity effect is important as the galaxies surveyed by the HST Key Project span a factor of 30 in oxygen abundance (Ferrarese et al. 2000).  Theoretical pulsation models lead to different results: linear models (e.g. Sandage et al. 1999, Alibert et al. 1999, Baraffe \\& Alibert 2001) predict a moderate dependence, while non-linear convective models (Bono et al. 1999, Caputo et al. 2000) predict a larger dependence, and in the sense that metal-rich cepheids are fainter (see Table~7 in Groenewegen \\& Oudmaijer 2000, for values at typical periods).  Recently, Fiorentino et al. (2002) suggested that there is also a dependence on the Helium abundance.  On the observational side the results seem to indicate consistently that metal-rich cepheids are brighter, and various estimates have been given in the literature, $-0.88 \\pm 0.16$ mag/dex ($BRI$ bands, Gould 1994;), $-0.44^{+0.1}_{-0.2}$ mag/dex ($VR$ bands, Sasselov et al. 1997), $-0.24 \\pm 0.16$ mag/dex ($VI$ bands, Kochanek 1997), $-0.14 \\pm 0.14$ mag/dex ($VI$ bands, Kennicutt et al. 1998), $-0.25 \\pm 0.05$ mag/dex ($VI$ bands, Kennicutt et al. 2003), and $-0.21 \\pm 0.19$ in $V$, $-0.29 \\pm 0.19$ in $W$,$-0.23 \\pm 0.19$ in $I$,$-0.21 \\pm 0.19$ mag/dex in $K$ (Storm et al. 2004).  The potential drawback or caveat is that no individual abundance determinations of individual cepheids are being used in these studies but rather abundances of nearby H{\\sc ii} regions, or even a mean abundance of the entire galaxy.  The present paper aims at investigating the metallicity dependence from the observational side, but using cepheid individual metallicity determinations. This has become possible because of advances in abundance determinations for the Galaxy (Fry \\& Carney 1997, Andrievsky et al. 2002a,b,c, Luck et al. 2003) and LMC (Luck et al. 1998), as well as recent advances in individual distance estimates for Galactic cepheids based on surface-brightness relations (Fouqu\\'e et al. 2003), a Bayesian statistical analysis to solve the surface-brightness equations (Barnes et al. 2003), direct measurement of distances based on combining radial velocity data with interferometric observations (Kervella et al. 2003), and distance determinations of cepheids in open clusters (Tammann et al. 2003).  The paper is organised as follows. In Sect.~2 the datasets on individual distance and metallicity determinations for Galactic cepheids are presented and compared. The Magellanic Cloud sample is also described. In Sect.~3 the model and the results are presented. The conclusions and future prospects are outlined in Sect.~4. ",
        "conclusions": "An attempt has been made, using the currently available data of cepheids with direct metallicity determinations, to quantify the metallicity dependence of the cepheid $PL$-relation.  For a purely Galactic sample, the range in metallicity covered is too small to draw any firm conclusions. A metallicity effect of $-0.6 \\pm 0.4$ or $-0.8 \\pm 0.3$ mag/dex in $VWK$ is derived, depending on whether the only Galactic star with a significantly sub-solar metallicity is excluded or included.  For the combined sample of Galactic, SMC and LMC cepheids the problem is that most of the MC cepheids have such long periods, in the regime where the $PL$-relation is no longer linear. Restricting the sample in periods to $\\log P < 1.65$ a metallicity effect of about $-0.27 \\pm 0.08$ mag/dex in the zero point is found (in $VIWK$). However, the sample of MC cepheids is presently too small to additionaly solve for a metallicity dependence of the slope in $\\log P$. Formally, the result for the Galactic sample is in agreement with that comprising all three galaxies.  This result based on the sample including the MC cepheids is in agreement with other recent empirical estimates (most recently by Kennicutt et al. 2003 and Storm et al. 2004). However, it is stressed that this study is the first to use {\\em direct} cepheid metallicity determinations to arrive at this result. In agreement with Storm et al. (2004) we find that in the range $V$ to $K$ the metallicity effect seems not to depend on wavelength, contrary to theoretical calculations (e.g. Bono et al. 1999), which predicted a decreasing effect towards longer wavelengths.  The method here developed (Eq.~(1)) is general and can easily be extended to other functional forms, or to include other galaxies. In principle, the distance to a galaxy and the metallicity dependence can be derived independently. For the present restrictive sample this is not the case however, since the distribution over metallicity is correlated with the parent galaxy. In other words, $\\alpha_2$, $\\Delta_{\\rm LMC}$ and $\\Delta_{\\rm SMC}$ are strongly correlated: a change in the DM to the MCs can be 'compensated for' by a change in the metallicity effect. Therefore, it was decided to not solve for the distance moduli. However, in the second block of results in Table~\\ref{TAB-RES1} (no metallicity effect, but allowing for a quadratic term in $\\log P$) there is some indication that the best fitting DM might be slightly longer than 18.50 for the LMC and slightly shorter than 18.90 for the SMC (although there is some scatter depending on the wavelength). Simply adopting values of 18.55 and 18.80, respectively, within the error bars of the current best independent estimates (Feast 2003, Walker 2003), would result in a value for $\\alpha_2$ of $-0.38 \\pm 0.08$ mag/dex in $V$, that is, it would make the metallicity effect slightly stronger (and similarly for the other bands).  It is shown that the error in the independent variable metallicity is a significant contributor to the final error in the metallicity effect as it is larger than the error in the fitting in the dependent variable, but of little effect in the final error in slope and zero point. The mean error in metallicity of the stars in the Galactic plus MC sample is about 0.08 dex. By repeating the Monte Carlo simulations it was derived that for a uniform error bar of 0.05 dex the two sources of error become comparable, and that for a uniform error bar of 0.03 dex or less the error in the metallicity becomes insignificant compared to the fit error.   Although the present results seem to indicate a significant metallicity effect, more data are needed to quantify this effect better. An important issue is an internally consistent metallicity scale. The simulations even suggest an accuracy to a level of 0.05 dex or better. However, this seems very difficult to achieve in practice as the error in the metallicity not only depends on the S/N in the spectra and the number of lines, but also on the uncertainties in the adopted effective temperature, gravity and model atmospheres (e.g. Fry \\& Carney 1997). This implies that an improvement in the present result must come from an increased sample of stars with periods $\\less\\ 45 d$ spread over a range in metallicities as large as possible."
    },
    "0403/astro-ph0403494_arXiv.txt": {
        "abstract": " ",
        "introduction": "QSOs are intrinsically luminous and therefore can be seen rather easily at large distances; but they are rare, and finding them requires surveys over large areas. As a consequence, at present, the number density of QSOs at high redshift is not well known.  Recently, the Sloan Digital Sky Survey (SDSS) has produced a breakthrough, discovering QSOs up to $z=6.43$ \\cite{Fan03} and building a sample of six QSOs with $z>5.7$.  The SDSS, however, has provided information only about very luminous QSOs ($M_{1450} \\mincir -26.5$), leaving unconstrained the faint end of the high-$z$ QSO Luminosity Function (LF), which is particularly important to understand the interplay between the formation of galaxies and super-massive black holes (SMBH) and to measure the QSO contribution to the UV ionizing background \\cite{Madau99}. New deep multi-wavelength surveys like the Great Observatories Origins Deep Survey (GOODS), described by M.Dickinson, M.Giavalisco and H.Ferguson in this conference, provide significant constraints on the space density of less luminous QSOs at high redshift. Here we present a search for high-$z$ QSOs, identified in the two GOODS fields on the basis of deep imaging in the optical (with {\\em HST}) and X-ray (with {\\it Chandra}), and discuss the allowed space density of QSOs in the early universe, updating the results presented in \\cite{cristiani04}. ",
        "conclusions": "At $z>4$ the space density of moderate luminosity ($M_{1450} \\simeq -23$) QSOs is significantly lower than the prediction of simple recipes matched to the SDSS data, such as a PLE evolution of the LF or a constant universal efficiency in the formation of SMBH in DMH. A flattening of the observed high-$z$ LF  is required below the typical luminosity regime ($M_{1450} \\mincir -26.5$) probed by the SDSS. An independent indication that this flattening must occur comes from the statistics of bright lensed QSOs observed in the SDSS \\cite{Wyithe02} that  would be much larger if the LF remains steep in the faint end. A similar result has been obtained at $5 \\mincir z \\mincir 6.5$, by \\cite{Barger03}. The QSO contribution to the UV background is insufficient to ionize the IGM at these redshifts. This is an indication that at these early epochs the formation or the feeding of SMBH is strongly suppressed in relatively low-mass DMH, as a consequence of feedback from star formation \\cite{granato01} and/or photoionization heating of the gas by the UV background \\cite{haiman99}, accomplishing a kind of inverse hierarchical scenario."
    },
    "0403/hep-th0403138_arXiv.txt": {
        "abstract": "The effective equation of state of normal matter is changed in theories where the size of the compact space depends upon the local energy density.  In particular we show how the dilution of a fluid due to the expansion of the universe can be compensated by an increase of the effective coupling of that fluid to gravity in the presence of a potential which acts to reduce the size of the compact space.  We estimate how much cosmic acceleration can be obtained in such a model and comment on the difficulties faced in finding an appropriate potential. ",
        "introduction": "Current astrophysical observations suggest that there are at least two separate epochs in the history of the universe where accelerated expansion occurred.  The first is in the very early universe where a period of accelerated expansion could explain the uniform temperature of the cosmic microwave background radiation across the sky.  The second is the apparent acceleration of the universe deduced from observations of type 1a supernovae which seems to be occurring today and to have begun extremely recently in cosmological terms \\cite{sn1a}.  It is not possible to obtain accelerated expansion from normal matter or radiation - as the universe expands, their energy density shrinks too rapidly, so in order to explain the observations one is forced to consider other forms of stress energy.  It is very easy to obtain accelerated expansion from the potential of a self interacting scalar field, although typically one does not obtain a prolonged period of such behaviour without fine-tuning which makes it difficult to explain the inflation required in the early universe.  In the same way it makes it difficult to arrange a period of acceleration which began only very recently.  It is therfore worth considering other mechanisms which could give rise to acceleration to see if there are any reasonable alternatives to the orthodox mechanisms.  In theories with more than 3 spatial dimensions, the values of the couplings in the 3+1 dimensional theory depend upon the details of the compactification of the higher dimensions.  Perhaps the simplest example of this is the ratio between the Newton's constant which appears in the higher dimensional theory $G_{D}^{-1}=8\\pi M_F^{D-2}$ and that which appears in the 3+1 dimensional theory $G_{4}^{-1}=8\\pi M_{Pl}^2$, which is simply the volume of the compact space\\footnote{For a string theory the usual situation is $8\\pi G_{D}=M_F^{2-D}=g_s^{2}l_s^{D-2}/8$ where $g_s$ and $l_s$ are the string coupling and length respectively \\cite{witten}.}.  If one then allows that volume to vary, the 4D Newton's constant will also vary relative to the underlying higher dimensional length scale.  The effect of a varying volume will be to add a dynamical scalar in front of the 4D Ricci Scalar in the action ${\\mathcal{L}}=\\sqrt{-g}\\Phi R[g]+...$ but it is always possible to perform a conformal transformation on the metric to the Einstein frame such that the the action for gravity takes the form (see appendix)  \\begin{equation} S=\\frac{M_{Pl}^{2}}{2}\\int d^{4}x \\sqrt{-g}\\left\\{R-\\frac{1}{2}\\partial^\\mu\\phi \\partial_\\mu \\phi\\right\\} \\end{equation} where the scalar $\\phi$ represents the volume of the compact space (variously referred to as the radion, dilaton, modulus, breathing mode etc.).  If we start in $D$ dimensions and compactify to $3+1$ dimensions $\\phi$ is given by $\\cite{argurio}$ \\begin{eqnarray} \\phi(t)-\\phi_0=-\\frac{(D-4)}{2}\\sqrt{\\frac{D-2}{D-4}}\\ln\\left(\\frac{r(t)}{r_0}\\right)&&\\nonumber\\\\ \\equiv-\\frac{(D-4)}{\\beta}\\ln\\left(\\frac{r(t)}{r_0}\\right)&& \\label{beta} \\end{eqnarray} where $r_0$ is the radius of the compact (D-4) torus today so that $\\phi-\\phi_0$ parameterises the relative change in volume over time. The conformal transformation should not affect any physics derived from the Lagrangian, and indeed one finds that the effect of the varying compact space has been re-absorbed into a variation of the effective density of matter.  In the Einstein frame the left hand side of the field equations will have the same form as Einstein gravity for given space-time symmetries. However, the equivalence principle is broken (the gravitational attraction between two particles will depend on the local value of the field $\\phi$) and the density that one uses to solve the equations will be given by $\\rho_{eff}=e^{\\beta\\phi}\\rho$ $\\cite{lidsey}$.  If one assumes that the total density is the sum of the matter density and a potential of the (ad-hoc at this stage) form $V(\\phi)=\\bar{V}e^{-\\alpha\\phi}$ where $\\bar{V}$ and $\\alpha$ are constants, the total effective density in the Einstein frame is given by \\begin{equation} \\rho_{total}=\\bar{V}e^{-\\alpha\\phi}+\\rho e^{\\beta\\phi} \\label{mixeddensity} \\end{equation} and (for $\\omega<1/3$) the expectation value of $\\phi$ will be at the minimum of this effective density, i.e. \\begin{equation} \\langle\\phi\\rangle=\\frac{1}{\\alpha+\\beta}\\ln\\left(\\frac{4}{(1-3\\omega)}\\frac{\\alpha}{\\beta}\\frac{\\bar{V}}{\\rho}\\right)\\qquad ; \\qquad \\omega < 1/3 \\label{expectation} \\end{equation} so that although $\\phi-\\phi_0$ is always negative since $r\\ge r_0$, $\\phi$ is always positive as we will only be considering situations where $\\rho\\ll\\bar{V}$.  The expectation value of $\\phi$ and the coupling of the matter to gravity therefore depends upon the local density of stress-energy.  The authors of $\\cite{chameleon}$ refer to $\\phi$ as a chameleon field when it is behaving like this since the mass of the field changes according to the density of the local medium (see also \\cite{barrow}).  The fact that the expectation value of $\\phi$ is a function of the local density is more important for our purposes than its mass. Consider a universe with a spatially homogeneous distribution of matter.  As the universe expands, the energy density will dilute in the normal way but the expectation value of the field $\\phi$ will also change, increasing the coupling of that matter to gravity. In this way the effective dilution of the gravitating energy due to the expansion will be reduced and the effective equation of state of the energy will change. In particular we will be interested in finding out under what conditions this situation can lead to accelerated expansion.  In order for such an increase in coupling to compensate for dilution, we require a potential such as the one written down in equation (\\ref{mixeddensity}) which, in the absence of matter, would cause the compact space to shrink as the universe expands.  This is not easy to achieve.  The reader should be aware of some closely related previous work where the mass of particles changes with the cosmically varying expectation value of a scalar field $\\cite{vamps}$ (see also \\cite{rachel}) Here, we are changing the gravitational mass of the particles whilst their inertial mass stays a fixed fraction of $M_{F}$. We are considering only cases when the mass scale associated with the energy density, $\\rho^\\frac{1}{4}$ is much less than the inverse radius of the compact space.  It is therefore irrelevant as to whether the matter is confined to a brane or not, the size of the compact space will not make any difference to its bare stress energy in the Jordan-string frame since we will assume that any matter which does propagate in the bulk consists only of zero modes.  The low energy gauge coupling of those fields which are able to propagate in the compact space will change with the size of the compact manifold, but we will assume that no phase transitions occur because of this effect, and will only consider particles which are given mass via some Yukawa coupling rather than some confinement scale ('electron' like particles rather than 'baryon' like ones) In the next section we will find out how the effective equation of state behaves for the simple situation presented in equation (\\ref{mixeddensity}). We will then compare the situation with that of an exponential potential and no matter, in other words power law inflation.  Next we will set up some checks that our scenario must pass in order to be self consistent. Then we will find out how much expansion it is possible to obtain using matter with such a potential without violating our consistency checks.  After that we will discuss why it is very difficult to find a potential like the one used above but we will show that it is possible to find some well motivated potentials which contain regions where the compact space is dynamically driven to smaller radii.  Finally we will show that it is impossible to obtain acceleration using only the matter and the motion of the scalar field without another potential being present. ",
        "conclusions": "In this paper we showed that in models with compactified higher dimensions, changes in the size of the compact space change the effective equation of state of matter.  In particular, in the presence of a potential which tends to reduce the size of the compact space, one can obtain acceleration from normal matter.  Assuming a simple form for this potential we have calculated the maximum amount of expansion one could get from a such a model without non-adiabatically producing KK excitations around the compact space. If we did excite KK modes, their effect on the evolution of the compact and non-compact spaces would be extremely difficult to calculate.  For a prolonged period of acceleration, the compact space would have to start out very large compared to its size at the end of this period.  We have also tried to explain why it is difficult to find potentials of a suitable form to lead to acceleration, although we have shown that realistic potentials already exist which contain regions where the compact space would be pushed to smaller sizes as the matter or radiation becomes diluted.  We ended by pointing out that it is hopeless to try and get acceleration without the presence of a potential.  So far we have said nothing about Dark energy.  It has been pointed out recently that if one is able to find an appropriate potential of the form that we have used in this paper one can give a mass to $\\phi$ which changes according to the local density $\\cite{chameleon}$.  In particular, tests of tensor gravity could be passed in denser mediums such as the solar system or neutron star binary systems, while a scalar-tensor theory would describe gravity in the low density voids between clusters of galaxies.  The equation of state of dark matter, could be reduced in these regions as it is in other models of varying mass particles (VAMPS) \\cite{vamps}\\cite{rachel}.  In this way, one might hope to shed light on the coincidence between the energy density of dark energy and dark matter."
    },
    "0403/astro-ph0403480_arXiv.txt": {
        "abstract": "We address the fine-tuning problem of dark energy cosmologies which arises when the dark energy density needs to initially lie in a narrow range in order for its present value to be consistent with observations. As recently noticed, this problem becomes particularly severe in canonical Quintessence scenarios, when trying to reproduce the behavior of a cosmological constant, i.e. when the dark energy equation of state $w_Q$ approaches $-1$: these models may be reconciled with a large basin of attraction only by requiring a rapid evolution of $w_Q$ at low reshifts, which is in conflict with the most recent estimates from type Ia Supernovae discovered by Hubble Space Telescope. Next, we focus on scalar-tensor theories of gravity, discussing the implications of a coupling between the Quintessence scalar field and the Ricci scalar (``Extended Quintessence''). We show that, even if the equation of state today is very close to $-1$, by virtue of the scalar-tensor coupling the quintessence trajectories still possess the attractive feature which allows to reach the present level of cosmic acceleration starting by a set of initial conditions which covers tens of orders of magnitude; this effect, entirely of gravitational origin, represents a new important consequence of the possible coupling between dark energy and gravity. We illustrate this effect in typical Extended Quintessence scenarios. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403355_arXiv.txt": {
        "abstract": "We present the results of mapping the planetary nebula NGC\\,6369 using multiple long slit spectra taken with the CTIO 1.5m telescope. We create two dimensional emission line images from our spectra, and use these to derive fluxes for 17 lines, the H$\\alpha$/H$\\beta$ extinction map, the [SII] line ratio density map, and the [NII] temperature map of the nebula. We use our photoionization code constrained by these data to determine the distance, the ionizing star characteristics, and show that a clumpy hour-glass shape is the most likely three-dimensional structure for NGC\\,6369. Note that our knowledge of the nebular structure eliminates all uncertainties associated with classical distance determinations, and our method can be applied to {\\it any spatially resolved emission line nebula}. We use the central star, nebular emission line, and optical+IR luminosities to show that NGC\\,6369 is matter bound, as about 70\\% of the Lyman continuum flux escapes. Using evolutionary tracks from \\cite{B95} we derive a central star mass of about 0.65\\,M$_{\\odot}$. ",
        "introduction": "The planetary nebula (PN) NGC\\,6369 -- shown in Fig.\\,\\ref{n2ima} in the light of H$\\alpha$\\,+\\,N[II]$\\lambda$658.4nm -- is an object with a complex morphology, consisting of a main bright annulus of diameter 40\\arcsec, and fainter, curved outer structures on two sides in the E-W direction. Deep narrow band H$\\alpha$\\, and \\,N[II]$\\lambda$6584 images of NGC\\,6369 were obtained by \\cite{CSSP03} in a survey to look for faint outer halos around PNe. Apart from being deeper, their image is not significantly different from the one shown in Fig.\\,1, and no large faint halo was found.  The HST image of NGC\\,6369 available at http://heritage.stsci.edu/2002/25/index.html also shows the same main features as our image but has additional resolved details, and tags the H$\\alpha$, [OIII]500.7nm, and [NII]658.4nm lines with the colors red, blue and green respectively. In the section on observational results we discuss this composite HST image in more detail.  \\clearpage \\begin{figure} \\includegraphics[scale=0.65]{f1.eps} \\caption{Narrow band filter image of NGC\\,6369 in H$\\alpha $+{[}NII{]} from \\cite{SCM92}. Note the bright ring and outer faint ansae to the E and W of the object. N is up, E to the left, and the plate scale is 0.26 ''/pix. \\label{n2ima}} \\end{figure} \\clearpage  The equatorial and galactic coordinates for NGC\\,6369 are $\\alpha$ = 17$^h$\\,29$^m$\\,20$^s$.40 and $\\delta$ = 23$^o$\\,45\\arcmin\\,37\\arcsec.9 (ICRS2000), and l = 2.43$^o$ b = +5.85$^o$ respectively. The extreme distances determined for NGC\\,6369 are 0.33\\,kpc obtained by \\cite{AGNR84}, and 2.00\\,kpc obtained by \\cite{GPG86}, out of a total of 10 listed by \\cite{AO92}, from which we obtain a mean value for the distance to NGC\\,6369 of 1039$\\pm$ 538 pc. Note that these ten distances were determined using several statistical methods, each with their own systematic biases and assumptions, so that the average probably is a reasonably unbiased estimator, within the (large) error.  The temperature of the central star obtained from the Zanstra method by \\cite{GP89} is $T(H)_\\mathrm{Z}=67\\,600~\\mathrm{K}$. The star has been classified as a WC4 star by \\cite{TAS93}.  Spectra of NGC\\,6369 were obtained by \\cite{AKSJ1989}, using an IDS (Image Dissector Scanner) as detector with double 4\\,\\arcsec~apertures, by \\cite{AK87} using an ITS (Image Tube Scanner) with a $2\\arcsec\\,\\times\\,10\\arcsec$ slit, and by \\cite{PPPP} using an IDS and a combination of different apertures not bigger than $13\\arcsec\\,\\times\\,13\\arcsec$. An expansion velocity of $41.5$ km s$^{-1}$ has been determined by \\cite{MWF88}. The system radial velocity was reported to be -106\\,km.s$^{-1}$ by \\cite{W53}, and -101\\,km.s$^{-1}$ by \\cite{MWF88}.  The spectral energy distribution (SED) of NGC\\,6369 is shown in Fig.\\,\\ref{sed} and has been computed using flux values from \\cite{AO92} and \\cite{SK02} over a range of wavelengths from 0.44 to 1100\\,$\\mu$m. Note that we plot $\\lambda$\\,F$_{\\lambda}$ against wavelength and below we will use the F$_{\\lambda}$ values to determine the observed optical+IR luminosity of NGC\\,6369.  \\clearpage \\begin{figure} \\includegraphics[width=\\columnwidth]{f2.eps} \\caption{The SED for NGC\\,6369 between 0.44 and 1100\\,$\\mu$m using flux values from the literature (see text).\\label{sed}} \\end{figure} \\clearpage  Planetary nebulae are generally classified according to their appearance, and such classification is then used for studies of the formation and evolution of these objects (for example \\citep{B87}). Such a classification, based on the observed two dimensional (2-D) brightness distribution in a given line or lines can be quite misleading. This could be the case for NGC 6369, for which a round morphology is evident from the observed images. Different 3-D geometries can produce the same observed 2-D morphology as has been shown by \\cite{M00}. Furthermore, images produced by ions of different ionization degree can be very different, due to the radiation transfer within the nebulae. Only detailed modeling, which reproduces the brightness intensity distribution of different lines, as well other observed parameters, can provide more reliable information about the 3-D geometry of the gas distribution. The determination of the actual 3-D gas distribution in planetary nebulae is essential for understanding their formation and evolution, as well as that of the ionizing star.  By determining the 3-D structure of nebulae, we eliminate the large uncertainties that have plagued classical statistical distance determination methods for over 5 decades. Assumptions about the filling factor, constancy of ionized mass or diameter, mass-radius relationships etc. are not needed here: we {\\it know} what the structure and ionized mass are, and can therefore determine distances to much greater accuracy than before. In this paper we explain in detail how we can determine this 3-D structure from long slit spectra and our 3-D photo-ionization model, and apply it to the case of NGC\\,6369  In summary, we obtain the spatial structure of the object along with its chemical composition, ionizing source temperature and luminosity, mass, as well as an {\\it independent} distance, in a self-consistent manner. In \\S2 we discuss the observations and basic reduction procedures, as well as the details of the image reconstruction technique used to obtain the line intensity maps. In \\S3 the results obtained from these maps are discussed: the reddening correction of the images, total line fluxes and the computed temperature and density maps. In \\S4 we present the model results of the 3-D photoionization code, and we discuss the derived quantities Finally, in \\S5 we give our conclusions. ",
        "conclusions": "We have presented spectrophotometric maps of NGC\\,6369. These maps provided spatially resolved information for many emission lines and precise total fluxes for the whole nebula. The images produced with this technique were used to study the nebula with the usual diagnostic ratios. Each image was corrected for reddening pixel by pixel using the H$\\alpha $/H$\\beta$ image. This correction lead to some significant differences between the corrected and uncorrected images, as was shown in Figs.\\,5 \\& 7.  The H$\\alpha $/H$\\beta$ map shows some interesting features. In Fig.\\,4 we show that the extinction is not uniform across the nebula, as it varies between 6 and 18. The structure follows the main nebular morphology, which could indicate that dust and neutral material are present. This is to be expected since the nebula shows a very clumpy structure, which may cause shadowing within the gas, allowing for the survival of grains and neutral material. The prominent features on either side of the nebula do not show significant differences in extinction when compared to the main region.  The temperature map derived from the observed integrated line of sight intensities showed some structure but due to the low signal to noise ratio, this should be treated with caution. Within the errors and resolution of our observations the temperature can be considered to be constant across the nebula.  We also show that the density map indicates a decrease in density for the central regions. This decrease, as has been discussed by \\cite{M00}, is not compatible with a closed shell structure. Based on this map we propose an hour-glass structure for the main nebula, which has reproduced all the observational features of NGC\\,6369. This indicates that the one needs to use the [SII] ratio in addition to images to distiguish between open and closed structures.  The position of the outer condensations or ansae --which are off-set by 30$^o$ from the main nebular symmetry axis-- was obtained by matching model images with the observed line images (Fig.\\,\\ref{maps_full}) as well as matching the model [SII] density map with the observed map (Fig.\\,\\ref{dens_map}). The presence of these condensations can be accounted for by earlier ejection of matter and precession can account for their deviation from the symmetry axis. Many authors have dealt with these issues; for more complete discussions and models see for example \\cite{CFJ96}, \\cite{CFLJ95}, \\cite{G97}, and \\cite{LP97}.   Some of the parameters we have derived (and use as a sanity check) are:  The total luminosity of the observed lines, L$_l$\\,=\\,1150\\,L$_{\\odot}$; that derived from the SED by integrating the F$_{\\lambda}$ curve, L$_o$\\,=\\,7.1$\\,\\times$\\,10$^{-4}$\\,d$^2$ where d is the distance to NGC\\,6369 in pc, resulting in L$_o$\\,=\\,1700\\,L$_{\\odot}$ . Correcting this value using the method of Myers et al. (1987) we obtain 2550\\,L$_{\\odot}$. The luminosity of the central star is L$_{tot}$\\,=\\,8100\\,L$_{\\odot}$ so that the ratios of the line and optical+IR luminosities to the total luminosity are respectively: L$_l$/L$_{tot}$\\,=\\,0.14 and L$_o$/L$_{tot}$\\,=\\,0.3. Assuming that the absorbed flux is re-radiated in the IR, mainly the IRAS bands, the integrated optical+IR luminosity indicates that about 70\\% of the UV flux (Lyman continuum) escapes from the nebula. This is not implausible as the nebula is open at it's ``poles'' and is clumpy, allowing the radiation to pass through the many ``holes''.  From the input matter distribution and abundances used in the model, we calculate the mass of the nebular ionized gas to be $M_{neb}\\,=\\,1.8\\,M_{\\odot}$. If we now use our values of luminosity and temperature for the central star and compare them with the evolutionary models of \\cite{B95} we can obtain the mass of the central core. Using this procedure we can see from Fig.\\,\\ref{evol_track} that the best fitting track for our data corresponds to a mass somewhat higher than $M_{core}\\,=\\,0.625 M_{\\odot}$, say, $\\approx\\,0.65\\,M_{\\odot}$, and to an initial mass of $M_{0}=3 M_{\\odot}$. If we sum our nebular and core masses, we obtain $M_{core}\\,+\\,M_{neb}\\,=\\,2.4\\,M_{\\odot}$ approximately. If we use the line intensity errors as a measure of the goodness of the model results and combine the luminosity and temperature errors we obtain approximately 20-25\\% for the uncertainty in $M_{core}\\,+\\,M_{neb}$, placing the obtained initial stellar mass within the uncertainty of our determination. The value for $M_{neb}$ calculated from the model input density is a lower limit estimate of the total nebular mass as at least some material is present in the form of dust, as indicated qualitatively by our extinction map structure. Additionally, there may be neutral gas for wich we have no constraints.  Note that typical parameters for a WC4 PN central star from \\cite{A64} are T\\,=\\,90kK, and R\\,=\\,0.4\\,R$_{\\odot}$, very close to our derived values of 91000\\,K and 0.4\\,R$_{\\odot}$. Using g\\,=\\,G\\,.\\,M\\,/\\,R$^2$ in CGS units, we obtain log(g)\\,=\\,5.1, which is much lower than the value of 5.8 derived from the depth of our He break. This is due to the fact that the central star has an extended atmosphere, again showing the consistency of the model results.  \\clearpage \\begin{figure} \\includegraphics[width=\\columnwidth]{f12.eps} \\caption{Comparison of model temperature and luminosity obtained for the central star with model tracks calculated by \\cite{B95}.\\label{evol_track}} \\end{figure} \\clearpage  Using our photoionization code and the proposed structure we obtained a complete 3-D model for the NGC\\,6369. The fitted model line intensities show excellent agreement with the observed values. The obtained distance of d=1550 pc, is well within the range present in the literature obtained from different methods. The model temperature for the ionizing star is similar to the Zanstra (He) value discussed in \\S3.3. The temperature and luminosity values obtained from the model as well as the total nebular plus central star mass show good agreement with the stellar evolution models of \\cite{B95}.  Using multiple long slit spectroscopy, we can determine accurate distances, 3-D structures, abundances, ionized masses, central star masses, luminosities, and temperatures to any spatially resolved emission line nebula, assuming there are no strong shocks or extreme morphologies involved."
    },
    "0403/astro-ph0403025_arXiv.txt": {
        "abstract": "{We report the first likely spectroscopic confirmation of a $z\\sim 10.0$ galaxy from our ongoing search for distant galaxies with ISAAC/VLT. Galaxy candidates at $z \\gtapprox 7$ are selected from ultra-deep $JHKs$ images in the core of gravitational lensing clusters for which deep optical imaging is also available, including HST data. The object reported here, found behind Abell 1835, exhibits a faint emission line detected in the $J$ band, leading to $z=10.0$ when identified as \\lya, in excellent agreement with the photometric redshift determination. Redshifts $z < 7$ are very unlikely for various reasons we discuss. The object is located on the critical lines corresponding to z$=$9 to 11. The magnification factor $\\mu$ ranges from 25 to 100. For this object we estimate $SFR(\\lya) \\sim  (0.8-2.2)$ \\msun yr$^{-1}$ and $SFR(UV) \\sim (47-75)$ \\msun\\ yr$^{-1}$, both uncorrected for lensing. The steep UV slope indicates a young object with negligible dust extinction. SED fits with young low-metallicity stellar population models yield (adopting $\\mu=25$) a lensing corrected stellar mass of $M_\\star \\sim 8 \\times 10^6$ \\msun, and luminosities of $2 \\times 10^{10}$ \\lsun, corresponding to a dark matter halo of a mass of typically $M_{\\rm tot} \\ga 5 \\times 10^8$ \\msun. In general our observations show that under excellent conditions and using strong gravitational lensing direct observations of galaxies close to the ``dark ages'' are feasible with ground-based 8-10m class telescopes. ",
        "introduction": "Spectacular progress during the last decade has permitted direct observations of galaxies and quasars out to redshifts $z \\sim 6.6$ (Hu et al.\\ 2002, Fan et al.\\ 2003, Kodaira et al.\\ 2003), or in other words over more than 90 \\% of cosmic time. With the exception of the cosmic microwave background, direct exploration of more distant objects has so far been hampered by the need for observations beyond the optical domain, and technical difficulties related to near-IR observations of such faint objects. However, by combining a strong magnification of background galaxies by well known gravitational lensing clusters with present-day near-IR instruments on 8-10m class telescopes such an attempt appears feasible (Schaerer \\& Pell\\'o\\ 2002; Pell\\'o et al.\\ 2003). We present here the first spectroscopic results obtained from our prototype observational program with ISAAC/VLT on one of our best $z \\ga 7$ galaxy candidates.  In Sect.~\\ref{observations} we summarize the observational strategy adopted in the present study, and we present the photometric and spectroscopic characteristics of this object. The redshift identification is discussed in Sect.\\ \\ref{s_z}. The physical properties of this galaxy and implications of our finding are briefly discussed in Sect.~\\ref{results}. ",
        "conclusions": ""
    },
    "0403/astro-ph0403096_arXiv.txt": {
        "abstract": "The study of the universe at energies above 100 GeV is a relatively new and exciting field.  The current generation of pointed instruments have detected TeV gamma rays from at least 10 sources and the next generation of detectors promises a large increase in sensitivity.  We have also seen the development of a new type of all-sky monitor in this energy regime based on water Cherenkov technology (Milagro).  To fully understand the universe at these extreme energies requires a highly sensitive detector capable of continuously monitoring the entire overhead sky.  Such an instrument could observe prompt emission from gamma-ray bursts and probe the limits of Lorentz invariance at high energies. With sufficient sensitivity it could detect short transients ($\\sim$15 minutes) from active galaxies and study the time structure of flares at energies unattainable to space-based instruments.  Unlike pointed instruments a wide-field instrument can make an unbiased study of all active galaxies and enable many multi-wavelength campaigns to study these objects. This paper describes the design and performance of a next generation water Cherenkov detector.  To attain a low energy threshold and have high sensitivity the detector should be located at high altitude ($>$ 4km) and have a large area ($\\sim$40,000 m$^2$).  Such an instrument could detect gamma ray bursts out to a redshift of 1, observe flares from active galaxies as short as 15 minutes in duration, and survey the overhead sky at a level of 50 mCrab in one year. ",
        "introduction": "\\label{sec:Introduction}  The past 15 years have seen large advances in the capabilities of ground-based gamma ray detection, from the pioneering observation of the Crab Nebula by the Whipple observatory in 1989\\cite{weekes1989} to the new generation of air Cherenkov telescope arrays such as HESS\\cite{hofmann2003}, VERITAS\\cite{veritas}, and CANGAROO\\cite{yoshikoski1999}  and large area air Cherenkov telescopes such as STACEE \\cite{hanna2002}, CELESTE\\cite{pare2002}, and MAGIC\\cite{lorenz2003}.  There are now at least 10 known sources of very-high-energy (VHE) gamma rays\\cite{horan2003}. The physics of these objects is astounding: from spinning neutron stars to super-massive black holes, these objects manage to accelerate particles to energies well in excess of 10 TeV. How this acceleration occurs is not well understood and there is not universal agreement on what particles are being accelerated in some of these sources.  At lower energies EGRET has detected over 270 sources of high-energy gamma rays\\cite{hartman1999} and GLAST is expected to detect several thousand sources.  In addition there are transient sources such as gamma-ray bursts that have to date eluded conclusive detection in the VHE regime (despite some tantalizing hints\\cite{atkins2000}).  The paucity of VHE sources can be traced to the nature of the existing instruments: they are either narrow field instruments that can only view a small region of the sky at any one time and can only operate on clear moonless nights (Whipple, HEGRA, etc.) or large field instruments with limited sensitivity (Milagro, Tibet Array).  The Milagro observatory has pioneered the use of a large area water Cherenkov detector for the detection of extensive air showers.  Since an extensive air shower (EAS) array directly detects the particles that survive to ground level it can operate continuously and simultaneously view the entire overhead sky.  With the observation of the Crab Nebula and the active galaxies Mrk 421 and Mrk 501, Milagro has proven the efficacy of the technique and its ability to reject the cosmic-ray background at a level sufficient to detect sources\\cite{atkins2003}. At the same time the Tibet group\\cite{amenomori1999} has demonstrated the importance of a high-altitude site and what can be accomplished with a classical scintillator array at extreme altitudes.  A detector with the all-sky and high-duty factor capabilities of Milagro, but with a substantially lower energy threshold and a greatly increased sensitivity, would dramatically improve our knowledge of the VHE universe. Reasonable design goals for such an instrument are: \\begin{itemize}% \\item Ability to detect gamma-ray bursts to a redshift of 1.0 \\item Ability to detect AGN to a redshift beyond 0.3 \\item Ability to resolve AGN flares at the intensities and durations observed by the current generation of ACTs \\item Ability to detect the Crab Nebula in a single transit \\end{itemize}  This paper describes a design for a next generation all-sky VHE gamma-ray telescope, the HAWC (High Altitude Water Cherenkov) array, that satisfies these requirements.  To quantify the definition of observing ``short'' flares from AGN, previous measurements of flare intensities and durations by air Cherenkov telescopes can be used.  To date the shortest observed flares have had $\\sim$15 minute durations with an intensity of 3-4 times that of the Crab\\cite{gaidos1996}.  The low energy threshold needed to accomplish these goals requires that the detector be placed at extreme altitudes (HAWC would be situated at an altitude of $\\sim$4500 meters) and the required sensitivity demands a large area detector - of order 40,000 m$^2$.  Section \\ref{sec:Particle_Detection} discusses the limiting performance of an EAS array based on the properties of the EAS, section \\ref{sec:Detector_Description} gives a physical description of the HAWC and section \\ref{sec:Detector_Performance} details the expected performance of HAWC. ",
        "conclusions": "\\label{sec:Conclusions}  A design for the next generation all-sky VHE gamma-ray telescope has been presented.  This instrument, HAWC, would be over 20 times more sensitive than the Milagro detector. With the ability to continuously view the entire overhead sky HAWC will be an excellent complement to both GLAST and the coming generation of air Cherenkov telescopes (HESS, MAGIC, VERITAS, and CANGAROO III).  With comparable sensitivity to GLAST it will be the only instrument capable of monitoring the many thousands of sources that GLAST is expected to detect at higher energies.  In addition to searching the sky for galactic sources (the VHE complement to the 150 EGRET unidentified objects), and active galaxies, the low-energy sensitivity of a detector placed at high altitude ensures that such an instrument will detect any VHE emission from gamma-ray bursts.  Perhaps most importantly an open aperture instrument with this level of sensitivity could discover completely new and unexpected phenomena that have so far eluded detection."
    },
    "0403/astro-ph0403575_arXiv.txt": {
        "abstract": "s{ We report results from {\\it XMM-Newton} observations of thirteen X-ray bright BL Lacertae objects, selected from the {\\it Einstein} Slew Survey sample.   The spectra are generally well fit by power-law models, with four objects having hard ($\\alpha<1; F_\\nu \\propto \\nu^{-\\alpha}$) spectra that indicates synchrotron peaks at $>5$ keV.  None of our spectra show line features, indicating that soft X-ray absorption ``notches'' must be rare amongst BL Lacs, rather than common or ubiquitous as had previously been asserted.  We find significant curvature in most of the spectra.  This curvature is almost certainly intrinsic, as it appears nearly constant from  0.5 to 6 keV, an observation which is inconsistent with the small columns seen in these sources.} ",
        "introduction": "The nature of the X-ray emission and absorption from BL Lacs is still an open question.  Most BL Lac objects have X-ray spectra that can be fit by power-laws within  smaller bandpasses, such as that of {\\it ROSAT} ([12],[17]) More recent results from {\\it ASCA} [7] and {\\it BeppoSAX} ([1],[11],[19]) have generally confirmed this spectral morphology, and also added onto it the possibility of intrinisic spectral curvature across a wider bandpass ([4],[6],[9],[16]). In addition, earlier missions had indicated that some BL Lac objects showed a deficit in soft X-rays below a power-law model, which had been interpreted by invoking X-ray absorption features at 0.5--0.8 keV ([3],[8],[14],[15]).  But not all bright BL Lacs were found to require such features  (e.g., [5]). ",
        "conclusions": ""
    },
    "0403/astro-ph0403269_arXiv.txt": {
        "abstract": "We present mass and radius derivations for a sample of very young, mid- to late M, low-mass stellar and substellar objects in Upper Scorpius and Taurus.  In a previous paper, we determined effective temperatures and surface gravities for these targets, from an analysis of their high-resolution optical spectra and comparisons to the latest synthetic spectra.  We now derive extinctions, radii, masses and luminosities by combining our previous results with observed photometry, surface fluxes from the synthetic spectra and the known cluster distances.  These are the first mass and radius estimates for young, very low mass bodies that are {\\it independent} of theoretical evolutionary models (though our estimates do depend on spectral modeling).  We find that for most of our sample, our derived mass-radius and mass-luminosity relationships are in very good agreement with the theoretical predictions.  However, our results diverge from the evolutionary model values for the coolest, lowest-mass targets: our inferred radii and luminosities are significantly larger than predicted for these objects at the likely cluster ages, causing them to appear much younger than expected.  We suggest that uncertainties in the evolutionary models - e.g., in the choice of initial conditions and/or treatment of interior convection - may be responsible for this discrepancy.  Finally, two of our late-M objects (USco 128 and 130) appear to have masses close to the deuterium-fusion boundary ($\\sim$9--14 Jupiters, within a factor of 2).  This conclusion is primarily a consequence of their considerable faintness compared to other targets with similar extinction, spectral type, and temperature (difference of $\\sim$ 1 mag).  Our result suggests that the faintest young late-M or cooler objects may be significantly lower in mass than current theoretical tracks indicate. ",
        "introduction": "The last few years have witnessed a dramatic swelling in the ranks of objects at the bottom of the Main Sequence, and in the substellar regime beyond.  Hundreds of ultra-low-mass stars and brown dwarfs have been uncovered, both in the field and in star-forming regions.  Studies of young clusters even suggest the presence of isolated planetary mass objects \\citep{Zapa00, Lucas00}.  The existence and properties of all these low-mass bodies have profound implications for a host of issues, ranging from the dominant mechanisms for star and planet formation \\citep{Boss01, Reipurth01, Bate02, Padoan02}, to the birthline and early evolution of low-mass objects \\citep{Hartmann03}, to the shape of the initial mass function \\citep{Briceno02}.  A reliable determination of mass is intrinsic to the ultimate resolution of these questions.  Presently, masses (and ages) are most widely inferred by comparing observables such as temperature and luminosity to the predictions of theoretical evolutionary tracks.  However, these models remain largely unverified for very low masses.  The simplest test is to derive dynamical masses for the components of binary (or higher-order) systems with known orbital parameters, and compare them to the theoretical values derived from other, directly observed quantities (e.g., $\\lbol$ and \\teff).  Unfortunately, this is impeded for very low-mass stars and substellar objects by the current paucity of suitable multiple systems.  In most known cases, one can either deduce dynamical masses but not theoretical ones (because the components are not directly detected), or vice versa (because the orbital parameters remain indeterminate).  The one exception is HD 209458, in which both are available \\citep{Charbonneau02}.  The comparison of theory to observations in this case does reveal some large uncertainties in the former, and underlines the usefulness of such tests (Baraffe et al. 2003; Burrows, Sudarsky \\& Hubbard 2003).  However, it is not very illuminating as a general evaluation of the models: the proximity between the planetary companion and the star in this instance engenders special insolation effects, precluding an extension of the results to free-floating brown dwarfs and planetary mass objects (or to planets with larger orbital radii).  The situation is likely to improve in the near future, at least in the field - several promising systems with directly detected, probably substellar components have now come to light; dynamical masses should be obtained fairly soon \\citep{Close02, Potter02, Lane01}, allowing checks on the theoretical models for field brown dwarfs.  In young clusters and star-forming regions, however, no suitable systems have emerged yet.  This is especially troubling since even the identification of objects as substellar currently depends, at these early ages, on the theoretical tracks (empirical tests of substellarity that depend on Lithium detection or minimum Main Sequence temperature are largely inapplicable to very young objects).  Moreover, the low-mass tracks are most uncertain precisely at such early times \\citep{Baraffe02}, so testing them for young objects is particularly crucial.  To address this issue, we have developed a technique for calculating masses for young cluster objects from surface gravity measurements, independent of theoretical evolutionary models.  The essential idea is simple: derive surface gravity and effective temperature (\\teff) by comparing the observed spectrum to the latest synthetic ones, then derive radius (and extinction) by combining the observed photometry and known cluster distance with the surface fluxes predicted by the synthetic spectra (for the inferred \\teff and gravity), and finally derive mass by combining radius and gravity.  The sticking point, of course, is the derivation of sufficiently accurate surface gravities from the spectra, which has long been one of the major goals in the study of ultra-low mass objects.  However, we have shown in a previous paper (Mohanty et al. 2003a; henceforth Paper I) that the current generation of highly detailed synthetic spectra is equal to the task.  Employing these, we have derived gravities to within $\\pm$ 0.25 dex (and \\teff to within $\\pm$ 50K) in a sample of very low-mass objects in the Upper Scorpius and Taurus clusters (Paper I).  We now derive masses and radii for these, using our Paper I results together with photometry and distance estimates.  We will show that our analysis allows mass to be determined to within a factor of $\\sim$ 2, and radius to within $\\sim$ 30\\%.  These errors are much larger than those associated, for example, with dynamical mass and radius measurements in eclipsing binaries.  Nevertheless, we will demonstrate that they are sufficient for first order tests of the theoretical evolutionary tracks.  Though our analysis is independent of the evolutionary models (and thus serves as a check on the latter), it is clearly dependent on the validity of the synthetic spectra we use.  The accuracy of these were discussed in Paper I, and will be addressed further in this work.  However, we point out that our derivation of physical parameters using spectral synthesis alone does not constitute a great leap of faith, any more than employing evolutionary models for this purpose does.  There are two reasons for this.  First, the $P$-$T$ structure of the deep atmosphere (which forms the inner boundary of the spectral calculations) acts as the outer boundary condition of the interior calculations; i.e., the evolutionary models are anchored with the same (deep) atmospheric modeling as the synthetic spectra.  Second, in order to compare observations to the evolutionary predictions, synthetic spectra are crucial: either to convert an observed spectral type to \\teff (when placing objects on a theoretical \\teff-luminosity H-R diagram), or to convert predicted effective temperatures and luminosities to photometric colors and magnitudes (when placing objects on a theoretical color-magnitude diagram).  We have only taken the dependence on synthetic spectra a step further, by using them to derive surface gravities as well; the accuracy of our gravities is discussed at length in Paper I.  The advantage of this route lies in our avoiding (and thereby testing) what are perhaps the greatest uncertainties in the theoretical models for young objects: the initial conditions and still-relevant effects of accretion and collapse during the formation stage. The main disadvantage of eschewing the evolutionary models is that we cannot independently estimate ages.  This is compensated for by our ability to independently estimate the mass, and thus provide a check on the evolutionary model predictions for a group of objects that belong to the same cluster, and are thus likely to be (nearly) coeval.  Finally, we elucidate the system of nomenclature we have adopted here for very low-mass objects. There are considerable differences within the community at present regarding this issue: naming conventions based on fusion (or equivalently, mass), origins and location (isolated or in orbit around a star) have all been proposed.  Controversies arise because the different definitions do not yield the same grouping of objects; which of these systems is finally adopted is a matter for future arbitration.  However, given that a consensus is currently lacking, and that our primary concern in this paper is mass, we adopt a fusion-based convention (since mass is most directly associated with the presence and type of fusion).  The term `brown dwarf' refers to all objects which never derive 100\\% of their luminosity from hydrogen burning (unlike stars), but which are nevertheless above the deuterium-burning mass limit.  Thus brown dwarfs are objects in the range $\\sim$ 0.012--0.080 \\msun (12--80 \\mj).  We contract the term `planetary mass object' to the less cumbersome `planemo', and use it to refer to all objects below the deuterium-burning limit (i.e., mass $\\lesssim$12 \\mj), regardless of whether they are free-floating or in orbit around a star.  When a distinction is required between the two cases (e.g., when referring specifically to the recent isolated planemo candidates), it will be made explicitly.  Since neither planemos nor brown dwarfs undergo stable hydrogen fusion (i.e., reach the Main Sequence), both are included under the rubric of `substellar objects'.  Since both stars and brown dwarfs undergo at least some fusion, they will collectively be termed `fusors'; in this context, planemos are non-fusors.  A broader discussion of these terms and nomenclature issues is presented in Basri 2003. ",
        "conclusions": "  {\\it (1)} Both radius and \\teff decrease less rapidly with diminishing mass, at a given young age, than predicted by the theoretical evolutionary models.  Specifically, in the mass-radius plane the lowest mass objects ($\\lesssim$30 M$_J$) remain much larger (i.e., contract more slowly with age) than the models suggest, while the higher masses have radii in good agreement with the model predictions.  In the mass-\\teff plane, the higher masses are substantially cooler than predicted, while the lowest masses have \\teff either in better agreement with, or hotter than, the model values.  The combination of these two trends implies that luminosity also falls off less dramatically with mass, at a given age, than the evolutionary models indicate.  {\\it (2)} The lowest masses in our Upper Sco sample are near the deuterium fusion boundary.  Because of the importance of both conclusions, we have taken considerable pains to consider possible sources of error, both observational and systematic. These include conversion of colors to extinctions, temperature scales for pre-main sequence objects, problems with the gravity measurements, and the effects of starspots or binarity. Our extinctions are consistent with an analysis of the same region using low dispersion spectra. Our temperature scale is in good agreement with recent photometric work in the field. The range of masses we find within a few spectral subclasses is perhaps surprising, but we show that some of our basic conclusions can be drawn just from the observations (without recourse to theory at all).  We conduct a comparative analysis with a more extensively-studied young (GG Tau B) binary system to further test our conclusions, and find comparable discrepancies with theory in that case as well.  Finally, our derived relationships between radius, mass, \\teff and luminosity all agree (within the measurement uncertainties) with certain basic theoretical predictions that are likely to be correct regardless of evolutionary model uncertainties: younger objects have larger radii than older ones of the same mass; less massive objects are cooler and generally smaller than more massive ones at a given age; and luminosity decreases with both diminishing mass (at fixed age) and increasing age (at fixed mass).  There does not appear to be any a priori physical basis, therefore, for discarding our results.  We also point out that (like the Lyon models) our mass-radius, mass-\\teff, radius-\\teff and mass-luminosity relationships are smooth, without any sharp breaks or discontinuities.  The precipitous drop in gravity at low temperatures that we found in Paper I, which might seem remarkable at first sight, is due (if our analysis is correct) simply to a relatively slow change in radius and \\teff (compared to the Lyon predictions) over a significant range in mass.  The weight of the evidence suggests that substantially more work should go into the measurement of physical parameters of young substellar objects, the validity of the evolutionary tracks, and, without doubt, further testing and confirmation of our results.  An especially important conclusion of our work is that agreement with the evolutionary models in any single two-parameter plane (e.g., mass-radius) does not guarantee agreement in all parameters (e.g., \\teff, luminosity).  In order to ascertain the veracity of the models, their predictions must be checked for all the parameters, not just a selected few.  As a corollary, comparing an object to the evolutionary models over one set of parameters (e.g., \\teff-luminosity), in order to estimate other quantities (e.g., mass), is an exercise that is not always justified.  Such translations, which are common practice in current studies of young low-mass objects, may lead to spurious mass and radius estimates, and must be undertaken with great caution.  Similar conclusions have been reached by other authors, in the context of evolutionary model comparisons to higher-mass (solar-type) PMS stars (e.g., Torres \\& Ribas 2002).  In Paper I, we pointed out some specific areas of concern for theory, such as accretion effects and the treatment of convection and deuterium fusion.  In particular, we noted that if deuterium fusion begin at an earlier time than predicted, the discrepancies in radius and gravity between the theoretical tracks and our measurements, for the lowest masses, may be resolved.  This is a testable hypothesis, as we outlined in Paper I, and bears closer examination.  In this paper, we have also identified discrepancies in the theoretical \\teff predictions (assuming our derived temperatures are accurate), especially for the higher mass objects in our sample.  The underlying physical basis for temperature uncertainties in the evolutionary models is unclear; it is possible that remaining inadequacies in the treatment of convection are at fault.  Finally, while the model atmospheres and synthetic spectra that lie at the heart of our analysis are tremendously improved from earlier generations, they still suffer from certain shortcomings.  Specifically, they reproduce the photometry of field M dwarfs in some, but not all, of the optical and infrared bands.  While we have gone to great lengths to account for, and exclude, any attendant uncertainties in our analysis, further improvements in the atmospheric modeling - particularly in the linelists and opacities (most importantly, of H$_2$O) - would be tremendously useful for future studies of field and PMS low-mass objects.  We have implemented methods that have long been used for normal stars. They provide a means of testing theoretical isochrones and obtaining fundamental stellar parameters for very young, very low-mass objects. This methodology (which highlights the importance of high resolution spectroscopy and model atmospheres) should also be extended to higher mass objects and other star-forming regions with different ages. Extensive programs of this nature are now both desirable and feasible."
    },
    "0403/astro-ph0403119_arXiv.txt": {
        "abstract": "{ We present simultaneous multifrequency radio observations for a complete subsample of 26 XBLs from the {\\it Einstein} Extended Medium-Sensitivity Survey, obtained with the Very Large Array (VLA). Spectra are computed using fluxes at 20, 6 and 3.6 cm. Unlike many radio selected samples, the EMSS did not impose any criterion on the radio spectrum to identify BL Lac objects. It is therefore possible to investigate the intrinsic radio spectral slope distribution and to determine the effect produced by this selection criterion. We find that 15\\% of the observed objects do not meet the flat-spectrum criterion imposed on some other BL Lac samples. A dataset that includes non-simultaneous data (that are also taken with different VLA configurations) shows an even higher percentage of steep spectrum sources. This effect can be ascribed to a larger fraction of extended flux detected with the more compact VLA configuration. Possible biases introduced by the flat--radio-spectrum criterion in the radio-selected BL Lac samples cannot explain the discrepancies observed in the evolutionary properties of Radio and X-ray selected samples of BL Lacs. ",
        "introduction": "BL Lac objects are an enigmatic class of active galactic nuclei (AGN) characterized by strong radio, optical and X-ray variability, relatively high optical and radio polarization and featureless optical spectra (e.g. Urry \\& Padovani, 1995). The properties of the class can be explained in terms of the {\\it{relativistic beaming}} scenario wherein the observed emission is dominated by Doppler-boosted non-thermal radiation from a relativistic jet aligned with the line of sight (Blandford \\& Rees, 1978, Antonucci \\& Ulvestad, 1985).  Owing to their featureless spectra and lack of UV excess in many BL Lacs, optical search techniques (e.g. by excess colors, emission line strength), used to search for other AGN, failed to find BL Lacs in large number (e.g. Fleming et al., 1993). Since BL Lacs are both radio-loud and X-ray-loud, surveys in these frequency bands discover them with high efficiency (e.g. the ``1 Jansky'' sample: Stickel et al., 1991; the {\\it Einstein} Extended Medium-Sensitivity Survey (EMSS) sample: Morris et al., 1991; Rector et al., 2000; the Deep X-ray Radio Blazar Survey (DXRBS) sample: Perlman et al., 1998, Landt et al. 2001; the ROSAT All Sky Survey (RASS) sample: Bade et al., 1998; the ``sedentary'' survey: Giommi, Menna \\& Padovani, 1999; the Radio Emitting X-ray (REX) survey sample: Caccianiga et al., 1999, Caccianiga et al., 2002). This splits BL Lacs into two empirical subclasses, namely {\\it radio-selected BL Lacs} (RBLs) and {\\it X-ray-selected BL Lacs} (XBLs). Padovani \\& Giommi (1995) proposed to distinguish these two classes on a more physical basis, according to the spectral energy distribution (SED). For ease of description here we will refer to RBL and XBL only.  RBLs and XBLs show different behaviors with respect to properties like cosmological evolution, polarization, variability, core-to-extended radio flux ratio, extended radio morphology, spectral features (e.g. % Wolter et al., 1994; Stickel et al., 1991 Morris et al., 1991;  Jannuzi, Smith \\& Elston, 1993, 1994; Perlman \\& Stocke, 1993 to mention just a few) which have not been satisfactorily explained yet. This led to the suggestion that the observed discrepancies are at least partially caused by selection effects due to different criteria used in radio and X-ray surveys.  In particular, the standard selection technique for RBLs requires a flat-radio-spectrum criterion for which the radio spectral slope must be less than 0.5 ($\\alpha_r\\leq0.5$; $S_\\nu \\propto\\nu^{-\\alpha_r}$). This criterion has not been used for classical X-ray selected samples of BL Lacs like the EMSS. If surveys with this criterion select only part of the BL Lac population, the RBL samples could be biased and the differences between RBLs and XBLs like the ones mentioned above could be at least partially explained.  It is already known that a number of objects with steep spectra ($\\alpha_r>0.5$), classified as radio galaxies in radio surveys like the ``1 Jansky'' (Owen, Ledlow, \\& Keel, 1996, Perlman et al. 1996), otherwise meet the BL Lacs selection criteria, based on optical properties (equivalent width of emission lines; Ca~II break contrast) of March\\~a et al. (1996), and have broadband properties (radio luminosity and overall spectral energy distributions) that agree with those of BL Lacs (e.g. Rector et al., 2000).  To test if the flat--radio-spectrum criterion really introduces biases in the BL Lacs selection and to estimate the incompleteness degree of RBLs samples, we studied the radio spectral indices of a complete subsample of XBLs in which the identification criteria were not based on the radio spectral properties. Previous multi-frequency studies had already been performed on XBLs. For instance, Laurent-Muehleisen et al. (1993) collected non-\\-simultaneous data for several objects of this class, that however do not make a complete sample. We prefer to deal with the largest subsample for which we could have both simultaneous and non-simultaneous data.   The sources of this subsample are extracted from the EMSS (Gioia et al., 1990; Stocke et al., 1991), a % catalog whose BL Lacs sample has been thoroughly studied for completeness (Rector, Stocke \\& Perlman, 1999). A pilot study on the radio spectra of 8 XBLs from the EMSS and the {\\it HEAO-1 A-2} all-sky survey (Piccinotti et al., 1982) was performed by Stocke et al. (1985). Three objects showed a spectral slope which exceeds the limit for flat spectrum and two more were marginally steep. Preliminary analysis of the EMSS BL Lacs sample, by using non-simultaneous data at 6 and 20 cm from the EMSS and the {\\it NRAO VLA Sky Survey} (NVSS: Condon et al., 1998) respectively, showed that about 30\\% of the objects have a spectral slope steeper than 0.5. This is the motivation that led us to investigate the EMSS sources by obtaining simultaneous spectra, which are not affected by variability.  We briefly describe the sample of objects used for this work in \\S\\ref{The Sample}; in \\S\\ref{Data analysis} we present the data reduction process as well as flux densities and radio spectra for each object; in \\S\\ref{Results} we show the results obtained; our conclusions are summarized in \\S\\ref{Conclusions}. Throughout the paper we used $H_0=50$ km s$^{-1}$ Mpc$^{-1}$ and $q_0=0$, but no conclusions are dependent upon this choice of cosmology. ",
        "conclusions": "\\label{Conclusions} By using fluxes at 20, 6 and 3.6 cm, simultaneously measured with the VLA, we computed the radio spectra for a complete subsample of 22 XBLs from the EMSS (Gioia et al., 1990, Stocke et al., 1991). To this sample we added 4 sources whose simultaneous spectra were obtained by Stocke et al. (1985). The aim was to study the spectral slope distribution at radio frequencies without possible biases introduced by the flat-radio-spectrum criterion often used in selecting BL Lac samples. We found that about 15\\% of the sources have steep spectra.  We considered the archived non-simultaneous data as well, finding an even higher percentage of steep spectra.  This effect could be ascribed to variability, but other factors seem more significant in determining the measured slope of spectra. The higher percentage of steep spectra objects ($\\sim$ 38\\%), obtained by using the 20cm NVSS observations, implies the presence of a large fraction of extended flux that is preferentially detected by the VLA D configuration. This seems also supported by a possible correlation between spectral index and 20cm extended luminosities, computed either with simultaneous or non-\\-simultaneous data.   We find that the possible biases introduced by the flat--radio-spectrum criterion in the RBLs samples cannot easily account for the discrepancies observed in the $\\langle V_e/V_a \\rangle$ values of RBLs and XBLs samples and therefore the issue of evolution of different BL Lac classes is still open. The percentages of steep spectrum BL Lacs found for the EMSS sample cannot be applied straightforwardly to RBL samples, since a) the selection by radio spectral index is sometimes computed using non-simultaneous data; b) the XBLs properties are different from those of RBLs in the sense that RBLs are more variable and possibly more core dominated. If the observed spread in $\\alpha_r$ is mostly due to variability, we should expect a larger fraction of steep RBLs than XBLs. If most of the effect is due, as the data seem to suggest, to the presence of a significant extended component, that is lesser in RBL, the influence on the radio selected objects should be less severe.  We stress here that the $\\alpha_r$ distribution in the EMSS BL Lacs simultaneous data is a continuous one and so any division in two subsamples is somewhat arbitrary and does not have a physical underlying support.  We are tempted to suggest a different cut-off value of $\\alpha_r=0.7$ (as done e.g. by Perlman et al. 1998 in DXRBS) that would then exclude $\\leq$ 10\\% of all BL Lacs based upon our findings here. But, because this would also include a very large number of ``normal'' radio galaxies, some of which with optical featureless spectra (see e.g. discussions in Rector \\& Stocke, 2001, Perlman et al. 1996), other information would have to be used to define them as BL Lacs (e.g. variable optical polarization).   However, we have shown that the selection effects applied by using a cut-off value in the radio spectral index of $\\alpha_r=0.5$ are not so severe. The steeper BL Lacs that are excluded from this criterion seem to have the same overall properties of the flat ones, especially for what concerns the cosmological distribution. We can therefore conclude that the discrepancies found by using different samples cannot be ascribed to the flat radio spectrum criterion."
    },
    "0403/astro-ph0403433_arXiv.txt": {
        "abstract": "Given the scientific goals of VSOP-2, including the possibility of observations of the shadows of black holes, we have investigated the fidelity of the recovered images given a typical {\\em uv}-coverage. We find that we can achieve a dynamic range of better than 1000:1. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403009_arXiv.txt": {
        "abstract": "{With the advent of the present and future spatial X-ray missions, it becomes crucial to model correctly the line spectrum of X-ray emitting media such as the photoionized plasma observed in the central regions of Active Galactic Nuclei (AGN), or in X-ray binaries. We have built a photoionization code,  Titan, solving the transfer of a thousand lines and of the continuum with the ``Accelerated Lambda Iteration\" method, which is one of the most efficient and at the same time the most reliable for line transfer. In all other photoionization codes the line intensities are computed with the so-called ``escape probability formalism\", used in its simplest approximation. In a previous paper (Dumont et al. 2003), it was shown that this approximation leads to a wrong estimation of the emitted X-ray line intensities, especially in the soft X-ray range. The errors can exceed one order of magnitude in the case of thick media (Thomson thickness of the order of unity). In the present paper, we show that it also happens, but for different reasons, in the case of moderately thin media (Thomson thickness of 0.001 to 0.1), characteristic of the Warm Absorber in Seyfert 1 or of the X-ray emitting medium in Seyfert 2. Typically, the errors on the line fluxes and line ratios are of the order of 30$\\%$ for a column density of 10$^{20}$ cm$^{-2}$, and a factor five for a column density of 10$^{23}$ cm$^{-2}$, in conditions giving rise to the spectra observed in these objects. We explain why this problem is less acute in cooler media, like the Broad Line Region of AGN. We show some examples of X-ray spectra appropriate for Seyfert 2 and for the Warm Absorber of Seyfert 1.  We conclude that though it is quite important to introduce numerous accurate X-ray data in photoionization codes, it should be accompanied by more elaborate methods than escape probability approximations to solve the line transfer. ",
        "introduction": "Since more than three decades, photoionization codes have been developed to compute the structure and the spectrum of photoionized media, such as HII regions, planetary nebulae, supernova remnants, envelopes of novae, Narrow Line Regions of  Active Galactic Nuclei (AGN), etc\\ldots At the end of the seventies, these codes have been extended to denser and thicker media, like the Broad Emission Line Region (BLR) in quasars and AGN or the emission regions of X-ray binaries and cataclysmic variables. For this purpose, the formalism of ``escape probability\" has been introduced to take into account self-absorption in lines while avoiding to solve the line transfer (Netzer 1975), and a new type of photoionization code has begun to be built. These codes - Cloudy (Ferland et al. 1988), XSTAR (Kallman \\& Krolik 1995), ION (Netzer 1993), for instance - have now reached a high degree of sophistication, including very accurate atomic data, a large number of ions and transitions, and all the necessary processes allowing to use them in various physical conditions.  With the advent of the X-ray missions Chandra and XMM Newton, splendid spectra of various types of objects have been obtained in the soft X-ray range, showing tens of emission lines which can be used as diagnostics of the physical state of the emitting region. The best examples are Seyfert 2 galaxies, which display a rich X-ray line spectrum, most probably produced by the external part of the ``Warm Absorber\" of Seyfert 1, photoionized by the intense central continuum and seen in emission because the central continuum is hidden from our view (Antonucci \\& Miller 1985). Typical column densities of this medium are 10$^{21-23}$ cm$^{-2}$ (Sako et al. 2000, Kinkhabwala et al. 2002, Ogle et al. 2003).  Most naturally, the photoionized codes developed for the BLR have been used to model the X-ray emitting regions of different objects, like the atmosphere of cataclysmic variables, X-ray binaries, etc\\ldots. These codes are also used to model the absorption lines observed in the UV and X-ray range in quasars and AGN. Numerous new X-ray atomic data have been introduced in the codes, in order to obtain the best possible accuracy on the X-ray spectrum. But the ``escape probability approximation\" still lies at the center of the computation of the line intensities.  This formalism, developed in the sixties and in the seventies, can be very useful to perform rapid approximate computations, but it does not lead to a correct estimation of the source function, especially in the case of strongly interlocked transitions, including continuum ones. Moreover it uses as a local quantity a global one, computed by an integration over the whole medium; this is a dramatic extrapolation for very inhomogeneous media like photoionized plasma. These aspects have been completely overlooked during the last twenty years, but they become now crucial in the context of X-ray emitting media.  In a previous paper, Dumont et al. (2003, referred as D03)  have shown that escape probability approximations, at least as they are used in the present codes, are unable to compute correct line intensities, within factors of ten, when the Thomson thickness of the medium is of the order of a few units, typical for the irradiated atmospheres of accretion discs in AGN and X-ray binaries. In the present paper, we extend this study to less optically thick media, and we show that the intensities of emission lines are also not accurately computed for parameters typical of the Warm Absorber of Seyfert 1 galaxies and of the region giving rise to the X-ray spectrum in Seyfert 2 galaxies \\footnote{One should note that modeling an absorption spectrum is much easier than an emission one. It requires only a correct computation of the thermal and ionization equilibrium, which give the fractional ion abundances, i.e. the populations of the ground levels, and thus the equivalent widths of the resonance lines, which can be compared to those deduced from the observations through a curve of growth analysis implying no line transfer. However it can happen (and this is indeed the case for the Warm Absorber), that emission lines are also produced by the absorbing medium, and one should thus worry about the line transfer.}.  In the next section, we recall very briefly the essence of the problem, and we give a few examples in Section 3. ",
        "conclusions": "We have shown that the use of the escape probability approximation leads to large errors in the computed line fluxes and line ratios, in conditions which are typical of the X-ray emitting regions of Seyfert 2 nuclei, and of the Warm Absorber of Seyfert 1, i.e. for a Thomson thickness of the order of 0.001 or larger. This completes the previous paper (D03), where the same study was performed for thicker media, and where it was shown that the use of the escape approximation leads to errors by more than one order of magnitude on the line fluxes and line ratios. We find here that the errors are of the order of 30$\\%$ for a column density of 10$^{20}$ cm$^{-2}$, and can reach a factor five for a column density 10$^{23}$ cm$^{-2}$. We confirm that they are almost always in the direction of an overestimation of the most intense line intensities, especially of the L$\\alpha$ line of H-like ions, and of the resonance w term of He-like ions. We explain why such large errors occur for X-ray emitting media, and not for cooler media like the BLR.  The comparison between the escape approximation and the transfer treatment (performed through the Accelerated Lambda Iteration method with our photoionization code Titan), is made in such a way that no other possible explanation of the discrepancies than the use of the escape probability approximation can be invoked. So one is led to conclude that unless a real transfer of the lines is introduced in the codes for modelling X-ray spectra, the results cannot have an accuracy better than that given by the approximation, even when the treatment of atomic physics is highly sophisticated.  Whatever the discrepancy between the results of the escape probability and the transfer treatment, an uncertainty remains concerning the real intensity of the resonance lines. Here they were computed assuming complete redistribution within Doppler core, which mimics partial redistribution within a Voigt profile. PRD cannot be taken into account accurately with the escape probability formalism, but it can be done with the transfer treatment, provided the implementation in the code of another substantial time consuming procedure. We are presently studying such an improvement.  \\bigskip  \\noindent {\\bf APPENDIX: Equations used for the escape approximation in the optically thin case}  We compute the escape probability towards the surface $P_{e}(\\tau_0)$ of a line  as (cf. Collin-Souffrin et al. 1981):  \\begin{equation} P_{e}(\\tau_0) ={\\rm max}(f1,f2), \\label{eq-esc-1} \\end{equation} for all subordinate and high resonance lines, and \\begin{equation} P_{e}(\\tau_0) =f1 \\label{eq-esc-2} \\end{equation} for the first resonance lines of H- and He-like species, with \\begin{eqnarray} f1 ={0.5\\over 1+2\\tau_0\\sqrt{\\pi ln(\\tau_0+1)}} && f2= {2\\over 3}\\sqrt{{a\\over\\sqrt{\\pi}\\tau_0}}, \\label{eq-esc-3} \\end{eqnarray} where $\\tau_0$ is the optical thickness at the line center between the emission point and the surface (we recall that it is taken along the photon path, so it multiplied by a factor $\\sqrt{3}$ to take into account the fact that the direction of the emitted photon is at random, and to be able to compare the escape approximation to the transfer treatment), and $a$ is the usual damping constant. $f1$ corresponds to the Doppler core, and $f2$ to the Lorentz wings of the Voigt profile.  The total escape probability $P_{\\rm esc}$ is the sum of the escape probability towards the illuminated side, $\\beta_{\\rm ref}=P_{\\rm e}(\\tau_0)$,  and towards the back side, $\\beta_{\\rm out}=P_{\\rm e}(T_0-\\tau_0)$, where $T_0$ is the total optical thickness of the slab depth at the line center.  In the equations for the level populations, the net radiative rate from an excited level is replaced by $A_{\\rm ul}\\beta_{\\rm pop}$, with \\begin{equation} \\beta_{\\rm pop}={\\rm min}[1, P_{\\rm esc}(\\tau) \\times (1+{\\kappa_{\\rm c}\\over \\kappa_l}F({\\kappa_{\\rm c}\\over \\sqrt{\\kappa_{\\rm l}}}))], \\label{eq-betapop} \\end{equation} where $\\kappa_{\\rm c}$ and $\\kappa_{\\rm l}$ are respectively the absorption coefficient in the continuum and in the line, and $ F$ is the operator given by Hummer (1968) to account for destruction by continuum absorption in one line scattering: \\begin{equation} F(X)=\\int_{-\\infty}^\\infty {\\phi (x)\\over X+\\phi (x)} dx \\label{eq-FX} \\end{equation} where $\\phi (x)$ is the absorption line profile, $x=\\delta \\nu /\\delta \\nu _{\\rm D}$, and $\\delta \\nu _{\\rm D}$ is the Doppler width.  A term $B_{\\rm lu}J_{\\nu}^{\\rm inc.att.}$ is added to the net radiative rate to take into account excitations by the attenuated incident radiation $J_{\\nu}^{\\rm inc.att.}=2\\beta_{\\rm ref}/(2\\pi)F_{\\nu}^{\\rm inc}$, according to the definition of the flux, $F_{\\nu}^{\\rm  inc}$ being the incident flux at the frequency $\\nu$ (the corresponding deexcitation rate is negligible).  The local cooling for each line is: \\begin{equation} \\Lambda_{\\rm line}=(n_{\\rm u}A_{\\rm ul} P_{\\rm esc}-n_{\\rm l}B_{\\rm lu}J_{\\nu}^{\\rm inc.att.}) \\ {h\\nu\\over n_{\\rm e}n_{\\rm H}} . \\label{eq-lambda-line} \\end{equation}  The reflected flux in a line is computed as: \\begin{equation} F_{\\rm ref}= \\int{n_{\\rm u}A_{\\rm ul}h\\nu \\beta_{\\rm ref}\\ {\\rm exp}[-\\tau_{\\rm e}] dz}, \\label{eq-Fref} \\end{equation} where $\\tau_{\\rm e}$ is the effective optical thickness  of the slab in the continuum at the line frequency, between the current point and the illuminated surface.  In the outward emitted line flux, one must take into account the photons absorbed in the incident continuum: \\begin{eqnarray} F_{\\rm out}&=& \\int{(n_{\\rm u}A_{\\rm ul}h\\nu \\beta_{\\rm out} {\\rm exp}[-(T_{\\rm e}-\\tau_{\\rm e})]} \\\\ \\nonumber && \\ \\ \\ -n_{\\rm l}B_{\\rm lu}J_{\\nu}^{\\rm inc.att.}) dz, \\label{eq-Fout} \\end{eqnarray} where $T_{\\rm e}$ is the effective total optical thickness of the slab in the continuum at the line frequency.  The ionization rate due to the lines, at the depth $z$, is equal to: \\begin{eqnarray} &&\\sqrt{3}\\kappa_{\\rm c} A_{\\rm ul}\\times \\\\ \\nonumber && \\left( \\int_0^z{(n_{\\rm u} \\beta_{\\rm out}{\\rm exp}[-\\tau_{\\rm e}+T_{\\rm e}]-n_{\\rm l}{\\rm B_{\\rm lu}\\over A_{\\rm ul}}J_{\\nu}^{\\rm inc.att.})} dZ \\right. \\\\ \\nonumber &&\\left. + \\int_H^z{n_{\\rm u} \\beta_{\\rm ref}{\\rm exp}[-T_{\\rm e}+\\tau_{\\rm e}]dZ}\\right). \\label{eq-ion} \\end{eqnarray} This expression is used also for the gains by photoionizations due to the lines.  The escape probability is slightly different in the case of a highly ionized and/or moderately thick medium, as one must take into account the shift of line photons by comptonization. Thus $\\beta_{\\rm ref}$ is replaced by $\\beta'_{\\rm ref}= \\beta_{\\rm ref}+{1\\over 2} (1-P_{\\rm esc}(\\tau)) {\\sigma\\over \\kappa_l\\sqrt{\\pi}+\\kappa_c+\\sigma}$, and $\\beta_{\\rm out}$ by $\\beta'_{\\rm out}= \\beta_{\\rm out}+ {1\\over 2} (1-P_{\\rm esc}(\\tau)) {\\sigma\\over \\kappa_l\\sqrt{\\pi}+\\kappa_c+\\sigma}$, and $P_{\\rm e}$ is replaced by $P'_{\\rm e}=\\beta'_{\\rm ref}+\\beta'_{\\rm out}$."
    },
    "0403/astro-ph0403523_arXiv.txt": {
        "abstract": "We present the long term X-ray light curves, detailed spectral and timing analyses of XTE J1908+094 using the Rossi X-ray Timing Explorer Proportional Counter Array observations covering two outbursts in 2002 and early 2003. At the onset of the first outburst, the source was found in a spectrally low/hard state lasting for $\\sim$40 days, followed by a three day long transition to the high/soft state. The source flux (in 2$-$10 keV) reached $\\sim$100 mCrab on 2002 April 6, then decayed rapidly. In power spectra, we detect strong band-limited noise and varying low-frequency quasi periodic oscillations that evolved from $\\sim$0.5 Hz to $\\sim$5 Hz during the initial low/hard state of the source. We find that the second outburst closely resembled the spectral evolution of the first. The X-ray transient's overall outburst characteristics lead us to classify XTE J1908+094 as a black-hole candidate. Here we also derive precise X-ray position of the source using Chandra observations which were performed during the decay phase of the first outburst and following the second outburst. ",
        "introduction": "The X-ray transient source \\xsrc was serendipitously discovered on 2002 February 19 during scheduled Rossi X-ray Timing Explorer (RXTE) Proportional Counter Array (PCA) observations of a Soft Gamma Repeater, SGR 1900+14 (Woods et al. 2002). Subsequent RXTE/PCA scanning observations of the region allowed the localization of the new source to RA: 19h08m50s, Dec: +09$^\\circ$22$\\arcmin$ 30$\\arcsec$ with an accuracy of 2$\\arcmin$. This placed the new source about 24$\\arcmin$ away from SGR 1900+14. Based on Very Large Array observations on 2002 March 21 and 22, a transient source was suggested as the radio counterpart candidate to \\xsrc (Rupen, Dhawan \\& Mioduszewski 2002a). Observations in the optical band on April 8 and 9 revealed no new sources near the radio position (Garnavich, Quinn \\& Callanan 2002), however Chaty, Mignani \\& Israel (2002) identified a near infrared counterpart to the new source and concluded that \\xsrc is in a low-mass X-ray binary system with a main sequence companion of spectral type later than K. Based on our preliminary analysis of the RXTE/PCA observations we concluded that \\xsrc is a new stellar mass black hole candidate (Woods et al. 2002).  There are currently 18 dynamically confirmed and 20 candidate stellar mass black hole systems in our Galaxy (see McClintock \\& Remillard 2003 for a recent review). Most of them are characterized by occasional transient outbursts (X-ray novae; Chen, Shrader \\& Livio 1997) as a result of sudden increase in the mass accretion rate possibly triggered by instabilities in the accretion disk (Cannizzo 1993, Dubus et al. 2001). During outbursts, these systems generally undergo various changes in their spectral characteristics, usually  in conjunction with changes in their timing behavior (see e.g., Homan et al. 2001). The most common BH spectral states are {\\it the low/hard state:} the spectrum is represented by a hard power law and usually accompanied by timing variability, and {\\it the thermal-dominant (high/soft) state:} a blackbody appears in the spectrum as the power law component gets steeper and timing features get weaker or completely disappear (Tanaka \\& Lewin 1995). Other spectral states characterized by more complicated spectral and timing properties are also observed (e.g., Homan et al. 2001).  In this study, we present the results of our spectral and timing analysis of the RXTE pointed observations of \\xsrc covering two outburst episodes. Additionally, we report observations with the Chandra X-Ray Observatory taken under Director's Discretionary Time (DDT). We describe our observations in \\S 2, and we present detailed data analyses in \\S 3. In \\S 4 we discuss and compare the different states of the source to other black hole candidates. ",
        "conclusions": "In the course of its X-ray activity XTE J1908+094 proceeds through a series of X-ray states characteristic of black hole binaries, which we shall discuss using the terminology proposed by McClintock  \\& Remillard (2003). The discovery outburst begins in the low/hard state, which lasted until about day 87. The source then enters an intermediate state during which the 2$-$5 keV flux peaks. This ends near  day 90 when the source enters the thermal-dominant (high/soft) state, which persists until  day 143, just before the secondary peak, where the source enters an intermediate state and then after  day 149 returns to the low/hard state. The remaining observations (in 1.5$-$12 keV) up to the peak of the second outburst show the source in the low/hard state (see Figure \\ref{fig:rates_obss}). The one observation following the second outburst peak is consistent with the thermal-dominant state.   In the low/hard state the energy spectra are dominated by a hard power-law component, and the power spectra by strong band-limited noise. During the low/hard state at the onset of the first outburst, the index of the power-law spectra began at $\\Gamma=1.4$, and then gradually softened to $\\Gamma=1.7$. The power spectra show band-limited noise with an rms amplitude which began near $r=30$\\% and gradually fell to $r=22$\\%. In addition there was a QPO with rms amplitude varying from 3\\% to 13\\%, which rose in frequency from 0.5 to 2.2 Hz. In the second interval of low/hard state ( days 149$-$295) the flux is lower, and the behavior of the power-law index is more complex. Due to the low flux, we could not make significant power-spectral measurements.  Outburst onsets in the low/hard state have been seen in a number of X-ray novae. Brocksopp et al. (2002) tabulate 13 sources with outbursts that began in the low/hard state, five of which never left this state. Strong low-frequency QPO with rising frequencies are common in these low/hard state onsets, and have been seen for GRO J0422+32 \\cite{vanderHooft99}, GRO J1719-24 \\cite{vanderHooft96}, XTE J1550-564 (Finger et al. 1998; Cui et al. 1999), 4U 1630-472 \\cite{Dieters00}, XTE J1859+226 \\cite{Markwardt99}, and XTE J1118+480 \\cite{Wood00}, among others.  \\begin{figure}[!b] \\centerline{\\includegraphics[scale=.50]{f16.eps}} \\vspace{0.0in} \\caption{\\baselineskip =0.5\\baselineskip The RXTE/PCA spectral fit for day 64 shown together with the BeppoSAX/PDA spectral fit for days 63$-$65 from in't Zand et al. (2002). \\label{fig:PDAspec}} \\end{figure}  Hard-X-ray and gamma-ray observations have shown that in the low/hard state the power-law spectra break in the 100 keV range \\cite{Grove98}. In Figure \\ref{fig:PDAspec} we show our spectral fit for the PCA data on day 64 along with the spectral fit for BeppoSAX/PDA data from days 62$-$65 (MJD 52342$-$52345) \\cite{intZand02}. We notice a break near 50 keV. The flux in the 30-250 keV range is $3.2\\times 10^{-9}~{\\rm erg~cm}^{-2}~{\\rm s}^{-1}$, which surpasses the flux in the 2.5-25 keV range of $2.63\\times 10^{-9}~{\\rm erg~cm}^{-2}~{\\rm s}^{-1}$.  The low/hard state is also associated with radio emission. Flat spectrum radio emission, associated with compact jets, is consistently observed during the low/hard state of X-ray novae \\cite{Fender03}. Indeed, during the onset of the first XTE J1908+094 outburst, Very Large Array observations (on days 74$-$75) led to the discovery of a radio counterpart to XTE J1908+094, with a flux of 0.85 mJy at 8.6 GHz \\cite{Rupen02a}. This was detected in additional observations until day 127 \\cite{Rupen02b}.   \\begin{figure}[!b] \\centerline{\\includegraphics[scale=.50]{f17.eps}} \\vspace{0.0in} \\caption{\\baselineskip =0.5\\baselineskip Flux, Power-law index, and variability evolution during the intermediate state. The top panel shows the 2.5-25 keV flux for the power-law (filled triangles), the disk black-body component (filled circles), and their sum (dashed). The middle panel shows the power-law index. The lower panel shows the variability amplitude.\\label{fig:IMstate}} \\end{figure}  Figure \\ref{fig:IMstate} shows the X-ray flux, power-law index and variability amplitude evolution during the transition interval between the low/hard and thermal-dominant state. This intermediate state begins near day 67, when the power-law flux begins to drop, the power-law component begins to rapidly soften, and the disk black-body flux begins to rise, and the flux variability begins to fall. The rise of the disk black-body flux occurs in four days, but the fall of the power-law flux takes 15 days to complete. BeppoSAX/MECS observations covering day 66.4 $-$67.8 (MJD 52366.4$-$52367.8) show the onset of this transition \\cite{intZand02}.  In this transition, the total flux in 2.5$-$25 keV band consistently falls. However, a significant fraction of the disk black-body flux is below this energy range. From our spectral fits we find that the total integrated flux of this thermal component rises to $\\sim 7\\times 10^{-9}~{\\rm erg~cm}^{-2}~{\\rm s}^{-1}$, implying that the bolometric flux may be constant or rising.   After day 90 the source is in the thermal-dominant state, with thermal disk flux dominating the spectrum, and low variability. The disk black-body normalizations average about 40, which is consistent with the inner disk being at the radius of the inner most stable circular orbit if \\begin{equation} (M/M_\\odot)D_{10kpc}^{-1}\\cos^{\\onehalf}\\theta \\approx 4.5 \\end{equation} where $M$ is the black-hole mass,  $D_{10kpc}$ the source distance in units of 10 kpc, and $\\theta$ the disk inclination to the line of sight. Note here that for any value of disk inclination angle, the mass of the central object is in the range of a black hole, if the source distance is of the order of 10 kpc.  During this thermal-dominant state, the flux steadily falls. This mainly occurs by the temperature decreasing. The thermal-dominant state ends on day 149 when a transition begins back to the low/hard state.   In the transition from the low/hard state to the thermal-dominant state starting day 87), the disk black-body normalization begins near 15 and rises, implying an increasing inner disk radius. The opposite occurs on the transition back to the low/hard state. This is counter to the expectation that the inner-disk radius is large during the low/hard state, and near the inner most stable orbit in the thermal-dominant state (e.g., Esin et al. 1997). This rise and fall may be due to systematic problems with our spectral fits: we detect only the high-energy tail of the thermal spectrum. In fits where the column density, disk temperature and flux are all free to vary, these parameters are, therefore, highly coupled. By fixing the column density to the value found with the BeppoSAX data, we have reduced this coupling, but could be biasing the solution.   There is a strong correlation between the flux associated with the broad line feature and that of the power law component during the early stages of the first outburst episode. One possible interpretation of this feature is that it is the fluorescent Fe K$\\alpha$ emission produced by the reprocessing of the hard X-ray photons by cooler material close to the central object. The line centroid energies were somewhat lower than what is expected for neutral iron (6.4 keV). This may suggest that what we observe is primarily the red wing of Doppler shifted neutral Fe K$\\alpha$ in a Keplerian accretion. This was seen also in 4U 1630$-$47 (Cui, Chen \\& Zhang 1999) and in XTE J1748$-$288 (Miller et al. 2001).   Esin et al. (1997) have presented a model to explain the states of X-ray novae. In the quiescent and low/hard state a thin accretion disk is present but truncated at a large inner radius. Within this radius there is an Advection Dominated Accretion Flow (ADAF) which is a hot and radiatively inefficient flow where most of the thermal energy generated is advected onto the black hole rather then being radiated. Above the accretion disk is a hot corona, which is a continuation of the advection dominated flow, which produces a power-law spectral component in the low/hard state by Comptonization.  This model does not incorporate the jets which are responsible for the radio emission now known to be associated with the low/hard state. Markoff, Falcke \\& Fender (2001) have proposed that these jets also produce the power-law component via synchrotron radiation. In their model a standard accretion disk transitions at an inner radius of $\\sim 10^3$ km to a hot ADAF-like flow which feeds the jet. The hard X-rays are synchrotron radiation produced in a shock acceleration region some $10^3$ km above the disk plane. The radio emission is from beyond this region.  While providing successful fits of a multi-wavelength spectrum, neither of these models yet consider dynamical changes in the flow or attempt to explain the power-spectra seen in the different states. The high amplitude variability seen in the low-hard state requires changes in emissivity that are spatially coherent over most of the emission region. It is tempting to associate the QPO's seen in the low/hard state with the Keplerian  frequency at the inner edge of the thin accretion disk. Yet while this is plausible in the low-hard state, the QPO in XTE J1908+094 persists into the intermediate state, where the inner disk radius inferred from spectral fits imply frequencies much larger than those observed."
    },
    "0403/astro-ph0403715_arXiv.txt": {
        "abstract": "We present spectroscopic confirmation of the Cl 1604 supercluster at $z \\sim 0.9$. Originally detected as two individual clusters, Cl 1604+4304 at $z = 0.90$ and Cl 1604+4321 at $z = 0.92$, which are closely separated in both redshift and sky position, subsequent imaging revealed a complex of red galaxies bridging the two clusters, suggesting that the region contained a large scale structure. We have carried out extensive multi-object spectroscopy, which, combined with previous measurements, provides $\\sim600$ redshifts in this area, including 230 confirmed supercluster members. We detect two additional clusters that are part of this structure, Cl 1604+4314 at $z = 0.87$ and Cl 1604+4316 at $z = 0.94$.  All four have properties typical of local clusters, with line-of-sight velocity dispersions between 489 and 962~{\\kms}. The structure is significantly extended in redshift space, which, if interpreted as a true elongation in real space, implies a depth of $93~h_{70}^{-1}~{\\rm Mpc}$. We examine the spatial and redshift distribution of the supercluster members. ",
        "introduction": "The Cl 1604 supercluster was initially detected as two separate clusters, Cl 1604+4304 at $z=0.90$ and Cl 1604+4321 at $z=0.92$, in the plate-based survey of \\citet{gho86}. Deeper imaging and multi-object spectroscopy taken with the Low Resolution Imaging Spectrograph \\citep[LRIS;][]{oke95} at the Keck telescopes by \\citet[hereafter O98]{oke98} yielded redshifts and preliminary velocity dispersions for each cluster \\citep[][hereafter P98,P01]{pos98,pos01}. Motivated by the close separation of Cl 1604+4304 and Cl 1604+4321 in both radial velocity (4300~{\\kms}) and position on the sky ($17'$), \\citet{lub00} performed deep multi-band imaging with the Palomar 5-m telescope covering the area between the two clusters. The imaging revealed an overdensity of red galaxies whose colors were consistent with early-type galaxies at $z \\sim 0.9$, suggesting that the clusters were part of a high-redshift supercluster.  To verify this conclusion, map the supercluster structure, and obtain a large sample of cluster galaxies with measured spectroscopic properties, we are conducting an extensive spectroscopic survey spanning the region between Cl 1604+4304 and Cl 1604+4321. In this Letter, we present spectroscopic confirmation of this large scale structure and a discussion of its properties based on the currently known supercluster members. We assume a $\\Lambda$CDM cosmology with $\\Omega_m=0.3, \\Lambda=0.7,$ and ${\\rm H_0}=70~h_{70}~{\\rm km~s^{-1}~Mpc^{-1}}$. ",
        "conclusions": "We have obtained the largest spectroscopic sample to date for a supercluster at $z\\sim0.9$. The initial results suggest that the four component clusters are part of a single, filamentary structure, with the individual clusters being massive (Abell Richness class 0--2). As seen in Figure~\\ref{gal_posns}, we may not yet have reached the outer limits of this structure, and future wide-area imaging and additional spectroscopy will be required to measure its full extent.  Already our dataset is comparable in size to studies of local superclusters (see \\S3.2). This structure is by far the best studied of the very few known superclusters at $z>0.3$ \\citep{con96,ros99,gav04}.  The spectroscopic data, especially from DEIMOS, are of sufficient resolution and quality to measure equivalent widths for both emission and absorption features. A large fraction of the galaxies in the cluster show evidence for star-formation (most notably strong [OII] emission). Examining the correlation of such features with galaxy location will yield insights into the dynamics and history of the supercluster system and shed light on the role of cluster-scale physical processes (e.g., shocks and ram-pressure stripping) in galaxy star formation. Proposed HST imaging and Chandra X-ray mapping of this structure will allow us to compare galaxy morphological and spectroscopic properties and correlate optical substructure with the location and temperature of the hot gas component. These data, combined with numerical simulations, may be able to predict future mergers and perhaps even the end-state of this system.  This survey will provide a view of a young supercluster rivaling the detailed studies of local structures."
    },
    "0403/astro-ph0403465_arXiv.txt": {
        "abstract": "Hard X-ray surveys have uncovered a large population of heavily obscured AGN. They also reveal a population of quasars with moderate obscuration at both visible and X-ray wavelengths. We use \\chandra selected samples of quasars from the ELAIS Deep X-ray Survey (EDXS) and the \\chandra Deep Field-North to investigate the obscuration towards the nuclei of moderately obscured AGN. We find an inverse correlation between the optical to X-ray flux ratio and the X-ray hardness ratio which can be interpreted as due to obscuration at visible and X-ray wavelengths. We present detailed optical and near-infrared data for a sample of optically-faint ($R>23$) quasars from the EDXS. These are used to constrain the amount of rest-frame UV/optical reddening towards these quasars. It is found that optically-faint quasars are mostly faint due to obscuration, not because they are intrinsically weak. After correcting for reddening, the optical magnitudes of most of these quasars are similar to the brighter quasars at these X-ray fluxes. Combining with gas column densities inferred from the X-ray observations we consider the gas-to-dust ratios of the obscuring matter. We find that the quasars generally have higher gas-to-dust absorption than that seen in the Milky Way -- similar to what has been found for nearby Seyfert galaxies. We consider the possible existence of a large population of X-ray sources which have optical properties of Type 1 (unobscured) quasars, but X-ray properties of Type 2 (obscured) quasars. However, we find that such sources only contribute about 6\\% of the 0.5-8 keV X-ray background. Finally we show that the observed distribution of optical-to-X-ray flux ratios of quasars at $z>1$ is skewed to low values compared to the intrinsic distribution due to the fact that the observed-frame $R$-band light is emitted in the UV and is more easily obscured than hard X-rays. ",
        "introduction": "The hard spectral shape of the X-ray background led to the idea that a large population of obscured active galactic nuclei (AGN) exist which fail to show up in optically-selected quasar surveys (Comastri et al. 1995).  Deep hard X-ray surveys with \\chandra and \\xmm have now revealed the sources responsible for the 0.5-10 keV X-ray background (e.g. Hornschemeier et al. 2001; Tozzi et al. 2001; Hasinger et al. 2001; Manners et al. 2003). As expected, the majority of the optical counterparts to hard X-ray sources are galaxies containing optically-obscured AGN.  Seyfert galaxies at low redshift often show complex absorption structures consisting of both cold (neutral) and warm (partially ionized) absorbers (Mushotzky, Done \\& Pounds 1993). A fraction of Seyferts appear to be Compton thick and their observed X-ray emission is reflection dominated. Due to the low signal-to-noise of the X-ray spectra of the sources responsible for the X-ray background, there is only limited knowledge of their absorption properties. It is unknown whether the orientation-based obscuration scheme which works well for low-redshift Seyferts can also be applied at higher redshifts. There are still only a handful of Type 2 narrow-line quasars found in deep X-ray surveys (Almaini et al. 1995; Norman et al. 2002; Stern et al. 2002; Crawford et al. 2002; Mainieri et al. 2002; Szokoly et al. 2004), seemingly at odds with simple unification schemes (however there are still many faint objects without redshifts in these surveys which could be Type 2 quasars). The relative lack of luminous Type 2 quasars and the dominance of obscured AGN at lower luminosities (Ueda et al. 2003; Hasinger et al. 2003) suggests that the fraction of obscured objects is a strong function of luminosity, as had previously been found for low-radio-frequency selected AGN (Simpson, Rawlings \\& Lacy 1999; Willott et al. 2000). The fraction of Compton thick AGN in the X-ray background sources is also quite uncertain (e.g. Fabian, Wilman \\& Crawford 2002).   The dust properties of the high redshift X-ray absorbing material have not yet been well studied. There are certainly objects at low redshift with dust absorption quite different from that expected from the observed X-ray absorption based on a Galactic gas-to-dust ratio (Simpson 1998; Maiolini et al. 2001a) and a few cases of such discrepancies at higher redshifts have been reported (Akiyama et al. 2000; Risaliti et al. 2001; Willott et al. 2003; Watanabe et al. 2004). In this paper we discuss the dust and X-ray absorption present in hard X-ray sources contained within the ELAIS Deep X-ray Survey and the implications for understanding the sources responsible for the X-ray background.  In Sec.\\,2 we discuss the observed correlations between optical and X-ray properties of X-ray selected quasars and what these suggest about obscuration. Sec.\\,3 presents near-infrared and optical data for optically-faint X-ray selected quasars. In Sec.\\,4 we fit model quasar spectra to the observations to constrain the reddening towards these quasars. In Sec.\\,5 we compare the obscuration in the UV/optical with that in X-rays to determine the gas-to-dust ratio of the obscuring material and compare this with values in the Milky Way, the Small Magellanic Cloud and low-redshift AGN. In Sec.\\,6 we discuss the effect that obscuration plays in altering the intrinsic optical-to-X-ray flux ratios of quasars to those observed.  We assume throughout that $H_0=70~ {\\rm km~s^{-1}Mpc^{-1}}$, $\\Omega_{\\mathrm M}=0.3$ and $\\Omega_\\Lambda=0.7$. ",
        "conclusions": "Deep X-ray surveys are capable of uncovering the AGN responsible for most of the supermassive black hole accretion history of the universe. Combined X-ray and optical/near-infrared observations of lightly reddened quasars from such X-ray surveys are a powerful probe of the physical conditions of the obscuring material. We have used optical and near-infrared photometry and spectroscopy to constrain the amount of reddening in the light from optically-faint quasars. Our main conclusions are:  \\begin{itemize}  \\item{Optically-faint ($R>23$) quasars at $f_{0.5-8}>2 \\times 10^{-15}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$ are mostly faint due to obscuration. De-reddening the observed $R$-band fluxes of these quasars gives them optical magnitudes similar to other quasars with these X-ray fluxes ($20<R<22$).}  \\item{Assuming that most of the spread in X-ray hardness ratios is due to absorption by gas, we find that gas-to-dust ratios in reddened X-ray selected quasars are typically a few times greater than the gas-to-dust ratio in the Milky Way and similar to those in Seyfert galaxies.}  \\item{The fraction of quasars responsible for the X-ray background which are Type 1 in the optical (broad permitted lines) and Type 2 in the X-ray (hard X-ray spectrum due to absorption) is only about 6\\%.}  \\item{The observed distribution of optical-to-X-ray flux ratios of quasars at $z>1$ is skewed to low values compared to the intrinsic distribution due to the fact that the observed-frame $R$-band light is emitted in the UV and is more easily obscured than the hard X-rays sampled by {\\it Chandra}.}   \\end{itemize}"
    },
    "0403/astro-ph0403186_arXiv.txt": {
        "abstract": "{We use Kuiper's test to detect periodicities in X-ray and gamma-ray observations. Like Rayleigh's test, it uses the individual photon arrival times, and is therefore well suited to the analysis of faint sources. Our method makes it possible to take into account the discontinuities in the observation, and to completely get rid of the contamination that results from them. This makes it particularly adapted to the search of periods long compared to the total observation duration. We propose a semi-analytical approach to determine the effective number of trial frequencies when searching for unknown periods over a frequency range. This approach can be easily adapted to other tests. We show that, using Kuiper's test, we can recover periods in frequency domains where other tests are completely confused by contamination.  We finally search the entire ROSAT Position-Sensitive Proportional Counter (PSPC) archive for long periods, and find 28 new periodic-source candidates. ",
        "introduction": "\\label{sec:intro} Important efforts have been devoted to the search of periodic signals throughout the electromagnetic spectrum.  Because of the idiosyncrasies of astrophysical observations, different methods must be used depending on the type of object and the wavelength range. Four test families seem to dominate the period-detection ``market''. The calculation of the Fourier power spectrum density \\citep[e.g.,][]{PresEtal-1993-NumRec} using a fast Fourier transform (FFT) is adapted to evenly spaced (or evenly binned) observations. The Lomb-Scargle periodogram \\citep{Lomb-1976-LeaSqu,Scar-1982-StuAstII,HornBali-1986-PrePer}, a discrete Fourier transform method, can be used for unevenly-spaced flux measurements. Epoch folding (EF) \\citep[e.g., ][]{LeahEtal-1983-SeaPul} can be used in the same conditions or for individual photons, but requires a binning according to the phase. Rayleigh's test \\citep[e.g.,][]{GibsEtal-1982-TraEmi,Fish-1993-StaAna} is particularly adapted for the analysis of individual photons.  Observations in the X- and gamma-rays usually have two important characteristics. First, independent, time-tagged photons are collected. A method requiring binning is therefore far from ideal, as it results in a loss of information.  Furthermore, binning is prohibited for sources detected with very few photons; for EF for instance, the required assumption of Gaussian distribution in each bin is not satisfied in this case. Moreover, the necessary assumptions on the number and sizes of the bins lower the performance of the test \\citep{Schw-1999-OptPer}.  Secondly, space observations are often interrupted by ``bad time'' periods, where no data are received. Fourier-based methods and Rayleigh tests are seriously affected by this problem. In practice, it means that only periods short compared to the durations of uninterrupted observation can be investigated.  In this paper we present in detail Kuiper's test \\citep{Kuip-1960-TesCon}. This test has been applied to the distribution of solar flares \\citep{JetsEtal-1997-LonDis}, and to the search for periodicities in Earth impacts \\citep{Jets-1997-HumSta,JetsPelt-2000-SpuPer}, but its unique suitability to X-ray and gamma-ray observations has been overlooked. Similarly to Rayleigh's test, it uses discrete events, and can be applied to very faint sources without any {\\em a priori} assumption. Similarly to EF, it takes into account non-uniform coverage of the phase domain, and can therefore be used when searching for periods long compared to the total observation duration\\footnote{In this paper, ``total duration'' means the time interval between the start and the end of an observation, including possible gaps.}.  We study in detail the properties of Kuiper's test for period detection, and particularly its significance level. We concentrate on two important issues: the treatment of discontinuous observations, and the determination of the effective number of trial frequencies when searching for unknown periods. We finally apply the algorithm to the entire archive of the ROSAT Position-Sensitive Proportional Counter (PSPC) archive. ",
        "conclusions": "\\label{sec:conc} Kuiper's test shows very interesting properties for the search of long-period periodic objects. Its ability to cope very naturally, without any hidden assumption, with complex GTIs is unique.  Compared to Rayleigh's, Kuiper's test performs better for narrow-peaked light curves. Kuiper's test is quite sensitive to both subharmonics and harmonics of the fundamental frequency, but usually identifies the fundamental correctly.  Kuiper's test is particularly adapted to X-ray missions, like XMM-Newton and Chandra, high-energy gamma-ray satellites like GLAST, and for Cherenkov telescopes.  The semi-analytical method we propose here to correct the false-positive probability in case of a search over a range of frequencies should be quite useful in practice, not only for Kuiper's test, but also for other tests, as its principle can be easily adapted.  It has the advantage of simplicity, and of being based on sound probability principles.  On the 28 candidate periodic sources, 6 could be cross-checked using other ROSAT PSPC observations. Good or partial confirmation of the existence of periodicities is found in 3 of these objects, and there is total absence of confirmation in 3 objects. This does not necessarily imply a ``confirmation of absence''. It must be reminded that X-ray sources are quite often strongly variable, and that a periodic signal may remain undetected in some observations, even though the observing conditions seem adequate. For instance, \\citet{IsraEtal-2000-BepCha} report the detection of a periodic signal in the X-ray pulsar \\object{2E~0053.2-7242} in only one out of nine ROSAT PSPC observations, the source having dimmed by a factor $>\\!\\!6$ between the different observations.  The possibility that extrinsic contamination, or statistical flukes explain some, or even most, of the candidate periods must be considered seriously. Firm identification of the candidates as periodic sources will be contingent upon the detection of the periods in independent data sets. The building up of important X-ray archives from XMM-Newton and Chandra makes it quite probable that new observations will be available for a fair number of these sources in the near future.  A C library implementing the algorithms discussed in this paper is available from the author."
    },
    "0403/astro-ph0403275.txt": {
        "abstract": "We combine the recently published CIZA galaxy cluster catalogue with the XBACs cluster sample to produce the first all-sky catalogue of X-ray clusters in order to examine the origins of the Local Group's peculiar velocity without the use of reconstruction methods to fill the traditional Zone of Avoidance.  The advantages of this approach are (i) X-ray emitting clusters tend to trace the deepest potential wells and therefore have the greatest effect on the dynamics of the Local Group and (ii) our all-sky sample provides data for nearly a quarter of the sky that is largely incomplete in optical cluster catalogues.  We find that the direction of the Local Group's peculiar velocity is well aligned with the CMB as early as the Great Attractor region 40 $h^{-1}$ Mpc away, but that the amplitude of its dipole motion is largely set between 140 and 160 $h^{-1}$ Mpc.  Unlike previous studies using galaxy samples, we find that without Virgo included, roughly $\\sim70\\%$ of our dipole signal comes from mass concentrations at large distances ($>60$ $h^{-1}$ Mpc) and does not flatten, indicating isotropy in the cluster distribution, until at least 160 $h^{-1}$ Mpc.  We also present a detailed discussion of our dipole profile, linking observed features to the structures and superclusters that produce them.  We find that most of the dipole signal can be attributed to the Shapley supercluster centered at about 150 $h^{-1}$ Mpc and a handful of very massive individual clusters, some of which are newly discovered and lie well in the Zone of Avoidance. ",
        "introduction": "Since the detection of a dipole anisotropy in the Cosmic Microwave Background (CMB, Kogut 1993), many question have been raised regarding the origin of the Local Group (LG) motion which is thought to give rise to such a signal.  Specifically at question have been the nature of the cosmic objects or structures most directly responsible for inducing the LG's peculiar velocity and, furthermore, the distance out to which inhomogeneities in the distribution of these objects continue to affect the LG's dynamics.  Interest in these questions is motivated by the cosmological implications they carry; for example, if the motion of the Milky Way (MW) is induced entirely nearby, then the current density of the universe would need to be considerable in order for nearby matter concentrations to adequately accelerate the MW over such relatively small scales.  On the other hand, if the LG's peculiar velocity is induced by more distant structures, then we know that anisotropies in the large-scale matter distribution must exist to at least those structures and that the universe becomes isotropic only at larger distances.  To answer these questions, much effort has been spent examining the peculiar velocity that would be induced onto the LG by the distribution of various mass tracers, such as galaxies and clusters of galaxies, and comparing this motion with that inferred from the CMB.  Traditionally such analyses have made use of the linear theory of gravitational instability, which dictates that the peculiar velocity of a reference frame can be related to the gravitational acceleration induced by the mass distribution surrounding it via (Peebles 1976, see our Appendix A)  \\begin{equation} \\textit{\\textbf{v}}_{p} \\hspace{.1in} = \\hspace{.1in} \\frac{H_{o}\\beta}{4\\pi \\bar{n}} \\int \\frac{n(r)}{r^{2}} \\textit{\\textbf{\\^{r}}}\\ dr \\end{equation}  \\noindent where $\\beta = \\Omega^{0.6}_{0}/b$ and $b$ is the biasing factor relating the mass-tracers to the underlying mass distribution they represent, and $\\bar{n}$ is the average mass-tracer number density.  In other words, Equation 1 tells us that the dipole moment of a mass-tracer distribution can be directly related to the peculiar velocity that sample would induce on the LG.  Within this framework, the comparison of the LG's peculiar velocity as inferred from the CMB dipole anisotropy to that produced by a mass-tracer distribution can shed light on the role of the sample on producing the LG's motion, provide an estimate of the depth at which inhomogeneities in the distribution of the sample affect the LG's dynamics (i.e. the convergence depth, $R_{conv}$), and place constraints on how the sample traces the underlying matter distribution in the form of the biasing parameter $\\beta$.  Performing this type of analysis is a nontrivial matter for numerous reasons: (i) for linear perturbation theory to be applicable, the dipole moment of the sample under study must be relatively well aligned with the observed peculiar velocity of the LG; (ii) the characteristic depth of the sample must be larger than the depth at which anisotropies in the sample can influence the LG's dynamics ($R_{*} > R_{conv}$); (iii) the number of mass-tracers must be sufficient to accurately sample the underlying density field and avoid the introduction of shot-noise errors.  Wary of these concerns, this type of dipole analysis has been extensively applied to the LG.  While studies have used samples ranging from optical galaxies (Lahav, Rowan-Robinson \\& Lynden-Bell 1988;  Lynden-Bell, Lahav \\& Burstein 1989; Shaya et al. 1992; Hudson 1993) to X-ray selected AGN (Miyaji \\& Boldt 1990), much of the early dipole work focused on IRAS galaxies due to their considerable sky coverage (Meiksin \\& Davis 1986; Yahil, Walker \\& Rowan-Robinson 1986; Strauss \\& Davis 1988; Yahil 1988; Rowan-Robinson et al. 1990, 1991; Strauss et al. 1992; Plionis, Coles \\& Catelan 1993; Basilakos \\& Plionis 1998).  Despite the wide employment of such galaxy samples, they are all plagued by steeply declining selection functions with distance, leading to a significant incompleteness at large distances and characteristic depths of 80-100 $h^{-1}$ Mpc.  Analyses of galaxy samples have produced dipoles that are generally well aligned with the CMB dipole direction, but many differ regarding their implied convergence depth.  Strauss et al. (1992) and Hudson (1993) found an $R_{conv}$ of roughly $\\sim$$50 h^{-1}$ Mpc, implying that all of the peculiar acceleration of the LG is induced relatively nearby.  On the other hand, Plionis (1988), Plionis, Coles \\& Catelan (1993), Basilakos \\& Plionis (1998), Branchini et al. (1999), Schmoldt et al. (1999), and Plionis et al. (2000) found evidence for contributions ranging from 100 to 150 $h^{-1}$ Mpc. The credibility of many of these early results must be questioned since the convergence depth of each sample shows a strong dependence on the sample's characteristic depth, a feature to be expected if anisotropies beyond that characteristic depth continued to contribute to the motion of the LG.  This implies that galaxy samples alone do not probe the mass fluctuation field deep enough to account for all the anisotropies that affect the LG's dynamics.  The use of clusters of galaxies overcomes some of the problems faced by galaxy samples since clusters are luminous enough for samples to be volume-limited out to larger distances.  Although clusters only sparsely sample the underlying density field, which introduces shot-noise errors, they trace the peaks of the density fluctuation field, which have the greatest effect on the dipole amplitude and direction.  Using the Abell/ACO cluster catalogue (Abell 1958, Abell, Corwin \\& Olowin 1989, hereafter ACO), which has a characteristic depth of over 240 $h^{-1}$ Mpc, evidence has been found for contributions to the LG dipole from depths of $\\sim$160 $h^{-1}$ Mpc (Scaramella, Vettolani \\& Zamorani 1991,1994; Plionis \\& Valdarnini 1991; Brunozzi et al. 1995; Branchini \\& Plionis 1996).  These results were confirmed by Plionis \\& Kolokotronis (1998) using the X-ray Brightest Abell-type Cluster catalogue (XBACs, Ebeling et al. 1996), which is optically selected, but X-ray confirmed, thus eliminating projection effects which may have enhanced the dipole amplitude obtained with Abell/ACO (Sutherland 1988).  In the context of the dipole analysis, the primary limitation of the XBACs and Abell/ACO catalogues are their incompleteness at low Galactic latitudes.  This is because traditional optical cluster searches suffer from severe extinction and stellar obscuration in the direction of the Milky Way (MW), leading to catalogues with poor coverage in a $40^{\\circ}$ wide strip centered on the plane of the Galaxy, known as the Zone of Avoidance (ZOA).  This is particularly troubling since Shaya suggested as early as 1984 that large-scale structures obscured by the ZOA could have a significant effect on the peculiar motion of the LG.  More recently, several studies have found renewed evidence for a significant bulk motion toward a vertex behind the plane of the MW (Riess et al 1997; Hudson et al. 1999), rekindling the idea of a Great Attractor (Lynden-Bell et al. 1988).  A variety of techniques have been used to reconstruct the ZOA, ranging from a 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).  The value of these reconstruction techniques is, however, limited if the MW does indeed obscure dynamically significant regions, as has been suggested.  With the recent compilation of the X-ray selected CIZA cluster catalogue (named for Clusters in the Zone of Avoidance; Ebeling, Mullis \\& Tully 2002), it has become possible to add real cluster data to the region behind the MW.  In this study we combine the CIZA and XBACs samples, with appropriate weightings, to produce the first all-sky catalogue of X-ray luminous clusters and provide a dipole analysis of the LG peculiar velocity without the use of reconstruction methods to fill the ZOA.  We proceed in the following manner: in section 2 we give an overview of the XBACs and CIZA samples, with a thorough look at the systematic effects in each. Section 3 describes the details of the dipole analysis and our results are put forward in section 4.  Finally we summarize our primary conclusions in section 5.  Throughout this paper we assume an Einstein-de Sitter universe with $q_{0}=0.5$ and $H_{0}=100$ $h$ km s$^{-1}$Mpc$^{-1}$, so that our results are directly comparable to those of previous studies.   \\begin{figure}[t] \\hspace*{-0.15in} \\centerline{\\psfig{file=f1.eps,width=3.25in,angle=0}} \\caption{Aitoff projection of the XBACs and CIZA cluster catalogues in Galactic Coordinates.  The dashed lines represent the traditional Zone of Avoidance ($|b|<20^{\\circ}$).} \\end{figure} ",
        "conclusions": "One of the limitations to the use of clusters of galaxies to identify the structures inducing the LG's peculiar motion has been the incompleteness of cluster catalogues at low Galactic latitudes.  In this study we have combined the recently published CIZA catalogue with the XBACs cluster sample to produce the first all-sky catalogue of X-ray bright clusters in order to analyze the origin of the LG peculiar velocity without the need for reconstruction methods to fill the traditional ZOA.  We find that the cluster distribution becomes isotropic with respect to the LG at a distance of 160 $h^{-1}$ Mpc and that with Virgo excluded the asymptotic value of the dipole amplitude is largely set by a coherent signal near $\\sim150 h^{-1}$ Mpc.  While this agrees with previous results, our finding that $\\sim70\\%$ of the dipole amplitude is set at distances larger than $60 h^{-1}$ Mpc differs from results obtained using galaxy and optically selected cluster samples.  By examining the dipole profile on a cluster-by-cluster basis, we conclude that the Shapley concentration is the single supercluster most responsible for producing the dipole signal between 140 and 160 $h^{-1}$ Mpc.  We also find that the cluster dipole is fairly well aligned with the direction of the CMB dipole at relatively shallow distances, and that the competing effects of Virgo cluster and the PC concentration essentially counteract each other's effects on the final dipole pointing.  These facts, coupled with the significant contributions to the LG peculiar acceleration from larger distances, reaffirm the bootstrap theory which suggests that the aligned mass concentrations of the GA and Shapley regions largely set the direction of the LG's peculiar motion throughout our study volume.  Furthermore our analysis has identified four dynamically interesting CIZA clusters: C1324 in the GA region, C1410 in the Shapley supercluster and the two clusters C1652 and C1638 (Triangulum Australis) which lie behind the GA region at roughly the same distance as Shapley but well in the ZOA.  Lastly, using the LG's peculiar velocity as measured from the CMB anisotropy (corrected for Virgo-centric infall) and the amplitude of the cluster dipole without Virgo included, we find the average value of the $\\beta$ parameter to be $\\beta = 0.24 \\pm 0.02$, in agreement with previous values determined from the XBACs catalogue alone.  %On a final cautionary note, we remind the reader that including Virgo in our analysis highlighted the fact that our sample does note trace the mass distribution in the local neighborhood very well. %On a final cautionary note, we remind the reader that since our sample does not include the most nearby clusters such as Virgo, it does not constrain the dipole zero-point very well. %On a final cautionary note, we remind the reader that our sample does not include some of the nearest clusters such as Virgo, which causes our dipole zero-point to be ill-constrained.  On a final cautionary note, we remind the reader that the sparseness of X-ray bright clusters causes them to be noisy tracers of the mass distribution when very small volumes are considered, which inturn causes the dipole zero-point to be ill-constrained.  Therefore one should bear in mind that our conclusions regarding the relative contributions of different distance scales to the final dipole amplitude are valid only in the range we sample well: $40$ to $240 h^{-1}$ Mpc.  That being said, we can think of two methods by which our results can be extended to shallower distances: (i) since truly X-ray selected cluster catalogues such as the BCS have been shown to better sample the nearby cluster distribution, the use of an all-sky, X-ray selected cluster sample would allow a dipole analysis to be performed without the need for weights to compensate for the residual incompleteness in the XBACs sample.  The construction of such a sample will be possible in the near future when the REFLEX catalogue becomes available; its combination with the BCS and CIZA samples will produce the first entirely X-ray selected, all-sky data set.  (ii) Nearby galaxy catalogues can be used to establish a zero-point for the X-ray cluster dipole, from which the relative contributions of different distance scales can truly be determined.  This galaxy dipole normalization has been applied to optical cluster catalogues (Scaramella et al. 1994), but has yet to be implemented on a X-ray selected cluster sample."
    },
    "0403/astro-ph0403379_arXiv.txt": {
        "abstract": "By integrating the relativistic hydrodynamic equations introduced by Taub we have determined the exact EQuiTemporal Surfaces (EQTSs) for the Gamma-Ray Burst (GRB) afterglows. These surfaces are compared and contrasted to the ones obtained, using approximate methods, by \\citet{pm98c,s98,gps99}. ",
        "introduction": "The recent explanation of the observed luminosity in X- and $\\gamma$-ray energy bands in Gamma-Ray Bursts (GRBs), as well as the comprehension of their spectral properties, depends in an essential way on the determination of the EQuiTemporal Surfaces (EQTSs) in the afterglow era \\citep{rbcfx02_letter,Spectr1}. Here we compare and contrasts the exact determination of the EQTSs with the approximate expressions presented by \\citet{pm98c,s98,gps99}.  A great deal of consensus has been reached concerning three basic issues in the description of GRB afterglows: \\\\ {\\bf a)} Their origin is generally traced back to the interaction of an ultrarelativistic baryonic matter pulse with the InterStellar Medium (ISM). It is also agreed that the relativistic hydrodynamic equations introduced by Abe \\citet{Taub} are the correct theoretical framework to describe such an interaction.\\\\ {\\bf b)} The general definition of the EQTSs is also agreed upon by everyone.\\\\ {\\bf c)} The necessity of determining the boundary conditions of the baryonic matter pulse in the early phases of the afterglow era is also generally recognized.  We illustrate in the following sections the equations needed for the description of these three basic issues and also point out a major difference between our approach and the ones in the current literature concerning the solutions of the Taub equations. We identify in this difference the reason for the results on the EQTSs presented in Fig.~\\ref{eqts_comp_ad} and Fig.~\\ref{eqts_comp_rad}. Our treatment assumes the exact integration of the Taub equations.  We recall that all the GRB observable quantities depend essentially on the EQTSs. Due care is therefore needed in their correct determination. ",
        "conclusions": "The consequences of using the approximate formula given in Eq.(\\ref{gamma_app}) to compute the expression $t \\equiv t(r)$, instead of the exact solution of the Taub Eqs.(\\ref{Taub_Eq}), are clearly shown in Figs.~\\ref{eqts_comp_ad}--\\ref{eqts_comp_rad}. The EQTSs represented in these figures are computed at selected values of the detector arrival time both in the early ($\\sim 35$ s) and in the late ($\\sim 4$ day) phases of the afterglow. Both the fully radiative and fully adiabatic cases are examined. Note the approximate expression of the EQTS can only be defined for $\\gamma < \\gamma_d$ and $r > r_d$. Consequently, at $t_a^d=35$ s the approximate EQTSs are represented by arcs, markedly different from the exact solution (see the upper panels of Figs.~\\ref{eqts_comp_ad}--\\ref{eqts_comp_rad}). The same conclusion is found for the EQTS at $t_a^d=4$ days, where marked differences are found both for the fully radiative and adiabatic regimes (see the lower panels of Figs.~\\ref{eqts_comp_ad}--\\ref{eqts_comp_rad}).  All the observational properties of GRBs, starting from the analysis of the prompt radiation \\citep{rbcfx02_letter}, to the luminosity in X- and $\\gamma$-ray bands, to their spectral distribution \\citep{Spectr1} as well as inferences on the possible presence or absence of beaming in GRBs, depend essentially on the structure of the EQTSs. In turn the determination of the EQTSs depends on the equations of motion of the baryonic pulse satisfying the Taub equations. The fact that the final results for the observable luminosity, spectral distribution, and substructures in the prompt radiation depend on $\\sim 10^8$ integration paths on different points on the EQTSs implies that the agreement between the theoretical predictions and the observations becomes a most stringent test for the validity of the equations of motion. The correct EQTSs are also essential for the identification of the energy source of the X and $\\gamma$ radiation in the GRB afterglow \\citep{rbcfx02_letter,Spectr1}.  In conclusion, the approximate treatments largely overestimate (underestimate) the size of the EQTSs in the early (late) part of the afterglow. The theoretical slopes of the observables as a function of the arrival time \\citep[see, e.g.,][and references therein]{p99,p00,vpkw00} are therefore incorrectly evaluated. In the meantime, analytic expressions for the EQTSs have been obtained, validating the above results and allowing the theoretical estimate of the observables in GRB afterglows \\citep{br04}."
    },
    "0403/astro-ph0403653_arXiv.txt": {
        "abstract": "{% In this paper we present a multifrequency and multiepoch study of \\object{PKS\\, 1502+106} at radio frequencies. The analysis is based on an EVN (European VLBI Network) dataset at 5 GHz and archive VLBA (Very Long Baseline Array) datasets at 2.3, 8.3, 24.4 and 43.1 GHz over a period of 8 years. The source is characterized by a multi--component one--sided jet at all epochs. The high--resolution images at 5, 8.3, 24.4 and 43.1 GHz show a curved jet morphology in the source. The radio core brightness temperature approaches the equipartition limit. Superluminal motions of $37.3\\pm9.3\\;c$, $22.0\\pm15.5\\;c$, $10.5\\pm2.6\\;c$ and $27.9\\pm7.0\\;c$  are measured in four distinct components of the jet. Our analysis supports the idea that the relativistic jet in \\object{PKS\\,1502+106} is characterised by extreme beaming and that its radio properties are similar to those of $\\gamma$--ray loud sources. ",
        "introduction": "\\begin{table*} \\centering \\caption{Logs of the observations} \\begin{tabular}{ccclc} \\hline  Epoch &Freq  &BW &Array and Available Telescopes$^{a}$&$D_{uv}^{b}$  \\\\ &(GHz) &(MHz)    &                           &(km)          \\\\ (1) & (2)  & (3)     & (4)                       &(5)            \\\\\\hline  1994.52 &2.3 & 16      &VLBA(All 10)     & 8600  \\\\  &8.3 & 16      &VLBA(All 10)     & 8600  \\\\  1997.85 &5.0 & 28      & EVN(Ef Sh Jb Ht Mc Nt On Tr Ur Wb) & $\\sim$10000 \\\\  1998.11 &2.3 & 32      &VLBA(BR FD MK OV PT)  GC KK      &             \\\\  &    &         &VLBA(BR FD HN KP LA NL OV PT SC) GC WF GN &5600  \\\\  &8.3 & 32      &VLBA(BR FD HN KP LA NL OV PT SC) GC WF GN &5600  \\\\  2002.37 &24.4& 32      &VLBA(BR FD HN KP LA MK NL OV SC)     & 8600   \\\\  &43.1& 32      &VLBA(BR FD HN LA MK NL OV SC) & 8600  \\\\  2002.65 &24.4& 30      &VLBA(All 10)     & 8600  \\\\  &43.1& 32      &VLBA(BR FD HN KP LA MK NL OV PT) & 6600 \\\\  2002.99 &24.4& 64      &VLBA(FD HN KP LA MK NL OV PT SC) & 8600 \\\\  &43.1& 64      &VLBA(FD HN KP LA MK NL OV PT SC) & 8600 \\\\  \\hline\\end{tabular} \\label{obs} \\\\[0.3cm] \\raggedright $^{a}$ EVN telescope codes: Ef: Effelsberg, Sh: Shanghai, Ht: HartRAO, Jb: Jodrell, Mc: Medicina, Nt: Noto, On: Onsala, Tr: Torun, Ur: Urumqi, Wb: Westerbork Array; the VLBA observations at epoch 1998.11 are performed using subarrays made up of the VLBA antennas together with 4 geodetic antennas: GC (Fairbanks, AK USA), WF (Westford, MA USA), GN (Green Bank, WV USA) and KK (Kokee Park, HI USA);\\\\ $^{b}$ the longest baseline of the array, in unit of kilometer. \\end{table*}   \\begin{table*} \\centering \\caption{Parameters of the images} \\begin{tabular}{lccllccc} \\hline  Figure& Epoch & Freq  & Real Beam$^{a}$ & Restored Beam$^{b}$ &S$_{Peak}~^{c}$ & rms$^{d}$ &Contours\\\\ &       & (GHz) &Maj$\\times$Min(mas),P.A.($\\degr$)& Maj$\\times$Min(mas),P.A.($\\degr$)&(Jy/b) & (mJy/b)&(mJy/b)\\\\ (1) & (2) & (3) & (4) &(5) & (6) & (7) &(8) \\\\\\hline Fig.\\ref{fig1}a&1994.52 &2.3 & 7.18$\\times$3.78, $-1.67$&7.18$\\times$3.78, 0& 1.81&1.1 & 3.5$\\times$(-1,1,2,...,256)\\\\  Fig.\\ref{fig1}b&1998.11 &2.3 & 5.94$\\times$4.54, 13.7 &7.18$\\times$3.78, 0& 1.23&0.9 & 2.8$\\times$(-1,1,2,...,256)\\\\  Fig.\\ref{fig2}a&1997.85 &5.0 & 1.36$\\times$1.21,  64.7 &1.25$\\times$1.25, 0& 0.80&0.7 & 3.0$\\times$(-1,1,2,...,256)\\\\  Fig.\\ref{fig2}b&1994.52 &8.3 & 1.95$\\times$1.03, $-1.19$&1.25$\\times$1.25, 0& 1.37&0.9 & 2.8$\\times$(-1,1,2,...,256)\\\\  Fig.\\ref{fig2}c&1998.11 &8.3 & 1.62$\\times$1.23, 16.7 &1.25$\\times$1.25, 0& 0.68&0.6 & 2.0$\\times$(-1,1,2,...,256)\\\\  Fig.\\ref{fig3}a&2002.37 &24.4& 0.64$\\times$0.28, $-$1.08&0.64$\\times$0.28, 0& 1.26&0.8 & 2.1$\\times$(-1,1,2,...,256)\\\\  Fig.\\ref{fig3}b&2002.65 &24.4& 0.70$\\times$0.32, $-$3.77&0.64$\\times$0.28, 0& 0.73&0.9 & 2.8$\\times$(-1,1,2,...,256)\\\\  Fig.\\ref{fig3}c&2002.99 &24.4& 0.76$\\times$0.28, $-$6.0 &0.64$\\times$0.28, 0& 0.95&1.0 & 3.0$\\times$(-1,1,2,...,128)\\\\  Fig.\\ref{fig4}a&2002.37 &43.1& 0.37$\\times$0.16, $-$3.35&0.37$\\times$0.16, 0& 0.82&1.0 & 3.0$\\times$(-1,1,2,...,256)\\\\  Fig.\\ref{fig4}b&2002.65 &43.1& 0.66$\\times$0.24, $-$20.3&0.37$\\times$0.16, 0& 0.50&1.2 & 3.7$\\times$(-1,1,2,...,128)\\\\  Fig.\\ref{fig4}c&2002.99 &43.1& 0.44$\\times$0.16, $-$6.41&0.37$\\times$0.16, 0& 0.59&1.0 & 3.2$\\times$(-1,1,2,...,128)\\\\ \\hline \\end{tabular}\\label{image} \\\\[0.3cm] \\raggedright $^a$ the beam size directly measured from the visibilities (FWHM);\\\\ $^b$ the restored beam shown in the images (FWHM);\\\\ $^c$ the peak brightness in the image; \\\\ $^d$ the off--source rms noise level in the images. \\end{table*}   One of the most significant observational results of extragalactic $\\gamma$--ray active galactic nuclei (AGNs) is that all EGRET--identified objects are radio--loud sources (Mattox et al. \\cite{Mattox}). Relativistic beaming in the jet is used to explain the EGRET identification in radio--loud AGNs. The EGRET--identified sources have on average much faster apparent superluminal motions than the general population of radio--loud sources (Jorstad et al. \\cite{Jorstad}). From a statistical analysis of $\\Delta PA$ (position angle differences between parsec-- and kiloparsec--scale structures) of EGRET--identified AGNs, Hong et al. (\\cite{Hong}) concluded that $\\gamma$--ray loud radio quasars typically show aligned morphologies on parsec and kiloparsec scales. It is still a matter of debate if the $\\gamma$--ray emission in AGNs is related to higher beaming in these sources.  The radio--loud active galactic nucleus (AGN) \\object{PKS\\,1502+106} (4C 10.39, OR103), $z=1.833$ (Fomalont et al. \\cite{Fomalont}), is a highly polarized quasar (Tabara \\& Inoue \\cite{Tabara}). A high and variable degree of polarisation in the optical band is reported by Impey \\& Tapia (1988). It is known to be active and variable at radio, optical and X--ray wavelengths. In particular, the source exhibits intensity variations by a factor of 3 $\\sim$ 5 on timescales from weeks to months in the radio band; it shows intensity variations in the optical band with $m_{\\nu}$ ranging from 19.5 to 18.6, and in the X--ray band by a factor $\\geq 2$ at 1~keV (George et al. \\cite{George} and references therein).  \\object{PKS\\, 1502+106} exhibits a 'core--jet--lobe' structure at radio wavelengths. A Very Large Array (VLA) image at 1.64 GHz (Murphy, Browne \\& Perley \\cite{Murphy}) shows that a continuous jet extends to the southeast and leads to a lobe located $\\sim$~7 arcsecond from the core. Very Long Baseline Interferometry (VLBI) observations (Fey, Clegg \\& Fomalont \\cite{Fey}; Fomalont et al. \\cite{Fomalont}; Zensus et al. \\cite{Zensus}) show a well--defined jet starting to the southeast and sharply bending to the east at a distance of 3--4 mas from the core.  Our interest in \\object{PKS\\,1502+106} is related to the misalignment of the pc-- and kpc--scale radio structure in AGNs and its relation to the $\\gamma$--ray emission. Fichtel et al. (\\cite{Fichtel}) reported an EGRET flux density upper limit in PKS 1502+106 of $7 \\times 10^{-8}$ photons cm$^{-2}$ s$^{-1}$ from the Phase I results. However, the source was not detected in the following EGRET observations (Thompson et al. \\cite{Thompson}; Hartman et al. \\cite{Hartman}). To study the relation between the $\\Delta PA$ and $\\gamma$--ray emission, we observed a sample of $\\gamma$--ray blazars with the EVN at 5 GHz in Nov. 1997, in which PKS 1502+106 was observed as a gamma--ray source candidate (Hong et al. \\cite{Hong04}).  In the next section we present the EVN observation of \\object{PKS\\,1502+106} at 5 GHz, and the data reduction. The analysis of the 5 GHz EVN image is performed in Sect. 3, where we also re--analyzed five epochs of Very Long Baseline Array (VLBA) datasets at 2.3, 8.3, 24.4 and 43.1 GHz obtained from the public archive of Radio Reference Frame Database (RRFID) for a consistent study with our data. To investigate the relativistic beaming in this source, we estimate the core physical parameters and measure the jet proper motions in Sect. 4. Conclusions and summary are given in Sect. 5. $H_0$= 65 km\\, s$^{-1}$Mpc$^{-1}$ and $q_0$=0.5 are used in this paper.   \\begin{figure}\\centering \\resizebox{6.5cm}{!}{\\includegraphics{fig-1.eps}} \\vspace{-1mm} \\caption{VLBA images of \\object{PKS\\, 1502+106} at 2.3 GHz. A: epoch 1994.52; B: epoch 1998.11.} \\label{fig1} \\end{figure} \\begin{figure}\\centering \\resizebox{6.5cm}{!}{\\includegraphics{fig-2.eps}}\\vspace{-1mm} \\caption{VLBI images of \\object{PKS\\, 1502+106}. A: EVN image at 5 GHz at epoch 1997.85; B: VLBA image of at 8.3 GHz at epoch 1994.52; C: VLBA image at 8.3 GHz at epoch 1998.11. } \\label{fig2} \\end{figure} \\begin{figure}\\centering \\vspace{3mm} \\resizebox{6.5cm}{!}{\\includegraphics{fig-3.eps}} \\vspace{-1mm} \\caption{VLBA images of \\object{PKS\\, 1502+106} at 24.4 GHz. A: epoch 2002.37; B: epoch 2002.65; C: epoch 2002.99.} \\label{fig3} \\end{figure}  \\begin{figure}\\centering \\vspace{3mm} \\resizebox{6.5cm}{!}{\\includegraphics{fig-4.eps}} \\vspace{-1mm} \\caption{VLBA images of \\object{PKS\\, 1502+106} at 43.1 GHz. A: epoch 2002.37; B: epoch 2002.65; C: epoch 2002.99.} \\label{fig4} \\end{figure} ",
        "conclusions": "In this paper we carried out a multifrequency and multiepoch study of \\object{PKS\\, 1502+106} at VLBI resolution. The source morphology is highly variable and the jet structure is very complex on this scale.  We analyzed the structure of the jet, as seen in projection on the sky, and the changes in the position angle of the various jet components with the radial distance from the core. The results suggest that the jet components in \\object{PKS\\, 1502+106} trace the same curved path. The jet undergoes two major bends, the first implies a change in the P.A from $\\sim70\\degr$ to $\\sim110\\degr$ within 0.5 mas, while in the second bend the P.A. changes from $\\sim130\\degr$ to P.A.$\\sim80\\degr$ at 3--4 mas. Beyond that distance the VLBI jet points to the east at 7--15 mas.  Based on the model fitting results on the VLBI core, we obtain a weighted mean value of the Doppler factor, $\\delta=4.5\\pm1.9$. The radio core brightness temperature in the source rest frame, T$_r = (2.0\\pm0.5) \\times10^{11}$K, approaches the equipartition limit.   We detect superluminal motion in four components along the jet. The derived apparent speeds range between $10.5\\pm2.6\\,c$ and $37.3\\pm9.3\\,c$. Doppler boosting plays a major role in determining the observed properties of the source. In particular, the apparent speeds we derive suggest that the source is viewed under an angle to the line of sight $\\theta < 5^{\\circ}$, and that the bulk flow velocity is $\\beta_{intr}$ is $\\sim$ $0.999\\,c$.  The $\\Delta PA$ between the pc-- and kpc--scale structure is about 30\\degr{}, indicating that PKS\\,1502+106 belongs to the aligned population. The superluminal speeds in 1502+106 are much higher than the average value in radio loud quasars. We therefore conclude that PKS\\,1502+106 is more beamed than the overall population of radio loud quasars, and that its radio properties are more similar to the $\\gamma$--ray loud quasars, although it is unclear if \\object{PKS\\, 1502+106} is a $\\gamma$--ray loud source to date. In particular, the superluminal speeds in this source are in the range found for the high redshift $\\gamma$--ray loud quasars in other high frequency radio surveys (Jorstad et al. 2001).  A confirmation of $\\gamma$--ray emission from this source would be highly valuable for our understanding of the $\\gamma$--ray loudness phenomenon in radio loud quasars."
    },
    "0403/hep-ph0403199_arXiv.txt": {
        "abstract": "Flat directions in the minimal supersymmetric standard model are known to deform into non-topological solitons, $Q$-balls, which generally possess both baryon and lepton asymmetries. We investigate how $Q$-balls evolve if some of the constituent fields of the flat direction decay into light species. It is found that the $Q$-balls takes a new configuration whose energy per charge slightly increases due to the decay. Specifically, we show that all the stable $Q$-balls eventually transform into pure $B$-balls via the decay into neutrinos. ",
        "introduction": "\\label{sec:intro}  Non-topological solitons, such as $Q$-balls~\\cite{Coleman:1985ki} and $I$-balls~\\cite{Kasuya:2002zs}, play important roles in the particle cosmology, partly because they are (quasi-)stable objects. Their stability comes from the conservation law. The $Q$-ball is composed of a complex scalar field with an $U(1)$ charge, the conservation of which enables the $Q$-ball to be stable or long-lived. Similarly, $I$-balls are composed of a real scalar field (or a complex scalar field in a nearly straight-line motion), whose dynamics conserve the adiabatic charge.  Since we can only see either these objects in their final state or their decay products in the present universe, it is important to study their evolution during the course of the expanding universe.  The baryon asymmetry is one of the key observational results to uncover the history of the universe. Among many baryogenesis scenarios proposed so far, the mechanism proposed by Affleck and Dine~\\cite{Affleck:1984fy}  can be realized in the minimal extension of the standard model, that is, the minimal supersymmetric standard model (MSSM).  The mechanism makes use of one of flat directions, along which there is no classical potential in the supersymmetric limit.  Although the flat direction is parametrized by a set of chiral superfields, it is convenient and even sufficient in the usual situation to express it in terms of a single complex scalar field dubbed `Affleck-Dine (AD) field,' $\\Phi$, which generally has nonzero baryon and lepton charges. The dynamics of the AD field not only generates the desired baryon asymmetry, but allows non-topological solitons to be formed. The properties of $Q$-balls in the context of the AD mechanism have been extensively studied~\\cite{Kusenko:1997si,Enqvist:1998en,Kasuya:1999wu, Kasuya:2000wx}.  $Q$-balls can be stable or unstable depending on the situation. Stable $Q$-balls can be dark matter~\\cite{Kusenko:1997si,Kasuya:2000sc, Kasuya:2001hg}, which might solve the coincidence problem between the abundance of baryons and that of dark matter.  Unstable but long-lived $Q$-balls can keep the baryon and lepton asymmetries stored inside them from being subject to the sphaleron effects.  For instance, they can successfully generate both large lepton asymmetry and small baryon asymmetry~\\cite{Kawasaki:2002hq}. It is even possible to realize the late-time baryon appearance after the relevant BBN epoch with use of $Q$-balls~\\cite{Ichikawa:2004pb}. $I$-balls are also formed in the AD mechanism as an excited state of $Q$-balls.  In the examples mentioned above, the single-field parametrization is simple and valid. However, if we examine the evolution of $Q$-balls closely, we have to pay much more attention to the fact that the flat direction is actually composed of several scalar fields.  Indeed, the slepton condensate always decays into a pair of neutrinos, while the squark condensate can decay into hadrons only if its energy per one baryon charge exceeds the nucleon or meson masses. Therefore, if a $Q$-ball has nonzero baryon and lepton asymmetries and its energy per unit charge is small enough, the leptonic component decays, leaving only the baryon asymmetry inside $Q$-balls ({\\it i.e.}, pure $B$-balls). However, since the $Q$-ball solution is obtained under the assumption that the $U(1)$ charge, $Q$, is conserved, it is not trivial how the $Q$-ball configuration changes during the course of the partial decay.  Do $Q$-balls come apart? Or do they change their shape to satisfy the modified $Q$-ball solution?  It is the purpose of this study to answer these questions.  In this paper, we investigate how $Q$-balls evolve if some of the constituent fields decay into lighter species. In the next section, we review the flat direction of the MSSM and the $Q$-ball solution. In section~\\ref{sec:model} we study the partial decay of $Q$-balls with a toy model of the flat direction.  The results obtained there are confirmed with use of numerical calculation presented in section~\\ref{sec:nume}.  Finally we give discussions and conclusions in section~\\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss} So far we have presented the analytical and numerical study of the partially decaying $Q$-balls with use of a toy model of flat directions. Let us consider now the implication of the results presented above. We have assumed that only a part of constituent fields decay, but our arguments apply to more generic case where there is a hierarchy between the decay rates, say, $\\Gamma_1 \\gg \\Gamma_2$. The asymmetry of the $Q$-ball continuously changes during the course of the decay. Most of the flat directions in the MSSM are composed of both squarks and sleptons, therefore they have both baryon and lepton number. If such a flat direction deforms into $Q$-balls and they are stable, our argument can be applied since the slepton condensate can always decay into a pair of neutrinos through the exchange of the gauginos. All through the decay processes, the $Q$-balls sequentially transform themselves toward pure $B$-balls with no lepton number. In the usual scenario considered so far, it is a flat direction composed only of squarks that forms pure $B$-balls. Therefore our analysis indicates that the pure $B$-balls are generated for most of flat directions, as a result of the decay into neutrinos. Specifically, all the $Q$-ball dark matter must be pure $B$-balls. Interestingly enough, the resultant $B$-ball might become unstable due to the increase of its energy per unit charge, leading to a complete decay of the $Q$-ball. In connection with the transformation of $Q$-balls into $B$-balls, the  leptogenesis  might become possible even for a flat direction with $B-L =0$. Another interesting possibility is that the $I$-ball  can be naturally altered to Q-ball, if one of the constituent fields decay into something else. In particular, the $I$-ball composed of  a flat direction that includes a leptonic field  is generally transformed into a $Q$-ball. In fact, since the $I$-ball is a quasi-stable object, its conversion  to  a $Q$-ball was not  known. The results of this study affords some new perspectives on the decay or stabilization processes of non-topological solitons made of flat directions.  In summary, we have investigated the evolution of  $Q$-balls, assuming that some of the constituent scalar fields decay into light particles. With a toy model of the flat direction, we have obtained the analytical solution of the $Q$-ball in the final state, and found that the spatial shape of the $Q$-ball does not change, but the orbit of the scalar field inside the $Q$-ball does change in order to compensate for the decay. In particular, the energy per unit charge generally increases due to the partial decay, which might induce further decay of the remnant scalar fields. Also we have performed numerical calculations to confirm these results. Again we stress that $Q$-balls remain $Q$-balls throughout the partial decay,  due to both the charge conservation of the remnant scalar fields and the $D$-flat condition.     \\subsection*"
    },
    "0403/astro-ph0403135_arXiv.txt": {
        "abstract": "We present a method for computing the 6-dimensional coarse-grained phase-space density $f(\\Bx,\\Bv)$ in an N-body system, and derive its distribution function $v(f)$.  The method is based on Delaunay tessellation, where $v(f)$ is obtained with an effective fixed smoothing window over a wide $f$ range.  The errors are estimated, and $v(f)$ is found to be insensitive to the sampling resolution or the simulation technique.  We find that in gravitationally relaxed haloes built by hierarchical clustering, $v(f)$ is well approximated by a robust power law, $v(f) \\propto f^{-2.5 \\pm 0.05}$, over more than 4 decades in $f$, from its virial level to the numerical resolution limit. This is tested to be valid in the $\\Lambda$CDM cosmology for haloes with masses $10^9-10^{15}\\msun$, indicating insensitivity to the slope of the initial fluctuation power spectrum.  By mapping the phase-space density in position space, we find that the high-$f$ end of $v(f)$ is dominated by the ``cold'' subhaloes rather than the parent-halo central region and its global spherical profile.  The value of $f$ in subhaloes near the virial radius is typically $>100$ times higher than its value at the halo centre, and it decreases gradually from outside in toward its value at the halo centre.  This seems to reflect phase mixing due to mergers and tidal effects involving puffing up and heating.  The phase-space density can thus provide a sensitive tool for studying the evolution of subhaloes during the hierarchical buildup of haloes.  It remains to be understood why the evolved substructure adds up to the actual universal power law of $v(f) \\propto f^{-2.5}$.  It seems that this behaviour results from the hierarchical clustering process and is not a general result of violent relaxation. ",
        "introduction": "\\label{sec:intro}  The standard paradigm assumes that dark-matter haloes are the basic entities in which luminous galaxies form and live.  The haloes dominate the gravitational potential over a wide range of radii and they have a crucial role in determining the galaxy properties. While many of the systematic features of halo structure and kinematics have been revealed by $N$-body simulations, the origin of these features is still an open theoretical puzzle, despite the fact that they are due to simple Newtonian gravity.  The halo \\emph{density profile} $\\rho(r)$ is an example for such a puzzle.  It is found in the simulations to have a robust non-power-law shape (which we refer to in general as ``NFW\"), with a log slope $-3$ near the virial radius, flattening gradually toward $-1$ at about 1\\% of the virial radius, and perhaps flattening further at smaller radii. \\citep[e.g.][and references therein]{ref:Nav97,ref:Mo98,ref:Kly01, ref:Pow03,ref:Hay04}.  This density profile is insensitive (or at most weakly sensitive) to parameters of the cosmological model and the initial fluctuation power spectrum \\citep[e.g.,][]{ref:Col96,ref:Nav97,ref:Sub00, ref:Ric03, ref:Col03, ref:Nav04} indicating that its origin is due to a robust relaxation process rather than specific initial conditions.  In particular, violent relaxation \\citep{ref:Lyn67} may be involved in shaping up the density profile.  In addition, secondary infall may be important in the outer regions \\citep{ref:Gun72, ref:Dek81, ref:Fil84, ref:Hof85, ref:Whi92} while some argue that resonances may have a role in the inner regions (\\citealt{ref:Wei02}; but see \\citealt{ref:Val03}). Nevertheless, there is no clear understanding for why the haloes actually pick up the particular density profile they have.  The properties of the \\emph{velocity dispersion} tensor is another theoretical puzzle. The haloes tend to be rather round, with a velocity dispersion profile that is slightly rising at small radii and slightly falling at large radii but is rather flat overall \\citep{ref:Hus99a, ref:Hus99b}. The profile of the anisotropy parameter $\\beta(r)$, which is a measure of radial versus tangential velocities, indicates near isotropy at small radii, which gradually develops into more radial orbits at large radii \\citep{ref:Col96, ref:Hus99a, ref:Hus99b, ref:Col00a}.  For a spherical system in equilibrium, the $\\sigma(r)$ and $\\beta(r)$ are related to $\\rho(r)$ via the Jeans equation, but it is not at all clear why the haloes in the simulations choose the characteristic $\\sigma(r)$ or $\\beta(r)$ that they actually do.  An interesting attempt to address the origin of the halo profile has been made by \\citet{ref:Tay01}, who measured a poor-man phase-space density profile by $f_{\\rm TN}(r) =\\rho(r)/\\sigma(r)^3$, and found that it displays an approximate power-law behaviour, $f_{\\rm TN} \\propto r^{-1.87}$, over more than two decades in $r$.  Using the spheri-symmetric Jeans equation, they showed that this power law permits a whole family of density profiles, whose limiting case is a profile similar to NFW, which asymptotically approaches a slope $-0.75$ as $r \\rarrow 0$. The general power-law shape of $f_{\\rm TN}(r)$ is confirmed in the simulated haloes described below. This scale-free behaviour of $f_{\\rm TN}(r)$ is intriguing, and it motivates further studies of halo structure by means of phase-space density.  Simulations of the $\\Lambda$CDM cosmology also reveal a roughly self-similar \\emph{hierarchical clustering} process, where smaller building blocks accrete and merge into bigger haloes.  At every moment, every halo contains a substructure of subhaloes on top of a smooth halo component that has been tidally stripped from an earlier generation of substructure.  Some of the important dynamical processes involved in this hierarchical halo buildup are understood qualitatively. These include, for example, the dynamical friction which governs the decay of the satellite orbits, the tidal stripping of subhaloes due to the host halo potential, and the mergers and flyby interactions of the subhaloes among themselves.  However, a complete theoretical understanding of how these processes work in detail, and how they combine to produce the halo structure and kinematics, is lacking.  Attempts have been made to explain an inner density cusp using toy models of dynamical stripping and tidal effects during the halo buildup by mergers \\citep{ref:Syr98,ref:Nus99,ref:ddh03,ref:Dek03}. However, a similar halo density profile seems to be produced also in some of the simulations where substructure has been artificially suppressed \\citep{ref:Mo99b, ref:Hus99a, ref:Avila01, ref:Bul02, ref:Alvarez02}, indicating that the process responsible for the origin of this density profile might be a somewhat more general feature of gravity and not unique to the merger scenario.   The issue of halo substructure has become timely both because of its relevance to observations and its implications on other major issues in galaxy formation.  Tidal tails and streams associated with dwarf satellite galaxies have been observed in the haloes of the Milky Way and M31 \\citep{ref:Iba94, ref:Iba01a, ref:Iba01b}, and they start to allow detailed modelling of the halo history through the satellite orbits. Gravitational-lensing observations provide preliminary indications for the presence of substructure in haloes at the high level predicted by the dissipationless $\\Lambda$CDM scenario \\cite[e.g.][]{ref:Dal02} In contrast, the observed number density of dwarf galaxies seems to be significantly lower, thus posing a ``missing dwarf problem'' \\cite{ref:Kly99b, ref:Mo99a}. Also, the ``angular-momentum problem'' of disk galaxies \\citep[e.g.][]{ref:Nav00,ref:Bul01a} is probably associated with the evolution of substructure in haloes \\citep{ref:Mal02, ref:mds02}. While the dwarf and angular-momentum problems necessarily involve baryonic processes, understanding the gravitational evolution of substructure is clearly a key for solving them.  In order to better understand the origin the various aspects of halo structure and buildup mentioned above, we make here a first attempt at addressing directly and in some detail the \\emph{phase-space} structure of dark-matter haloes.  The fundamental quantity in the dynamical evolution of gravitating systems is the full, six-dimensional, coarse-grained, phase-space density $f(\\Bx,\\Bv)$, which intimately relates to the underlying Vlasov equation and lies behind any relaxation process that may give rise to the virialized halo structure \\citep[][ chapter 4]{ref:Bin87}.  Ideally, one would have liked to compute it free of assumptions concerning spherical symmetry, isotropy, or any kind of equilibrium. However, computing densities in a six-dimensional space is a non-trivial challenge which requires simulations of a very broad dynamical range.  The state-of-the-art N-body simulations, with more than million particles per halo, allow an attempt of this sort for the first time.  We describe below a successful algorithm for measuring $f(\\Bx,\\Bv)$, and study its relevant properties including the associated systematic and random uncertainties.  We then apply this algorithm to simulated virialized haloes in the $\\Lambda$CDM cosmology.  We report in this paper two surprising new results. First, we discover that the volume distribution function of the phase-space density, $v(f)$, displays a universal scale-invariant \\emph{power-law} shape, valid in all virialized haloes that form by hierarchical clustering. Second, we realise that this power law is not directly related to the overall density profile, but is rather driven by the halo \\emph{substructure}.  This implies that the phase-space density provides a useful tool for studying the hierarchical buildup of dark-matter haloes and the evolution of substructure in them.  In \\se{vf} we introduce $f(\\Bx,\\Bv)$ and $v(f)$.  In \\se{computing} we describe the method of computing $f$ and $v(f)$ from an $N$-body halo, and summarise its properties and the associated errors, which are addressed in more detail in Appendix~\\ref{sec:errors}.  In \\se{universal} we present the universal power-law shape of $v(f)$ based on several different simulations, and demonstrate its robustness to the mass scale and simulation technique.  In \\se{substructures} we display maps of phase-space density and show that the high-$f$ contributions to $v(f)$ come from substructures within the parent halo.  In \\se{conc} we summarise and discuss our results and future work. ",
        "conclusions": "\\label{sec:conc}  Using Delaunay tessellation, we developed a method for measuring the 6-dimensional coarse-grained phase-space density $f(\\Bx,\\Bv)$ in $N$-body systems.  We focused, in particular, on measuring the phase-space volume distribution function, $v(f)$.  We applied this technique to several simulated haloes of $\\sim 10^6$ particles, formed by hierarchical clustering in the standard $\\Lambda$CDM scenario, and obtained two striking new results.  First, $v(f)$ is well described by a power law, $v(f) \\propto f^{-2.5 \\pm 0.05}$, over 3 to 5 decades in $f$.  The power-law regime starts at an $f$ value which corresponds to the characteristic size of the virialized halo.  It ends at an $f$ value which is determined by the dynamical resolution limit of the specific simulation.  Therefore, the true power-law range may extend to $f \\rightarrow \\infty$.  This power law seems to be insensitive to the halo mass in the range $10^9-10^{15}\\msun$, indicating insensitivity to the exact slope of the fluctuation power spectrum, as long as the haloes are built by hierarchical merging of clumps bottom up.  Second, this power-law originates from substructures within the halo rather then the overall trend with radius. The substructure completely dominates the high-$f$ parts of the $v(f)$ distribution.  The infalling clumps seem to phase-mix --- by puffing up, heating and stripping --- as their orbits decay from the virial radius inwards toward the halo centre and they melt into the halo smooth background.  Our first worry is that these results could be numerical artifacts, or severely contaminated by such.  Based on our error analysis and tests with mock datasets, we argue that the $v(f)$ measured by the DTFE algorithm genuinely reflects the true phase-space properties of the given $N$-body system over a broad range of $f$.  The question is whether the phase mixing suffered by the subclumps as they approach the halo centre might be an artifact of numerical effects such as two-body relaxation, leading to underestimated inner densities and/or overestimated internal velocities.  A similar effect has been pointed out using a one-dimensional toy model \\citep{ref:Bin03}.  The apparent agreement between simulations run with different codes and different resolutions is encouraging.  In order to specifically address the effect of few-body relaxation, we intend to run twice a simulation of the same halo with the same number of particles but with a different force resolution (ongoing work with F. Stoehr).  Assuming that the simulations genuinely reflect the true physical behaviour under the Vlasov equation, the origin of the robust power-law shape of $v(f)$ from the merging substructure becomes a very interesting theoretical issue.  As demonstrated in \\se{substructures}, a simple model using the mass function and the scaled profiles of the general halo population in the $\\Lambda$CDM scenario does not reproduce the correct power law.  This, and the apparent trend of the $f$ spikes with radius, indicate that the structural and kinematic evolution of the subhaloes in the parent halo are important. Studies of tidal heating and stripping may be found useful in this modelling.  It would be interesting to follow the phase-space evolution and the contribution to the overall $v(f)$ by a single, highly resolved subhalo, or many of those, as they orbit within the parent halo and approach its centre. This may help us understand the nature of the interaction between the parent halo and its subhaloes, and the origin of the $v(f)$ power law (ongoing works with E. Hayashi and with B. Moore).  Another more general but speculative possibility is that the $f^{-2.5}$ power law represents some sort of a cascade of relaxation processed in phase-space, in which high phase-space densities turn into lower (coarse-grained) densities through the process of mixing. In general, the fact that our findings are expressed in terms of the fundamental concept of phase-space density should make them more directly accessible to analytical treatment. In this respect, it may prove beneficial to investigate more closely the time evolution of the $v(f)$ of a cosmological halo and its components.  This may shed light on the connection between the $v(f)$ power-law behaviour and the relaxation processes within the halo.  We saw that the power-law behavior of $v(f)$ is limited to the virial regime.  It would be interesting to learn how this shape evolves in time as the halo virializes.  A preliminary study (to be concluded and reported in another paper) indicates that in the intermediate-$f$ regime the $v(f)$ of a pre-virialized system is significantly flatter than $f^{-2.5}$, while in the high-$f$ regime it drops in a much steeper way.  The $f^{-2.5}$ behaviour seems to be a feature unique to virialized systems.  We learnt that in the haloes that are built by hierarchical clustering, the power-law behaviour $v(f)\\propto f^{-2.5}$ reflects the halo substructure. It would be interesting to find out whether this power-law behaviour actually requires substructure, or it is a more general phenomenon of virialized gravitating systems, valid independently of substructure.  One way to answer this question would be to analyse simulated haloes in which all fluctuations of wavelengths smaller than the halo scale were removed, resulting in a smooth halo formed by monolithic collapse, with no apparent substructure in the final configuration.  As described in \\se{intro}, such haloes are known to still have NFW-like density profiles in real space, and one wonders whether they also have the magic power-law $v(f)$.  There are preliminary indications for a steeper $v(f)$ in this case (Arad, Dekel \\& Moore, in preparation).  If confirmed, it would indicate that the $f^{-2.5}$ behaviour, while insensitive to the exact slope of the initial power spectrum, is unique to the hierarchical clustering process, and is not a general result of violent relaxation.  Our current results are just first hints from what seems to be a promising rich new tool for analysing the dynamics and structure of virialized gravitating systems. The analysis could become even more interesting when applied to haloes including the associated gaseous and stellar components."
    },
    "0403/astro-ph0403303_arXiv.txt": {
        "abstract": "{We present an X-ray study of the high metallicity young open cluster Blanco\\,1 based on {\\em XMM--Newton} data. X-ray spectroscopy of cluster members is presented for the first time as well as new X-ray distribution functions of late-type stars. We detected all known dF and dG stars in the EPIC field and 80\\% and 90\\% of dK and dM stars, respectively. The X-ray spectral analysis of the X-ray brightest cluster stars and X-ray color analysis  of a larger sample show that a model with two temperatures (at about 0.3 and 1 keV) explains the quiescent activity phase spectra.  We discuss also the nature of unidentified X-ray sources in the observed region and their X-ray spectral properties. ",
        "introduction": "% Open clusters are powerful laboratories to test the models of star formation and evolution as well as the metal enrichment in the Galaxy. In fact, they are naturally selected samples of stars with same age, composition and environmental formation conditions. In the last two decades, X-ray observations of open clusters have assumed great importance. Since the early '80s, the open cluster X-ray observations of first {\\em Einstein} and subsequently ROSAT, showed that young cluster stars are stronger X-ray sources than the Sun (\\citealp[see][]{Micela2002} and reference therein cited). The evolution of X-ray activity is correlated with the evolutionary angular momentum losses, older stars being less luminous than the younger ones, but other, more subtle, factors like pre-main-sequence history and chemical composition could also play a role. The comparison of the properties of different clusters  is essential for an understanding of the importance of these factors. It is expected that stellar chemical composition influences the processes of coronal emission and activity: X-ray spectra of stars with high metal content are expected to have enhanced line emission contribution with respect to low metallicity stars. Moreover, the extent of the convective zone in solar type stars is also determined by metal content, thus resulting in a possible enhanced dynamo efficiency. Investigations of open clusters with different metallicity are of great importance for the study of the relation between stellar structure and X-ray activity.  Blanco\\,1 is a young open cluster noticeably more distant from the Galactic Plane (about 240 pc) than the young open cluster scale  height ($\\sim 100$ pc). Its age, around 100 Myr, is very similar to that of the Pleiades and NGC\\,2516 clusters \\citep{deEpst85,West88}, while its metallicity ([Fe/H]=+0.23 dex) is significantly higher than that of the Sun and the Pleiades \\citep{Edv95,Pan97,Jef99}. The high metallicity of Blanco\\,1 offers a suitable benchmark for the connection between stellar structure, activity and convection. Measurements of Li abundance by \\citet{Jef99} do not fit standard mixing models for stars of this age and metallicity, thus implying some Li-depletion inhibition. Our previous work (\\citealp{Giusi99}; \\citealp{bl1hri}, Paper I) based on ROSAT data, showed that the cluster dG and dK stars have an overall X-ray emission similar to that of the Pleiades while the dM stars in Blanco\\,1 appeared more luminous in the X-ray band with respect to the Pleiades and $\\alpha$\\,Per clusters. Prior to Paper I, membership catalogues for stars down to late K types were based only on photometry \\citep{deEpst85,Pan97}, while a few radial velocity measurements are given in \\citet{Jef99} and \\citet{Edv95}. In Paper I the membership of the cluster stars in the ROSAT field of view was defined by means of proper motion analysis and it was possible to derive X-ray luminosity distributions (XLDs) for dF, dG, dK and dM types. In that work, due to the limited sensitivity of ROSAT, the XLD of dM stars was significantly affected by the large number of upper limits. Furthermore, the lack of spectral capability of the ROSAT-HRI camera, did not allow us to derive any spectral features of the cluster coronae. The {\\em XMM-Newton} observation addresses these two issues: the large effective area of its three X-ray telescopes allows us to detect essentially all cluster members in the observed region and the moderate spectral resolution of the EPIC camera allows us to explore the main characteristics of the coronal spectra.  In this work we present the X-ray spectroscopy and photometric analysis of the cluster region as follows: in Sect. 2 we describe the observations, the basic processing of the data and the source detection results. In Sect. 3 we discuss the X-ray spectral analysis of the cluster stars and their XLDs. In Sect. 4 we discuss the nature of the unidentified sources and in Sect. 5 we summarize our conclusions. ",
        "conclusions": "% We have presented our analysis of an {\\em XMM-Newton} observation of the young open cluster Blanco\\,1 addressing the issues of the completeness of detection among X-ray cluster sources and, for the first time, of low resolution X-ray spectroscopy of its members.  We have detected 190 X-ray sources by using a wavelet based algorithm developed specifically for the EPIC camera. Detections have been derived on the combined MOS 1, 2 and {\\em p--n} datasets, making full use of the EPIC camera capability and improving the detection sensitivity. Of the 190 detected sources, 36 are identified with cluster members, while 5 cluster members remain undetected. As compared with previous ROSAT observations, the {\\it XMM-Newton} rate of detection has increased from 61\\% to 88\\%.  Six cluster sources with more than 1000 counts have been the subject of spectral analysis. A 2-T model fits the main spectrum features well and allows us to identify a ``cool'' component around 0.3 keV and a ``hot'' component around 1 keV, both with similar emission measures. These values are in good agreement with the XMM description of coronal spectra for young active stars like the Pleiades.  A 1-T model at $\\sim 0.6$ keV explains the soft spectrum of an intermediate type star (spectral type $\\sim$ A8) of the cluster. We consider that the X-ray emission is intrinsic to the A star itself and not due to an unresolved late companion.  By means of two X-ray color indices, the X-ray spectral properties of a larger cluster sample were explored and found to be consistent with the above 2-T model.  The X-ray luminosity distribution function of dM-type stars of the cluster indicates a probable difference in emission levels with respect to the Pleiades. Saturation of $\\log \\mathrm{L}_\\mathrm{X}/\\mathrm{L}_\\mathrm{bol}$ is not clearly evident down to mid-K stars, slightly later than in Pleiades.  Appendix A lists non-cluster optical counterparts for 90 of the 190 X-ray sources, identified either with the GSC-II and  USNO-B1 optical catalogs or with the 2MASS infrared catalog. One star is suggested as new probable low mass cluster member.  Listed in Appendix B are 64 sources that remain unidentified: most of them show X-ray colors quite different from those of the cluster, suggesting a different nature. However, for two of them the color analysis in Sect. 4 suggests agreement of their spectrum  with those of cluster coronae."
    },
    "0403/astro-ph0403629_arXiv.txt": {
        "abstract": "We present   results of an  analysis of the structural and kinematical properties of a sample of elliptical-like objects (ELOs) identified  in four hydrodynamical,   self-consistent simulations run with the DEVA code (Serna et al. 2003). Star formation has been implemented in the code through a simple phenomenological parameterization, that takes into account stellar physics processes only implicitly through the values of a threshold gas density, $\\rho_{\\rm g,thres}$, and an efficiency parameter, $c_*$. The four simulations operate in the context of a $\\Lambda$CDM cosmological model consistent with observations, resolve ELO mass assembly at scales up to $\\simeq$ 2 kpc, and differ in the values of their star formation parameters. Stellar masses, projected half-mass radii and central l.o.s. velocity dispersions, $\\sigma_{\\rm los, 0}$, have been measured on the ELO sample and their  values compared with data from the Sloan digital sky survey. For the first time in self-consistent simulations, a good degree of agreement has been shown, including the Faber-Jackson  and  the $D_n - \\sigma_{\\rm los, 0}$ relations, among others, but only when particular values of the  $\\rho_{\\rm g,thres}$ and $c_*$ parameters are used. This demostrates the effect that the star formation parameterization has on  the ELO mass distribution. Additionally, our results suggest that  it is not strictly necessary, at the scales resolved in this work, to appeal to energy sources other than gravitational (as for example   supernovae  feedback effects) to account for the structure and kinematics of large ellipticals. ",
        "introduction": "One of the most challenging open problems in modern cosmology is the origin of the local galaxies of different Hubble types  we observe to-day. Among them,  ellipticals are the easiest to study. They form the  most homogeneous family and show the most precise regularities in the form of correlations among some pairs of their observable parameters. The Sloan digital sky survey (SDSS, see York et al. 2000) sample of early-type galaxies, containing to date 9000 galaxies from different environments, provides a new standard of reference for nearby elliptical  galaxies. Its analysis confirmed correlations previously established, such as those involving  structural and kinematical parameters (the $L - \\sigma_{\\rm los, 0}$ or Faber-Jackson relation, 1976; the surface-brightness - effective radius relation, Kormendy 1977; the $D_n - \\sigma_{\\rm los, 0}$ relation, Dressler et al. 1987; among others, see Bernardi et al. 2003a, 2003b, 2003c, 2003d, and references quoted therein). These  correlations, as well as the [$\\alpha/$Fe] ratio trend with $\\sigma_{\\rm los, 0}$ (Jorgensen 1999), demand  short formation time-scales and old formation ages for the bulk of the stellar populations of ellipticals. These requirements are naturally met by one of the scenarios proposed so-far to explain galaxy formation and evolution: the so-called {\\it monolithic collapse} scenario (ellipticals would form  at high $z$ in a single burst of star formation, and would passively evolve since then; Patridge \\& Peebles 1967; Larson 1974). The competing {\\it merging  scenario} (galaxy mass assembly takes place gradually through repeated mergers of smaller subunits; Toomre 1977; Kauffmann 1996) meets some difficulties at explaining the correlations above as well as other observations on ellipticals, see Peebles 2002 and Matteucci 2003 for  details and discussions. But the monolithic collapse scenario does not recover all the currently available observations on ellipticals either. Such are, for example,  the range in ages their stellar populations span in some cases or  their kinematical and dynamical peculiarities (Trager et al. 2000; Menanteau, Abraham \\& Ellis 2001), indicating that an important fraction of present-day ellipticals have recently experienced merger events.  In order to reconcile all this observational background within a formation scenario, it is preferable to study galaxy assembly from simple physical principles and in connection with the global cosmological model. Self-consistent gravo-hydrodynamical simulations are a very convenient tool to work out this problem (Navarro \\& White 1994, Tissera, Lambas \\& Abadi 1997, Thacker \\& Couchman 2000). The method works as follows: initial conditions are set at high $z$ as a Montecarlo realization of the field of primordial fluctuations to a given cosmological model in a periodic, homogeneously sampled box; then the evolution of these fluctuations is numerically followed up to $z =0$ by means of a computing code that solves the N-body plus hydrodynamical evolution equations. In this way, the uncertainties resulting from prescriptions on dynamics and gas  cooling and heating, present in other methods such as semi-analytical ones (Kauffmann et al. 1999, Mathis et al. 2002) can be removed, only star formation needs further modelling. Individual galaxy-like objects (GLOs) naturally appear as an output of the simulations, no  prescriptions  are needed as far as their mass  assembly processes are concerned. Moreover, self-consistent simulations directly provide detailed information on each individual GLO  at each $ z$, namely  its  six dimensional phase space structure, as well as  the  temperature and age distributions of its gaseous and stellar components, respectively. From this information, the parameters  characterizing  each GLO can be estimated and compared with observations (see, for example, S\\'aiz et al. 2001, concerning disk galaxies). The first step in the program of studying the origin of galaxies through self-consistent simulations, is to make sure that they form GLO samples of different Hubble types that have counterparts in the real local Universe. In particular, the possible simple correlations  involving  structural and kinematical  parameters must be recovered. A detailed  analysis of this kind was not yet available for ellipticals (see previous work in Kobayashi 2002; Sommer-Larsen, Gotz,  \\& Portinari 2002; Meza et al. 2003).  We present in this paper the results of an analysis of the structure and kinematics of a sample of elliptical-like objects (ELOs) identified in four  self-consistent hydrodynamical simulations run in the framework of a flat $\\Lambda$CDM model consistent with observations. We have used DEVA, an  AP3M-SPH code particularly designed to study galaxy assembly in a cosmological context. In DEVA, special attention has been paid that the implementation of conservation laws (energy, entropy and angular momentum) is as accurate as possible (see Serna, Dom{\\'{\\i}}nguez-Tenreiro, \\& S\\'aiz, 2003 for details, in particular for  a discussion on the observational implications of violating some conservation laws). Star formation (SF) processes have been implemented in the code through a simple  parameterization, similar to that first used by Katz (1992), that includes a  threshold gas density, $\\rho_{\\rm g,thres}$ and an efficiency parameter, $c_{*}$, determining  the SF timescales according with a Kennicutt-Schmidt law\\footnote{See Elmegreen (2003) for a discussion on the possibility that this law can be explained as a result of SF processes at the scale of molecular cloud cores, through an interstellar medium  (ISM) gas structure whose density, prior to SF, can be described by a log-normal probability distribution, as Wada \\& Norman (2001) found in their simulations } (Kennicutt 1998).  Galaxy-like objects of different morphologies have been identified in the simulations (S\\'aiz, Dom\\'{\\i}nguez-Tenreiro \\& Serna 2002; S\\'aiz 2003). The aim of this paper is to show that some of  those identified as ELOs, have, at a structural and kinematical level, counterparts in the local Universe, including parameter  correlations. Data have been taken from the SDSS as analyzed by Bernardi  et al.\\ (2003a, 2003b), Kauffmann et al. (2003a, 2003b) and Shen et al., 2003. A brief account on ELO assembly is as follows: it mainly  occurs through a multiclump  collapse at rather high $z$ involving many clumps; collapse takes the clumps closer and closer along filaments causing them to merge at very low relative angular momentum and, consequently, at short timescales. This results into  fast SF bursts at high $z$ that transform most of the available gas into stars. The frequency of  head-on  mergers decreases with $z$. ELO stellar populations are mostly old, and a trend exists with $\\sigma_{\\rm los, 0}$, as suggested by some observations (Thomas, Maraston \\& Bender, 2002). ",
        "conclusions": "The degree of consistency reported here between sizes, velocity dispersions and stellar masses of simulated ELOs and SDSS data is very good. The agreement is particularly outstanding when the simplicity of our working scheme is recalled: ELO assembly has been simulated in the context of a cosmological model roughly consistent with observations; Newton laws and hydrodynamical equations have been integrated in this context, with a standard cooling algorithm and a SF parameterization through a Kennicutt-Schmidt-like law, containing our ignorance about its details at sub-kpc scales. No further hypotheses to model  the assembly processes have been made. Our results suggest that, at least  for the more massive objects, say $\\ge 6 \\times 10^{10}$ M$_{\\odot}$, it is not strictly necessary to appeal to energy sources, {\\it at the scales resolved in this work}, that is, up to $\\sim 2$ kpc, other than gravitational to account for their sizes, velocity dispersions, stellar masses and their correlations. This result is consistent with the idea that SNe explosions are not relevant at these scales, but only at smaller scales, to make shells and trigger star formation in molecular cloud cores. Their effect would be accounted for in the SF parameters (Elmegreen 2003).  This work strongly suggests that the SF parameterization is a key ingredient to determine the compactness of elliptical galaxies. A good agreement has been found in the correlations addressed here with SDSS data, but in the case of those involving sizes, only when  particular  values of the SF  parameters are used. Our results push the problem of elliptical galaxy formation from understanding their mass  assembly  at scales $>$  kpc in a cosmological context (Dom\\'{\\i}nguez-Tenreiro et al. 2003, in preparation), to understanding how SF was regulated at high $z$ as the gas falls within collapsing volumes, or other shock locations, so as to proceed just with  the   efficiency necessary to  produce the sizes observed in  to-day ellipticals.  It is a pleasure to thank J. Silk and J. Sommer-Larsen for useful information on the topics addressed in this paper. This work was partially supported by the MCyT (Spain) through grant AYA-0973 from the Programa Nacional de Astronom\\'{\\i}a y Astrof\\'{\\i}sica. We also thank the Centro de Computaci\\'on Cient\\'{\\i}fica (UAM, Spain) for computing facilities.   \\clearpage"
    },
    "0403/astro-ph0403073_arXiv.txt": {
        "abstract": "The large scale anisotropies of WMAP data have attracted a lot of attention and have been a source of controversy, with many of favourite cosmological models being apparently disfavoured by the power spectrum estimates at low $\\ell$.  All of the existing analyses of theoretical models are based on approximations for the likelihood function, which are likely to be inaccurate on large scales.  Here we present exact evaluations of the likelihood of the low multipoles by direct inversion of the theoretical covariance matrix for low resolution WMAP maps.  We project out the unwanted galactic contaminants using the WMAP derived maps of these foregrounds. This improves over the template based foreground subtraction used in the original analysis, which can remove some of the cosmological signal and may lead to a suppression of power. As a result we find an increase in power at low multipoles.  For the quadrupole the maximum likelihood values are rather uncertain and vary between 140-220$\\mu {\\rm K}^2$. On the other hand, the probability distribution away from the peak is robust and, assuming a uniform prior between 0 and $2000\\mu {\\rm K}^2$, the probability of having the true value above $1200\\mu {\\rm K}^2$ (as predicted by the simplest $\\Lambda CDM$ model) is 10\\%, a factor of 2.5 higher than predicted by WMAP likelihood code.  We do not find the correlation function to be unusual beyond the low quadrupole value. We develop a fast likelihood evaluation routine that can be used instead of WMAP routines for low $\\ell$ values.  We apply it to the Markov Chain Monte Carlo analysis to compare the cosmological parameters between the two cases.  The new analysis of WMAP either alone or jointly with SDSS and VSA reduces the evidence for running to less than 1-$\\sigma$, giving $\\alpha_s=-0.022\\pm 0.033$ for the combined case.  The new analysis prefers about 1-$\\sigma$ lower value of $\\Omega_m$, a consequence of an increased ISW contribution required by the increase in the spectrum at low $\\ell$.  These results suggest that the details of foreground removal and full likelihood analysis are important for the parameter estimation from WMAP data.  They are robust in the sense that they do not change significantly with frequency, mask or details of foreground template marginalization.  The marginalization approach presented here is the most conservative method to remove the foregrounds and should be particularly useful in the analysis of polarization, where foreground contamination may be much more severe. ",
        "introduction": "Data analysis of cosmic microwave background maps is a challenging numerical problem. The question that we want to answer is the probability (or likelihood) of a theoretical model given the data.  In order to evaluate the exact likelihood of a theoretical power spectrum of CMB fluctuations given a sky map of these fluctuations it is necessary to invert the theoretical covariance matrix. This operation scales as $O(N^3)$, where $N$ is the length of the data vector and is currently limited by practically available computer technology to $N \\lesssim 10^4$. One is hence forced to use approximate estimators when inferring the power spectrum from data such as WMAP satellite \\citep{2003ApJS..148....1B}, which have 1-2 orders of magnitude more independent measurements. The most popular methods are the pseudo-Cl (PCL) method (see e.g. \\citep{2002ApJ...567....2H}) and the Quadratic Maximum Likelihood (QML) estimator (see e.g. \\citep{1997PhRvD..55.5895T}). Both of these methods produce as an intermediate step estimates of multipole moments $C_\\ell$ and approximate methods have been developed to describe their probability distributions as accurately as possible \\citep{2000ApJ...533...19B,2003ApJS..148..195V}.  These perform satisfactorily for high $\\ell$ values, where the central limit theorem guarantees a Gaussian distribution (in offset lognormal transformed variables) will be a good approximation.  Unfortunately, these methods are much less reliable at low multipoles, where the distributions are not Gaussian. The situation is complicated further by the masks applied to the data to remove the galactic foreground contamination and by the marginalization of unwanted components, all of which makes analytic approach unreliable. In \\cite{2003astro.ph..7515E} it was suggested to  use a hybrid approach using QML on degraded maps at low $\\ell$ and PCL at higher multipoles.  The issue of the exact values of multipole moments in WMAP data has attracted much attention since the original analysis by WMAP team \\cite{2003ApJS..148..175S}. Several unusual features have been pointed out already in the original analysis.  One of these was the correlation function, which appears to almost vanish above $60^{\\circ}$.  Another was the low value of the quadrupole. With the PCL analysis the value of the quadrupole was found to be $\\sim 123 \\mu {\\rm K}^2$, compared to the expected value of $\\sim 1200 \\mu {\\rm K}^2$ for the simplest $\\Lambda CDM$ model. The probability for this low value was estimated to be below 1\\%, depending on the parameter space of models. The discussion of the statistical significance of the low values of quadrupole and octopole in the WMAP data \\citep{2003MNRAS.346L..26E,2003JCAP...07..002C,2003JCAP...09..010C,PhLB..570..145L} has sparked a renewed interest into the so called estimator induced variance \\citep{2004MNRAS.348..885E} - the error in the likelihood evaluation arising due to the use of an estimator rather than the exact expression. In \\cite{2004MNRAS.348..885E} it has been argued that [QML estimator performs significantly better than the PCL estimator and that the true value of the quadrupole probably lies in the range around $170-250 \\mu {\\rm K}^2$.  However, only the maximum likelihood value was computed and not the full likelihood distributions so the statistical significance of this result and its effect on the parameter estimation remained unclear. In addition, the role of foregrounds and monopole/dipole removal has not been explored in detail.  In this paper we take a different approach.  We argue that the actual value of the best fitted quadrupole (and other multipoles) is not of the main interest, since it can be quite sensitive to the details of the foreground subtraction procedure, type of mask used and numerical details of the analysis (in fact, the various values proposed so far may even be statistically indistinguishable if the likelihood function at the peak is very broad). What is more important is the probability or likelihood of a model given the data, compared to another model that may, for example, fit the data better.  This is encapsulated in the likelihood ratio between models and within the Bayesian context is the only information we really need to asses the viability of cosmological models that belong to a certain class. In this paper we perform the exact likelihood calculation by a direct inversion of the covariance matrix for the low resolution maps, thus eliminating all the uncertainties related to estimator variance approximations. Since we use low resolution maps with less than 3000 pixels we can do the inversions with a brute force linear algebra routines.  This means we cannot do the analysis on all of the multipole moments, so we analyse low multipoles with the exact method and use PCL analysis for the higher multipoles, where the two methods agree with each other and where the approximate variance estimates developed for PCL analysis are likely to be valid.  Second issue we wish to address in this paper is the question of foreground subtraction. This is done in two steps. First, pixels with high degree of contamination are completely removed from the data. This results in the so called KP2 (less aggressive, 85\\% of the sky) and KP0 (more aggressive, 75\\% of the sky) masks \\citep{2003ApJS..148...97B}.  There remains contamination even outside these masks in individual frequency channels. This contamination can be further reduced using templates and/or frequency information \\citep{2003ApJS..148...97B}. In WMAP analysis the templates were fitted for and subtracted out of WMAP data.  Even with a perfect template there is a danger that this procedure can oversubtract the foregrounds, since one is essentially subtracting out the maximum amplitude consistent with the template which could include some of the signal.  Instead, here we do not subtract out the templates, but marginalise over them by not using any information in the data that correlates with a given template. This procedure has not been applied to WMAP data in previous analyses. It guarantees that there is no statistical bias caused by the foreground removal.  Some of the templates that were subtracted in WMAP analysis, particularly 408MHz Haslam synchrotron radiation map \\citep{1982A&AS...47....1H}, are of poor quality. WMAP produced a better set of templates applying Maximum Entropy Method (MEM) to WMAP maps in several frequency channels using templates as priors only \\citep{2003ApJS..148...97B}. In addition to the Haslam synchrotron map, they used \\cite{2003ApJS..146..407F} H-$\\alpha$ map as a tracer of free-free emission and the SFD dust template based on \\cite{1998ApJ...500..525S}.  This process resulted in three MEM derived foreground maps.  These, however, were not used to infer the power spectrum. Instead, the official power spectrum was determined from the integrated single frequency maps and the same templates that were used as priors for the MEM map making procedure, ignoring the MEM derived maps.  The MEM derived maps are likely to be the most faithful representation of the foregrounds. When used in power spectrum inference, however, they must be used with care due to complicated nature of their signal and noise correlations \\citep{2003ApJS..148...97B}. Nevertheless, on the largest scales, where receiver noise is negligible, they are probably the best available option. We therefore use the integrated single channel maps and the MEM derived foreground templates as a basis of our work. Note that in foreground marginalization procedure no template is actually removed from the data and there is no danger of introducing noise correlations that could significantly affect the power spectrum estimates. We perform this process on foreground unsubtracted maps of the V and W channels of the WMAP satellite. We use both KP2 and KP0 masks and project out the remaining galactic contamination using MEM inferred maps of dust, synchrotron and free-free foregrounds. We use the likelihood evaluated in this way to asses the statistical significance of departure from the concordant model at low multipoles and to perform the statistical analysis of cosmological models given the data.   WMAP team also produced the so called \\emph{Internal Linear Combination} (ILC) map of the CMB emission, by using internal maps at various frequencies to decompose them into CMB and foreground components.  This approach is not based on any templates and so uses less information than in principle available. While visually these maps appear to be relatively free of contamination outside the galactic plane, there are still artifacts within the plane. This means that one must be careful when projecting out monopole and dipole: one should not simply remove them from the all-sky map, since they could be contaminated by galactic emission at the canter and this would leave a residual offset outside the galactic plane, which could contaminate all of low multipoles. One must again apply the marginalisation over monopole and quadrupole on the masked map to eliminate any contamination in the final result. A similar approach has been taken by \\cite{2003PhRvD..68l3523T}  and \\cite{2004astro.ph..3098E}, who produced their own versions of ILC maps. Since we argue that the best method is to use single frequency maps together with correct templates and we use ILC map for illustration and cross-check purposes only, we do not consider these alternative ILC solutions further. ",
        "conclusions": "In this paper we have developed routines to calculate the exact likelihood of the low resolution WMAP data.  We have projected out unwanted foreground components by adding the foreground templates to our covariance matrix with large variance.  Both of these methods have not been applied to WMAP data before and should improve upon the existing analyses.  We have tested the robustness of our results by applying the method to many different combinations of observing frequency, mask, smoothing and templates and found consistent results among these various cases. In particular, we find consistent results if we marginalize only over dust in $W$ channel as opposed to all 3 foreground templates, if we use templates external to WMAP instead of WMAP MEM templates, if we use KP0 instead of KP2 mask, if we use ILC maps instead of individual V or W frequencies or if we use Healpix windows instead of gaussian smoothing. The two most important features of our procedure are thus marginalization over dust and exact likelihood analysis.  Important differences exist between our results and previous work. We find higher values of the lowest multipoles, which is partly a consequence of template subtraction method used in WMAP analysis. This procedure would certainly remove some of the real power, although it is difficult to estimate how much and the differences could also be just a statistical fluctuation.  For the maximum likelihood value of the quadrupole we find values between the original WMAP analysis and subsequent reanalysis by \\cite{2004MNRAS.348..885E}. The differences are within the estimated error of the foreground contamination and we argue that the actual value is not very reliable given how broad the likelihood is at the peak.  More important is the shape of the likelihood function, which we find to be broader than in the WMAP team provided likelihood evaluation, which underestimates the errors compared to our analysis.  This lowers the statistical significance of the departure of the data from the concordant model. Within a Bayesian context and assuming a flat prior on the distribution of quadrupoles we find the probability that a model exceeds the concordance model predicted quadrupole to be 10\\%.  We also do not find anything particularly unusual in the correlation function and in the joint quadrupole-octopole analysis.  We combine the full likelihood calculation with foreground marginalization at low $\\ell$ with the original WMAP PCL analysis at high $\\ell$ to generate Monte Carlo Markov Chains, whose distribution converges to the probability distribution of theoretical models given the data and assumed priors.  The main effect of the new analysis is on the running of the spectral index, for which the marginal 2 sigma evidence for $\\alpha_s<0$ present in the original analysis and in the recent analysis of WMAP+VSA \\cite{2004astro.ph..2466R} (see also \\cite{2004astro.ph..2359R}) is reduced to below 1 sigma. Using the exact WMAP likelihood analysis will be essential for attempts to determine the running of the spectral index by combining WMAP with either the small scale CMB data or with the upcoming Ly-$\\alpha$ forest analysis from SDSS. In all of these cases the exact method increases the value of the running by pushing up the CMB spectrum at large scales. Another parameter which is significantly affected is the matter density $\\Omega_m$ or, equivalently, the dark energy density $\\Omega_{\\Lambda}$. We find $\\Omega_m$ to be reduced by the new analysis because of the added power at low multipoles, which is most easily accounted for by an increase in ISW contribution.  We have shown that the effects of the improved likelihood analysis presented here can be significant for the determination of cosmological parameters. We expect the methods applied here will be equally important for the analysis of polarization data in WMAP, where the foregrounds play a much more important role and where a full likelihood analysis of joint temperature and polarization data is necessary to extract the maximum amount of information. Current analysis of temperature-polarization data is rather unsatisfactory, since it is based on the cross-spectrum information alone.  Without having access to the full polarization maps we cannot improve upon it here. Thus the results shown in tables 1-2 should still be viewed as preliminary regarding the optical depth, which is essentially determined by the polarization data. Upcoming WMAP 2-year analysis/release of polarization data should elucidate the current situation. The code developed here will be made available to the community at \\texttt{cosmas.org}."
    },
    "0403/astro-ph0403559_arXiv.txt": {
        "abstract": "We report the non-detection of the \\cii\\ \\fscii\\ 157.74 $\\mu$m transition in the $z=6.42$ quasar SDSS J1148+5251 after 37.5 hours of integration with the James Clerk Maxwell Telescope.  This transition is the main cooling line of the star-forming interstellar medium, and usually the brightest FIR line in galaxies. Our observed 1 $\\sigma$ RMS = 1.3 mK in the $T_{A}^{*}$ scale translates to $\\Lcii<2.6 \\times 10^{9}$ L$_\\odot$.  Using a recent estimate of the far-infrared continuum of this quasar, we derive for SDSS J1148+5251 $\\Lcii/\\Lfir<5\\times10^{-4}$, a ratio similar to that observed in local ultra-luminous infrared galaxies but considerably smaller than what is typical in nearby normal and starburst galaxies. This indicates that the small \\Lcii/\\Lfir\\ ratio observed locally in luminous far-infrared objects also persists at the highest redshifts. ",
        "introduction": "The most distant quasar in the universe known at the time of this writing is SDSS J114816.64+525150.3 (hereafter SDSS J1148+5251) at a redshift $z=6.42$ (Fan et al. 2003; Walter et al. 2003; Bertoldi et al. 2003a). With a bolometric luminosity of L$_{\\rm bol}\\sim10^{14}$ L$_\\odot$, this is an extremely luminous object powered by a $\\sim3\\times10^{9}$ M$_\\odot$ supermassive black hole at its core (Willott et al. 2003).  The high far-infrared (FIR) luminosity of this object implies that the host galaxy is forming stars at the prodigious rate of 3000 M$_\\odot$ yr$^{-1}$ (Bertoldi et al. 2003b). At this redshift the universe is only 840 million years old and both the black hole mass and star formation rate imply the existence of a very massive galaxy formed from one of the rarest high-density peaks in the matter distribution. Although there are hints that SDSS J1148+5251 may be weakly gravitationally lensed (White et al. 2003), we assume no lensing magnification in this paper.  The recent detection of bright molecular and FIR continuum emission from the host galaxy of SDSS J1148+5251 (Walter et al. 2003; Bertoldi et al. 2003a; Bertoldi et al. 2003b) prompted us to observe the \\cii\\ \\fscii\\ 157.74 $\\mu$m fine structure transition in this object. This is the main cooling transition of the star-forming interstellar medium (e.g., Tielens \\& Hollenbach 1985), and is commonly the single brightest emission line in galaxies.  For example, Stacey et al. (1991) found that most galaxies emit 0.1\\% to 1\\% of their far-infrared luminosities in this line alone.  Given its extremely high luminosity and its relationship to star formation activity, the redshifted \\cii\\ 158 $\\mu$m transition is attractive as a star formation indicator at high-$z$ (e.g., Stark 1997).  This transition is conveniently placed into the atmospheric 1~mm window for redshifts $z\\gtrsim6.2$, making it accessible to ground-based instrumentation.  Furthermore, a spectroscopic approach to detecting high-$z$ sources has the clear advantage of containing redshift information, unlike continuum observations, thus avoiding problems of source confusion that stem from the low angular resolution of single-dish radio telescopes equipped with bolometer arrays.  As early as 1997, however, Infrared Space Observatory (ISO) observations suggested potential problems with this method.  Indeed, Malhotra et al. (1997) and Luhman et al. (1998) found that the proportionality between the \\cii\\ and FIR continuum luminosities observed in nearby star forming galaxies broke down for luminous infrared galaxies, where the \\cii\\ luminosity appeared not to exceed $\\Lcii\\sim10^9$ \\Lsun. These conclusions, confirmed and expanded in a recent analysis of ISO data by Luhman et al. (2003), have since cast doubts on the usefulness of redshifted \\cii\\ to probe star formation in the distant universe. The observations presented here have bearing on this matter: does the paucity of \\cii\\ emission observed in local ultra-luminous infrared galaxies (ULIRGs) hold for objects in the early universe? ",
        "conclusions": "Our redshifted \\cii\\ observations at $\\nu\\sim256$~GHz have failed to detect the fine structure 158 $\\mu$m \\cii\\ transition in this high-$z$ quasar. We place a 1 $\\sigma$ limit on the \\cii\\ luminosity of SDSS J1148+5251 of $\\Lcii<2.6\\times10^9$ \\Lsun\\ (including 30\\% calibration uncertainty). Given the observed FIR flux of this source, this places an upper limit on its \\cii\\ to FIR ratio of $\\Lcii/\\Lfir<5\\times10^{-4}$, substantially lower than what is found in most local star-forming galaxies but similar to what is observed in nearby ultra-luminous IR galaxies.  If the cause of this \\cii\\ deficit in SDSS J1148+5251 was a large FIR component due to partial dust reprocessing of the AGN radiation, thus unrelated to star formation activity, the $\\sim$3000 M$_\\odot$ yr$^{-1}$ star formation rate obtained from the FIR continuum could be a substantial overestimate. Obtaining a better SED for the FIR continuum emission, to constrain better the FIR luminosity, as well as improving the sensitivity of the CO \\jone\\ observations, may help to elucidate this matter.  Will redshifted \\cii\\ observations open a new window onto high-$z$ galaxies?  Our results are based on observations of just one object at high redshift. A larger sample of $z>6$ sources with high FIR luminosities must be observed in the \\cii\\ transition to establish whether low $\\Lcii/\\Lfir$ ratios are a general phenomenon.  If a maximum luminosity $\\Lcii\\sim10^9$ \\Lsun\\ were confirmed for most sources, the Atacama Large Millimeter Array would still be able to detect the \\cii\\ line and spatially resolve its distribution and kinematics. The answer, then, is a hopeful {\\em yes}."
    },
    "0403/astro-ph0403245_arXiv.txt": {
        "abstract": "{White dwarfs in globular clusters offer additional possibilities to determine distances and ages of globular clusters, provided their spectral types and masses are known. We therefore started a project to obtain spectra of white dwarfs in the globular clusters NGC~6397 and NGC~6752.  All observed white dwarfs show hydrogen-rich spectra and are therefore classified as DA. Analysing the multi-colour photometry of the white dwarfs in NGC~6752 yields an average gravity of \\logg\\ = 7.84 and 0.53\\Msolar\\ as the most probable average mass for globular cluster white dwarfs. Using this average gravity we try to determine independent temperatures by fitting the white dwarf spectra.  While the stellar parameters determined from spectroscopy and photometry usually agree within the mutual error bars, the low resolution and S/N of the spectra prevent us from setting constraints stronger than the ones derived from the photometry alone. For the same reasons the white dwarf spectra obtained for NGC~6397 unfortunately do not provide an independent distance estimate of sufficient accuracy to distinguish between the long and short distance  scale for globular clusters. ",
        "introduction": "As white dwarfs are the final stage in the evolution of all low mass stars, many are expected to exist in globular clusters.  However, due to their faintness and occurrence in very crowded fields they were detected only after the refurbishment of HST with the WFPC2. Several candidate white dwarfs were then soon identified in M15 (de Marchi \\& Paresce \\cite{depa95}), $\\omega$ Cen (Elson et al. \\cite{elgi95}), NGC 6397 (Paresce et al. \\cite{pade95a}, Cool et al. \\cite{copi96}), M4 (Richter et al. \\cite{rifa95}, \\cite{ribr02}), 47 Tuc (Paresce et al. \\cite{pade95b}, Zoccali et al. \\cite{zore01}) and NGC 6752 (Renzini et al. \\cite{rebr96}). Renzini \\& Fusi Pecci (\\cite{refu88}) suggested to use the white dwarf sequence as a standard candle for determining the distance to nearby globular clusters, similarly to the traditional main sequence fitting procedure using local subdwarfs with known trigonometric parallax. In this case, the white dwarf sequence of the cluster is compared to a sequence constructed with local white dwarfs with accurate trigonometric parallax. The method was then applied to NGC 6752 (Renzini et al. \\cite{rebr96}) and 47 Tuc (Zoccali et al. \\cite{zore01}, where also the result from Renzini et al. \\cite{rebr96} for NGC~6752 was updated). However, while the updated distance modulus for NGC~6752 agrees well with all other distance determinations for this cluster, the white dwarf distance to 47~Tuc is considerably shorter than that found by Gratton et al. (\\cite{grfu01}, \\cite{grbr03}) via main sequence fitting.  While it may seem strange to use the faintest objects in a globular cluster to derive its distance white dwarfs offer some advantages as standard candles when compared to main sequence stars: \\begin{itemize} \\item They come in just two varieties - either hydrogen-rich (DA) or helium-rich (DB) -- {\\em independent of their original metallicity} and, in both cases, their atmospheres are virtually free of metals. So, unlike in the case of main sequence fitting, there is not the problem of finding local calibrators with the same metallicity as the globular clusters. \\item White dwarfs are locally much more numerous than metal-poor main sequence stars and thus allow to define a better reference sample. \\end{itemize}  However, the method has its own specific problems, which are discussed in great detail in Zoccali et al. (\\cite{zore01}) and Salaris et al. (\\cite{saca01}). Indeed, the location of the white dwarf cooling sequence depends on:  \\begin{itemize} \\item {\\em the white dwarf mass}\\\\ On theoretical grounds, given the observed maximum luminosity reached on the asymptotic giant branch (AGB), the mass of currently forming white dwarfs in globular clusters should be $0.53\\pm 0.02$\\Msolar\\ (Renzini \\& Fusi Pecci \\cite{refu88}, Renzini et al. \\cite{rebr96}). Unfortunately, there are no local white dwarfs in this mass range with directly determined masses (i.e. without using a mass-radius relationship). There is, however, a handful of local white dwarfs with spectroscopically determined masses near this value (cf. Table 1 in Zoccali et al. \\cite{zore01}), which allows to construct a semi-empirical cooling sequence for $M_{\\rm WD}$=0.53\\Msolar, once relatively small mass-dependent corrections are applied to each local white dwarf. The spectroscopic determination of the mass of white dwarfs in a globular cluster was first attempted by Moehler et al. (\\cite{mohe00}) for white dwarfs in NGC~6397. However, the low S/N of the spectra of these very faint stars did not allow to determine the mass with sufficient accuracy.  \\item {\\em the white dwarf envelope mass}\\\\ In the case of DA white dwarfs the cooling sequence location depends also on the mass of the residual hydrogen-rich envelope. This affects {\\it spectroscopically} derived masses (see above), with the resulting mass being $\\approx 0.04$\\Msolar\\ higher when using the {\\it evolutionary} envelope mass ($\\approx 10^{-4}$M$_{\\rm WD}$, Fontaine \\& Wesemael \\cite{fowe97}) as opposed to virtually zero envelope mass. This mass uncertainty corresponds to an uncertainty of \\magpt{0}{1} in the distance modulus and 1--1.5 Gyr in the age derived from the main sequence turnoff.  \\item {\\em spectral type}\\\\ DB stars are fainter than DA stars at a given colour, with the offset depending on the filter combination (i.e. the offset is greater in $V$ vs. $B-V$ than in $I$ vs. $V-I$). However, more massive DA white dwarfs are also fainter at a given colour. \\end{itemize}  The white dwarf sequence also allows to determine the age of a globular cluster if its faint end is detected. In that case one can derive the age of the globular cluster from the luminosity of its oldest and faintest white dwarfs. Aside from the observational difficulties and the uncertainties in the cooling tracks (see Chabrier et al. \\cite{chbr00} for more details) any error in the assumed mass affects the result. Recent very deep HST observations allowed to detect the white dwarf cooling sequence in M~4 to unprecedented depths of $V$$\\approx$30 (Richer et al. \\cite{ribr02}).  As a preliminary result Hansen et al. (\\cite{habr02}) derive an age of 12.7$\\pm$0.7 Gyr from the white dwarf luminosity function of M~4, consistent with other independent age estimates (their error bar, however, does not include errors due to the uncertainty of the white dwarf mass). Their result has recently been questioned by de Marchi et al. (\\cite{depa03}), who claim that the cluster membership of the white dwarfs can not be verified down to sufficiently faint limits to obtain more than a lower limit of the age. Richer et al. (\\cite{ribr04}) showed, however, that different methods of data reduction and analysis account for the different depths reached with the same data and thereby defended the original result of Hansen et al. (\\cite{habr02}).  In view of the relevance of the white dwarf masses and spectral types to the problems described above we decided to observe spectra of white dwarfs in NGC~6397 and NGC~6752 in order to determine their spectral types and get mass estimates. Pilot observations of the white dwarf candidates in NGC~6397 (63.H-0348) had already shown that the targets are hydrogen-rich DA white dwarfs (Moehler et al. \\cite{mohe00}), but did not allow much quantitative work. ",
        "conclusions": "We observed a sample of white dwarfs in the globular clusters NGC~6397 and NGC~6752 and showed that they are all hydrogen-rich DA. From multicolour photometry we determined an average mass of 0.53$\\pm$0.03\\Msolar\\ (uncertainty due to uncertainties of 0.02 dex in the average \\logg\\ and of \\magpt{0}{05} in the distance modulus of NGC~6752). This value (based on the assumption of a thick hydrogen layer) is identical to that assumed by Renzini et al. (\\cite{rebr96}) on theoretical grounds. Therefore both observational and theoretical arguments strongly advocate against the use of 0.6~\\Msolar\\ (the mean mass of the local white dwarfs) or even more for the mass of hot white dwarfs in globular clusters when comparing observations to theoretical tracks. However, the limited S/N in combination with the low resolution of the spectroscopic data prevented the independent determination of masses from spectroscopic fits alone.  Multi-colour photometry may be the better way to determine the physical parameters of white dwarfs in globular clusters once their spectral types are known. For spectroscopic observations our experience shows that crowding can severely limit the usefulness of the data due to problems with sky subtraction. We would therefore strongly recommend to look for white dwarf candidates in ground-based wide-field photometry to avoid the problems we encountered for NGC~6752. Also a better sensitivity in the blue would be useful for future spectroscopy.  \\begin{table}[!h] \\caption[]{Results for the brightest object in NGC~6397 with \\logg\\ as free parameter. We also give M$_V$ from the model fit (assuming a thick hydrogen layer), the observed $V$ corrected for a reddening of E$_{\\rm B-V}$ = \\magpt{0}{18}, and the derived distance modulus. \\label{par-spec3}} \\begin{tabular}{lrrrrrr} \\hline \\hline & \\teff & $\\Delta$\\teff & \\logg & $\\Delta$ \\logg & M$_V$ & (m-M)$_0$  \\\\ & [K] & [K] & & & &  \\\\ \\hline DK & 18800 & 340 & 7.72 & 0.06 & \\magpt{10}{47} & \\magpt{11}{72}\\\\ RN & 18700 & 520 & 8.09 & 0.12 & \\magpt{11}{02} & \\magpt{11}{17}\\\\ \\hline \\end{tabular} \\end{table}"
    },
    "0403/hep-ph0403293_arXiv.txt": {
        "abstract": "We compute the neutrino detection rates to be expected at a low-energy beta-beam facility. We consider various nuclei as neutrino detectors and compare the case of a small versus large storage ring. ",
        "introduction": "The pioneering experiment of R. Davis \\cite{davis} has started the era of neutrino astronomy. Because they only have weak interactions with matter, neutrinos are precious messengers of what happens in the interior of stars, like our sun, or in explosive phenomena, such as Supernova type II explosions. Such astronomical neutrinos therefore provide an important source of information for our understanding of the life and death of stars. Nuclei are commonly used as detectors in neutrino observatories as well as in various experiments aiming at studying intrinsic neutrino properties, such as their masses and mixings. A precise knowledge of neutrino-nucleus cross-sections is needed for the interpretation of these measurements and/or to study the feasibility of new projects. The understanding of neutrino-nucleus interactions is also of crucial importance for various astrophysical processes. Timely examples include neutrino nucleosynthesis \\cite{haxton,woosley}, or the nucleosynthesis of heavy elements during the so-called r-process \\cite{bahanunucl,gail,bahaastro,qian,goriely}. If the latter takes place during the explosion of Supernovae type II, where a gigantic amount of energy is emitted as neutrinos of all flavors, final abundances depend on several nuclear properties, among which the interactions with neutrinos.  According to existing simulations, the average energy of neutrinos emitted from core-collapse Supernovae is about $10$~MeV for electron neutrinos and about $20$~MeV for muon and tau neutrinos \\cite{raffelt}. Notice however that, due to oscillations, electron neutrinos can become hotter while traversing the star \\cite{hax1999,smirnov,nuPb}. The predicted spectra cover the $50$~MeV region and present a tail up to about $100$~MeV \\cite{raffelt}. Reactor and solar neutrinos have typical energies in the $10$~MeV energy range, while accelerator and atmospheric neutrinos cover the GeV and multi-GeV range. The various theoretical approaches employed to describe neutrino-nucleus interactions therefore involve nuclear as well as nucleonic degrees of freedom (for a review, see \\cite{revuekubo,jpg}). There are a number of open issues in this context. The A=2 system is the simplest case, for which the reaction cross sections can be estimated with high accuracy \\cite{kubodera}. However, there is still an important quantity, namely L$_{1,A}$, related to the axial two-body current, which dominates the theoretical uncertainty in neutrino-deuteron interactions. For heavier nuclei, in the tens of MeV energy range, the reaction cross sections are dominated by collective modes, like the Gamow-Teller resonance or the Isobaric Analog State, which have been extensively studied in the past \\cite{osterfeld}. As the neutrino impinging energy increases, transitions to states of higher multipolarity (such as the spin-dipole or higher forbidden transitions) become important \\cite{volpelead}. The latter also play an important role in the context of core-collapse Supernova physics \\cite{kolbe,gail,jon,volpelead}. Although some information on these states can be gathered through other probes, such as charge-exchange reactions \\cite{osterfeld}, muon capture \\cite{giai}, or inelastic electron scattering \\cite{electron}, the experimental information  is rather scarce. Note that the understanding of neutrino-carbon reactions with neutrinos produced from the decay in flight of pions is still an open issue, for most of the theoretical calculations over-estimate the experimental value \\cite{c12}. So far, measurements with low-energy neutrinos have been performed in a few cases only, namely deuteron \\cite{deut}, carbon \\cite{expc12}, and iron \\cite{iron}. Systematic studies would be of great importance both for what concerns the interpolation from the MeV to the GeV neutrino energy range and the extrapolation to neutron-rich nuclei, as required in the astrophysical context.  Neutrino-nucleus interaction studies were one of the main physics issues of the proposed ORLAND underground neutrino facility, which was based on a conventional neutrino source (pion and muon decays) \\cite{orland,jpg}. A smaller version of the ORLAND project  is now under study \\cite{efremenko}. At present, the MINER$\\nu$A project \\cite{minerva} includes the study of neutrino-nucleus interactions for neutrino energies in the GeV range. Here, we study the potential of a low-energy neutrino facility based on beta-beams, a novel method to produce neutrino beams \\cite{zucchelli}. This consists in boosting exotic ions which decay through beta-decay and produce pure, collimated and well-understood electron neutrino fluxes. Such a method could be exploited for a future facility at CERN \\cite{zucchelli,mats}. High energy beta-beams would be fired to a gigantic Cherenkov detector like UNO \\cite{uno}, located in an (upgraded) Fr\\'ejus underground laboratory to study, in particular, the possible existence of CP violation in the leptonic sector \\cite{zucchelli,mats,mauro}. The discovery potential with a very high $\\gamma$ and a longer baseline is discussed in \\cite{jj,tmms}.  It has recently been proposed to use the beta-beam concept for the production of low-energy neutrinos \\cite{lownu}. Several laboratories will produce intense exotic beams in the near future and could, therefore, be possible sites for a low-energy beta-beam facility. These include GANIL, CERN, GSI, as well as the EURISOL project. Low-energy neutrino beams would offer an interesting opportunity to study various neutrino properties, such as e.g.\\ the neutrino magnetic moment \\cite{munu}, as well as neutrino-nucleus interactions, of interest for nuclear physics, particle physics and astrophysics. In the former case, one would exploit the ions at rest as an intense neutrino source, whereas, in the latter case, one would use boosted ions, which would be collected in a storage ring \\cite{lownu}, as in the original high energy proposal. An important feature of such beta-beams is that the boost factor of the accelerated ions can be varied, allowing one to explore various neutrino energy ranges.  In this paper, we present for the first time charged-current neutrino-nucleus interaction rates achievable at a low-energy beta-beam facility. We consider two possible cases for the dimensions of the storage ring, for which we inspire ourselves of the one planned in the future GSI facility \\cite{gsi} and the one thought in the CERN baseline scenario \\cite{zucchelli,mats}. We consider various target nuclei as neutrino detectors, namely deuteron, oxygen, iron and lead, which are commonly used in existing or planned experiments \\cite{orland}. Related work in the case of lead can be found in \\cite{gailnew}. ",
        "conclusions": ""
    },
    "0403/hep-ph0403308_arXiv.txt": {
        "abstract": "We consider the infrared modification of gravity by  ghost condensate. Naively, in this scenario one expects sizeable modification of gravity at distances of order 1000 km, provided that  the characteristic time scale of the theory is of the order of the Hubble time. However, we argue that this is not the case. The main physical reason for the conspiracy is a simple fact that the Earth (and any other object in the Universe) has  velocity of at least of order $10^{-3}c$ with respect to the rest frame of  ghost condensate. Combined with strong retardation effects present in the ghost sector, this fact implies that no observable modification of the gravitational field of nearby objects occurs. Instead, the physical manifestation of ghost condensate is the presence of ``star tracks'' --- narrow regions of space with growing gravitational and ghost fields inside --- along the trajectory of any massive object. We briefly discuss the possibilities to observe these tracks. ",
        "introduction": "\\label{intro} Emerging evidence for the accelerated expansion of the Universe triggered interest in the non-standard theories of gravity in which gravitational interactions get modified in the infrared. To some extent the motivation is that in these theories, unlike in the case of pure cosmological constant, physics responsible for the cosmic acceleration may manifest itself in observations at smaller distance scales (e.g., in Lunar Ranging experiments \\cite{Lue:2002sw,Dvali:2002vf}).  Naively, the most straightforward way to modify gravity at distance scale $r_c$ would be to give a graviton a mass $m_g\\sim r_c^{-1}$. However, conventional massive gravity suffers either from the presence of ghosts or from the loss of predictivity because of strong coupling at  unacceptably low energy scales \\cite{Arkani-Hamed:2002sp}. Similar problems arise in  multi-dimensional models, and it is not clear whether there exists a consistent {\\it quantum} brane world theory where gravity is modified in the infrared and predictive power is not lost at unacceptably low energy scale.  Recently, an example of a theory which does not suffer from the above problems has been constructed \\cite{Arkani-Hamed:2003uy}. This theory, dubbed ``ghost condensate'', is somewhat similar to the Fierz--Pauli massive gravity. The difference is that the Fierz--Pauli mass term breaks reparametrization invariance completely, while in the ghost condensate theory only the time reparametrization invariance \\[ t\\to t+\\xi(t,x) \\] is broken, while the invariance under (possibly time-dependent) spatial diffeomorphisms is kept intact.  As a result,  the latter theory becomes formally reparametrization invariant after a single St\\\"uckelberg field is introduced, as opposed to four St\\\"uckelberg fields in the Fierz--Pauli gravity. The key difference between these two theories is that in the case of ghost condensate,  decoupling limit exists in which gravity is switched off while the St\\\"uckelberg sector is still described by a well defined low-energy effective theory valid up to a certain energy scale $M$.  The price  one pays is that the Lorentz invariance is not preserved by ghost condensate; for instance, at the quadratic level one effectively adds to the Einstein theory the ``mass term'' of the form \\be \\label{mass} \\int dtd^3x{1\\over 8}M^4h_{00}^2\\;. \\ee As a consequence of this violation of the Lorentz invariance, the dispersion law for the St\\\"uckelberg field\\footnote{In what follows, we somewhat loosely refer to this field as a ghost. It is worth stressing however, that this is not a ghost field in the usual sense, i.e. the sign in front of its kinetic term in the action is positive.} $\\pi$ has rather peculiar form, \\[ \\omega^2={\\alpha\\over M^2} k^4\\;, \\] where $\\alpha$ is a dimensionless parameter of the theory. Another manifestation of the fact that the Lorentz invariance is broken is that modification of gravity in the infrared is not characterized by a single length scale $r_c$. Instead, there is a length scale $r_c$, equal to \\be \\label{rc} r_c={\\sqrt{2}M_{Pl}\\over M^2} \\ee and a time scale $t_c$ given by \\be \\label{tc} t_c={2M_{Pl}^2\\over \\alpha M^3}\\;. \\ee As discussed in Ref.~\\cite{Arkani-Hamed:2003uy}, for a massive source {\\it at rest} the length scale $r_c$ determines the characteristic distance at which the gravitational potential starts to deviate from the Newtonian one, while $t_c$ determines the characteristic time needed for this deviation to show up. Naively, this implies that, assuming that $\\alpha\\sim 1$ and $t_c$  is of the order of the present age of the Universe, $t_U\\sim 15$ Gyr, one might expect a sizeable modification of gravity due to ghost condensate at length scales of order 1000 km.  However,  it was noted already in Ref.~\\cite{Arkani-Hamed:2003uy} that the retardation effects are very strong in the ghost sector. The point is that it is the whole history of a system that determines its actual gravitational potential in the presence of  ghost condensate. The purpose of this paper is to better understand this feature and thus reconsider possible observational signatures of ghost condensate. It is worth noting that for the moment, consequences of  ghost condensate with $t_c \\sim t_U$ for present day cosmology have not been elaborated yet (see, however, Ref.~\\cite{Arkani-Hamed:2003uz} where ghost condensate was used to construct a model of inflation with quite unusual perturbation spectra). Still, we believe that the question we address is of relevance, since ghost condensate is an interesting infrared modification of gravity whose consistency is beyond any doubt.  Surprisingly, we find that the above naive expectation is incorrect and it is not excluded that we  live in the Universe with $t_c\\sim t_U$ and $r_c\\sim 1000$~km and have not noticed that so far. To understand how that could be, one recalls a simple fact that objects in the Universe are not at rest. Instead, solar system and other stars in our Galaxy rotate around the center of the Galaxy with typical velocity of order $10^{-3}$, while the Galaxy itself moves in the local cluster of galaxies with the velocity of the same order of  magnitude. A well-known observational consequence of this motion is the dipole anisotropy of the cosmic microwave background.  This implies that all stellar objects have velocity of at least of the same order of magnitude with respect to the rest frame of  ghost condensate\\footnote{A priori one may think that this velocity can be significantly larger if the rest frame of the CMB has a finite velocity with respect to the rest frame of the ghost condensate. On the other hand we expect that one of the effects of the Hubble friction is to slow down this overall motion of the CMB. Throughout this paper we assume that the rest frame of the CMB coincides with the rest frame of the ghost condensate, leaving the study of whether it may be not the case for future.}. Now, it takes time of order $t_c$ for the modification of the gravitational potential to occur. Consequently, the effect of ghost condensate which we can observe on the Earth now is not a modification, say, of the gravitational field of the Sun, but the ``ghost'' tail of the potential of a star which was located nearby (in the rest frame of  ghost condensate) time $t_c$ ago. The Universe with ghost condensate can be thought of as a kind of a bubble chamber where all moving massive objects leave  long (and, as we will see later, narrow) tracks in which ghost field and gravitational potential are perturbed. The time delay between the moment when the object passes  a given point in space and the appearance of the track around this point is of order $t_c$.  This observation is our main result. The rest of this paper is organized as follows. In section \\ref{general} we derive  a general expression for the gravitational potential of a massive source in the presence of ghost condensate (in the Newtonian approximation). In section \\ref{specific} we consider specific examples, namely a potential of a moving point-like source and the effect of the finite size of the source. Section \\ref{evidence} contains preliminary discussion of phenomenological implications of our calculations. Technical details can be found in  Appendix. ",
        "conclusions": "\\label{evidence} Now we are at a point to discuss possible observational signatures of  ghost condensate, taking into account the effect of non-zero velocity of all stellar objects in the rest frame of ghost condensate. Let us start with a few preliminary remarks. From the phenomenological point of view, ghost condensate  (at the linearized level) has the characteristic time and length scales $t_c$ and $r_c$. These parameters are related to the two microscopic parameters,  mass scale $M$ and dimensionless parameter $\\alpha$, as written in Eqs.~(\\ref{rc}), (\\ref{tc}). The allowed deviation of the latter parameter from unity is determined by the amount of fine-tuning that one  is ready to introduce in the theory. We are going to be rather generous in this respect; in fact, our discussion is quite flexible and as large values of $\\alpha$ as $10^{10}$ will not affect it significantly.  An exhaustive phenomenological analysis of the theory would have resulted in the exclusion plot in the $(t_c,r_c)$-parameter space, and in the detailed discussion of the  characteristic observational signatures for different allowed regions. We believe that such an analysis deserves a separate publication. Our purpose here is to discuss qualitative features of the ghost condensate phenomenology, stressing the crucial role of the effect of finite velocity. Our  claims are the following. \\begin{enumerate} \\item It is very unlikely to observe ghost condensate with $t_c\\sim t_U$, where $t_U$ is the present age of the Universe. In other words, it is crucial for the observability of  ghost condensate that it enters the regime in which tracks with exponentially enhanced field are present. \\item The tracks of compact massive objects (stars) become pronounced earlier than the tracks of the supermassive objects of small density (galaxies). \\item The chance to observe ghost condensate is larger for larger values of $r_c$. \\item Relatively promising ways of searching for tracks in ghost condensate are: (i) search for ``mad'' stars which feel the gravitational field of the tracks of other stars; (ii) microlensing observations and (iii) gravitational wave experiments. \\item It may happen that $t_c$ is so small that tracks of some objects are already in the non-linear (quantum?) regime and, still, we have not noticed the presence of ghost condensate so far. Consequently, it may happen that to fully understand the phenomenology of ghost condensate one needs the details of the UV completion in the ghost sector. \\end{enumerate}  Let us first explain why it is unlikely to observe  ghost condensate for large characteristic time scales, $t_c\\gtrsim t_U$. In this regime, tracks with exponentially enhanced field did not have enough time to develop, so the only potentially observable effects are due to the gravitational potential in the wave zone. Let us first assume that the characteristic size $r_c$ is somewhat larger than the size of a typical star like the Sun, \\[ r_c\\gtrsim 10^6\\;\\mbox{km}\\; \\] This actually implies that the parameter $\\alpha$ is quite large, $\\alpha\\gtrsim 10^5$; however, as we will see, the chance to detect ghost condensate is even lower for smaller values of $r_c$. Then, to estimate the ``extra'' gravitational potential  $\\Delta\\Phi$ (see, Eq.~(\\ref{NDel})) of a star whose trajectory was at a distance $y_0$ to the current location of the Earth (in the rest frame of ghost condensate) we make use of the expression (\\ref{weaklimit}) for the potential of a point-like source in the wave zone, \\be \\label{weakstar} \\Delta\\Phi\\sim 10^{-20} \\l \\mu\\over M_{\\odot}\\r \\l{ 10^{-3}\\over v}\\r \\l{r_c\\over    y_0}\\r^2 \\sin{y_0\\Delta y\\over 2r_c^2} \\ee where we set $T=1$, expanded the phase of oscillations in Eq.~(\\ref{weaklimit}) near a given space-time point setting \\[ y=y_0+\\Delta y \\] neglected  extremely slow variation of this phase in time and in $z$-direction and  dropped  the constant shift of this phase. Equation (\\ref{weakstar}) applies to the rest frame of ghost condensate, while in the rest frame of the Earth the gravitational field has the form of a wave of the amplitude $\\sim 10^{-20}$  and frequency \\be \\label{nu} \\nu= {v_r y_0\\over 4\\pi r_c^2}\\simeq 2\\cdot 10^{-5}\\; \\mbox{Hz}\\l {v_r\\over 10^{-3}}\\r \\l{y_0\\over r_c}\\r\\l{10^6\\;\\mbox{km}\\over r_c}\\r \\ee where $v_r$ is the Earth velocity in the direction transverse to the star trajectory. Note that this is a scalar wave unlike the tensor waves of the Einstein theory. The gravity waves of such a low frequency are in principle accessible to the LISA project (see, e.g., Refs.~\\cite{Thorne:1995xs,Lobo:2002pr} for  reviews of the gravitational wave experiments), however, its sensitivity at these frequences is at the level\\footnote{It is worth noting that gravitational wave experiments are sensitive not to the amplitude $\\Delta\\Phi$ of the gravitational wave itself, but to the product $\\sqrt{n}\\Delta\\Phi$, where $n$ is a number of cycles produced in a logarithmic band about a given frequency. In the case at hand $n\\sim\\nu y_0/v_r=y_0^2/(4\\pi r_c^2)$. This comment is practically irrelevant for our discussion.} $\\Delta\\Phi\\sim 10^{-17}$.  Furthemore, it is straightforward to see that the probability $p(r_c)$ to have even this weak signal is extremely low. Indeed, to estimate this probability, note first that the probability ${\\cal P}(r_c)$ that  the distance from the current position of the Earth to the nearest trajectory of a star is smaller than $r_c$, is given by \\be \\label{probest} {\\cal P}(r_c)\\sim {N_{st}N_gvt_Ur_c^2\\over r_U^3}\\sim 10^{-15} \\ee where $N_{st}\\sim N_g\\sim 10^{11}$ are, respectively, the number of stars in a typical galaxy and the number of galaxies in the Hubble volume; $r_U\\sim 10^{28}$ cm, $v\\sim 10^{-3} $ and $r_c\\sim 10^6$~km. For long enough period of observation $t_o$, such that the Earth travels a distance $L_E$ much larger than $r_c$, the probability $p(r_c)$ is somewhat larger, \\[ p(r_c)\\sim {\\cal P}(r_c){L_E\\over r_c}\\sim 10^4\\cdot {\\cal P}(r_c)\\l {t_o \\over 1\\;\\mbox{yr} }\\r\\l{v\\over 10^{-3}}\\r\\l{10^6\\;\\mbox{km}\\over r_c}\\r \\] Since the amplitude in Eq.~(\\ref{weakstar}) rapidly decreases at $y_0 > r_c$, this expression determines the probability of having a non-negligible signal. Clearly, this probability is quite low.   The sensitivity of  gravitational wave detectors is higher in the higher frequency bands, so one might expect better signal for smaller $r_c$, and hence higher frequency $\\nu$. For instance, LISA will have  sensitivity at the level $\\Delta\\Phi\\sim 10^{-20}$ in the frequency range $10^{-3}\\div 10^{-1}$~Hz. The waves of amplitudes and frequencies in this range would be generated for $y\\sim r_c\\sim 1000$~km. In this case, in order to avoid the suppression of the amplitude due to the effect of the finite size of the source (section \\ref{thicksub}), one should consider very compact sources like neutron stars. But the above estimate shows that we should be extremely lucky to have a trajectory of a neutron star at a distance of order 1000 km from the LISA facility.  Let us now discuss the effects  from  objects of larger size, e.g. galaxies. One could expect two different types of signatures on these large distance scales. The first is the gravitational wave signals like those discussed above. The second is the modification of the gravitational dynamics at large scales due to extra contributions to the Newtonian potential. However, it is straightforward to see that the effect of averaging discussed in section \\ref{thicksub} kills all these signatures unless the value of the characteristic length scale $r_c$ (and, correspondingly, of the parameter $\\alpha$) is extremely large. Indeed, due to the large mass of a galaxy there is an extra factor of order $10^{12}$, as compared to a star, in the estimate (\\ref{weakstar}) for the gravitational potential. However, Eq.~(\\ref{wavebound}) tells us that there is an extra suppression by at least a factor \\be \\label{halofield} {r_c^2\\over L_g^2}\\sim 10^{-30}\\l{r_c\\over 1000\\;\\mbox{km}}\\r^2\\l{30\\;\\mbox{kpc} \\over L_g}\\r^2 \\ee where $L_g$ is the size of a galaxy. We see that the effects of ghost condensate are indeed very small for large objects of small density.  The above arguments show that  chance to detect ghost condensate is very low if the characteristic time scale $t_c$ is longer than, or equal to the present age of  the Universe. This forces us to discuss the regime $t_c\\lesssim t_U$. This regime is rather dangerous, as the gravitational and ghost fields grow exponentially  inside the tracks. Still, let us make a few general remarks,  postponing the detailed discussion of the ghost condensate phenomenology in this regime for future.  If the ratio $t_U/t_c$ is large enough, the tracks of  galactic halos are very pronounced (say, the gravitational potential in the track is comparable to the typical gravitational potential between two interacting galaxies). The fraction of the Universe filled by these tracks is  estimated as \\be \\label{halofraction} {\\Delta V_h\\over V_U}\\sim N_g{r^2_{halo}vt_u\\over t_u^3}\\sim 0.05 \\l{N_g\\over 10^{11}}\\r \\l{r_{halo}\\over 100\\mbox{ kpc}}\\r^2\\l{v\\over 10^{-3}}\\r \\ee where we estimated the size of a halo of a typical galaxy as 100 kpc. It is unlikely that this effect would have been  unnoticed. For instance, the dynamics of a sizeable number of galaxies would have been affected by these tracks.  For not so large $t_U/t_c$ the situation is more interesting. There is a range of parameters in which tracks of galaxies are unnoticeable, but star tracks are strong. As an example, one finds from Eqs.~(\\ref{weakstar}) and (\\ref{halofield}) (assuming $r_c=1000$~km for definiteness) that for \\[ {t_U\\over t_c}\\sim 2 \\ln{10^{20}}\\sim 92 \\] the gravitational potentials in the tracks of  neutron stars are of order one, while the gravitational potentials in the tracks of  typical galactic halos are still of order $10^{-18}$ which is more than 10 orders of magnitude smaller than the gravitational potential of a typical galaxy at a distance of order 1 Mpc. Consequently, there is an interesting range of the characteristic time scales, \\be \\label{interesting} 0.01 \\lesssim {t_c\\over t_U} <1 \\ee in which star tracks, but not galactic tracks, are pronounced (up to $\\Delta\\Phi\\sim 1$).  Normally, the existence of exponentially growing field would mean that an extreme fine-tuning of the parameters is needed to have potentially observable effects without ruling out the theory completely. Clearly, to have $t_c$ in the range (\\ref{interesting}) one needs some fine-tuning, but not that strong as one might expect.  We suggest here three potential signatures of the star tracks. First, one can search for  ``mad'' stars, intersecting the tracks of other stars, so that their motion is strongly affected by the gravitational field of the track (probably, just for a short period of time). Second, it seems possible to observe tracks in the microlensing experiments (see, e.g., Ref.~\\cite{Roulet:1996ur} for a review of microlensing experiments) due to the variation of the visible luminosity of a background star when the line of sight  crosses the track. Finally, the gravitational wave detetors may detect a signal from the track, as discussed above.  To observe  mad stars one needs a galaxy whose disc is intersecting a track of the disc of another galaxy at the moment of observation. Plugging the typical disc size $r_{disc}\\sim 10$~kpc instead of $r_{halo}$ into Eq.~(\\ref{halofraction}) we find that this happens for approximately one galaxy of a thousand. To get a feeling of numbers let us also estimate the fraction of the volume of such a galaxy filled by the star tracks, \\be \\label{starfraction} {\\Delta V_{s}\\over V_{galaxy}}\\sim N_s\\l {r_s\\over r_{disc}}\\r^2\\sim 10^{-12}\\;, \\ee where $r_s\\sim 10^6$~km is the radius of a typical star. This estimate is not very optimistic as it implies that there is about one mad star in the galaxy at each moment of time. Note however, that one can significantly enhance the success probability by performing  monitoring of stars for a long period of time. Also it would be interesting to check whether trapping of  a star by the gravitational field of a track is possible. Finally, the estimate (\\ref{starfraction})  assumes that the parameter $r_c$ is smaller than $r_s$ so that the diameter of the star track is equal to the size of the star. At larger values of $r_c$ this fraction is enhanced by a factor of $(r_c/ r_s)^2$.   Similar considerations show that chance to observe two other signatures (microlensing and gravity waves) are rather low  for $r_c<r_s$.  The above estimates show that it may happen, that while we have not noticed the presence of ghost condensate yet, the gravitational and ghost fields in the tracks of some dense objects are already very strong. Note that it follows from Eq.~(\\ref{action}) that gravitational field $\\Delta\\Phi\\sim 1$ corresponds to  ghost field $\\pi_c\\sim M$, so  it is unclear whether  non-linear classical dynamics will result in some stable non-linear classical track solution or  strong quantum effects will start operating at this point. [In particular, there is a danger that the points of the field space where excitations of  ghost condensate are ghosts (i.e., have wrong sign in front of kinetic term) become accessible inside the tracks.] It seems difficult to deduce what is happening in this regime without knowing the UV completed theory of  ghost condensate (in other words, without constructing a complete analog of the Higgs mechanism for gravity and not just a non-linear sigma-model).   To conclude, let us mention some other open questions and further directions of research.  First, in our treatment we did not take into account the expansion of the Universe. In the interesting case (\\ref{interesting}) this assumption seems justified because at the initial, rather short cosmological stage when the Hubble rate was higher than the rate of the development of instabilities in  ghost condensate $t_c^{-1}$, the Hubble friction prevented the instabilities to grow. Also, the Universe was essentially homogeneous during the fast period of expansion, while the tracks discussed here emerge due to the inhomogeneities. Still, we believe that this question deserves further study. Especially interesting would be to check the possibility to have sizeable tracks of inhomogeneities which could be present before inflationary stage.  Second, our discussion was purely non-relativistic.  However, one could argue that there is a possibility that the whole observed part of the Universe is moving with respect to the rest frame of ghost condensate with  velocity close to the speed of light. It would be interesting to check whether this is a consistent possibility, or one of the effects of the Hubble friction is to slow down the velocity with respect to the rest frame of  ghost condensate, so the non-relativistic approximation is always justified. In the latter case it would be natural to assume that the rest frame of  ghost condensate coincides with the rest frame of the CMB. If true, this assumption would imply that the motion with respect to  ghost condensate is entirely determined by the peculiar velocities of galaxies and galactic clusters. Then it would be interesting to reconstruct the actual map of galactic tracks at least in the local part of the Universe. Such a map would significantly decrease the uncertainties in the observational predictions of the model.  It is worth  noting also, that the  effect of finite velocity of an observer with respect to the rest frame of ghost condensate should also be taken into account while discussing the limits on the direct coupling between matter fields and ghost condensate,  e.g., coming from the spin-dependent forces~\\cite{Arkani-Hamed:2003uy} mediated by ghost condensate.  Finally, ghost condensate may be considered as a particular case of the $k$-essence models~\\cite{Armendariz-Picon:2000dh,Armendariz-Picon:2000ah}. It is interesting to check whether the effects similar to those discussed here can be present and pronounced in the models of $k$-essence with cosmologically relevant parameters."
    },
    "0403/astro-ph0403659.txt": {
        "abstract": "%[x]  We fit the surface-brightness profiles of 21 elliptical galaxies using both the S\\'ersic function and a new empirical model which combines an inner power law with an outer S\\'ersic function.  The profiles are combinations of deconvolved \\textit{HST} profiles from the literature and ellipse fits to the full WFPC2 mosaic images, and thus span a radial range from $\\sim 0\\farcs02$ to $\\sim$ twice the half-light radius.  We are able to accurately fit the entire profiles using either the S\\'ersic function or our new model.  In doing so, we demonstrate that most, if not all, so-called ``power-law'' galaxies are better described as ``S\\'ersic galaxies'' --- they are well modeled by the three-parameter S\\'ersic profile into the limits of \\textit{HST} resolution --- and that ``core'' galaxies are best understood as consisting of an outer S\\'ersic profile with an inner power-law cusp, which is a downward deviation from the inward extrapolation of the S\\'ersic profile.  This definition of cores resolves ambiguities that result when the popular ``Nuker law'' is fitted to the profiles of ellipticals and bulges, particularly at lower luminosities.  We also find that using the Nuker law to model core-galaxy nuclear profiles systematically overestimates the core radii by factors of 1.5--4.5 and underestimates the inner power-law slope by $\\sim 20$--40\\% or more. ",
        "introduction": "%[x] \\label{sec:intro}  The availability of high-resolution imaging with the \\textit{Hubble Space Telescope} (\\textit{HST}) has revolutionized the study of galaxy centers.  Following up on early work by \\citet{crane93}, \\citet{kormendy94}, \\citet{grillmair94}, \\citet{jaffe94}, and \\citet{ferrarese94}, a series of papers by the ``Nuker team'' \\citep{lauer95,byun96,gebhardt96,faber97} presented a detailed study of the central regions of early-type galaxies (specifically, ellipticals and the bulges of spiral galaxies).  They introduced a model for fitting the radial surface-brightness profiles: a double power-law with an adjustable transition region, dubbed the ``Nuker law'': \\begin{equation} I(r) = I_b \\, 2^{(\\beta - \\gamma)/\\alpha} \\left(\\frac{r}{r_b}\\right)^{-\\gamma} \\left[1 + \\left( \\frac{r}{r_b} \\right)^{\\alpha}\\right]^{(\\gamma - \\beta)/\\alpha}. \\end{equation} The inner and outer power law exponents are $\\gamma$ and $\\beta$, respectively; $I_{b}$ is the surface brightness at the core or ``break'' radius $r_b$, and $\\alpha$ controls the sharpness of the transition between the two power laws (larger $\\alpha$ = sharper transition).  They identified two distinct classes of galaxy centers: ``power-law'' galaxies, where the central surface brightness increases into the limit of resolution with something like a steep power-law profile; and ``core'' galaxies, where the luminosity profile turns over at a fairly sharp ``break radius'' into a shallower power-law. Ferrarese et al.\\ and Faber et al.\\ found evidence that global parameters of early-type galaxies correlated with their nuclear profiles: core galaxies tend to have high luminosities, boxy isophotes, and pressure-supported kinematics, while power-law galaxies are typically lower-luminosity and often have disky isophotes and rotationally supported kinematics.  The Nuker-law parameterization of galaxy centers has subsequently enjoyed a great deal of popularity, including extensive studies using WFPC2 and NICMOS \\citep[e.g.,][]{rest01,quillen00,ravindranath01,laine03}, and extensions to early- and late-type spirals \\citep[e.g.,][]{carollo98,seigar02}.  These more recent studies have, however, suggested that the clear core/power-law dichotomy found by the Nuker team may not be so clear after all.  In addition, almost all the studies using \\textit{HST} data and Nuker-law fits have left unanswered a key question: how does the nuclear part of a bulge or elliptical, seemingly well fit by a double power-law, connect to the outer profiles of such systems, which are generally well fit by the \\citet{sersic} $r^{1/n}$ function?  In our first paper \\nocite{paper1}(Graham et al.\\ 2003a, hereafter Paper~I), we discussed some of the systematic problems and ambiguities which can arise when using a double power-law model to fit galaxy light profiles, and suggested a new hypothesis and a new model which might resolve some of these problems.  The hypothesis has two parts: first, that the nuclear (\\textit{HST}-resolved) profiles of most lower-luminosity hot systems, including the power-law galaxies, are simply inward extensions of each galaxy's outer profile, best modeled with a S\\'ersic function; second, that core galaxies are best modeled with our new function, an outer S\\'ersic function with a break to an inner power-law.  In this paper, we make an empirical test of this proposed solution, by modeling the \\textit{entire} light profiles of a sample of elliptical galaxies.  In what follows, we first review some of the problems stemming from the use of the Nuker law, including the problem of how best to identify genuine cores in galaxies (Section~\\ref{sec:nuker-problems}); readers familiar with these issues can probably skip this section. We then discuss our sample selection, data reduction and analysis, and the source of the profiles used (Section~\\ref{sec:data}).  In Section~\\ref{sec:models}, we discuss the S\\'ersic model and our new model for core-galaxy profiles.  Section~\\ref{sec:fits} presents criteria for identifying core galaxies, and for discriminating between core and S\\'ersic profiles.  We also present the results of our fits to the galaxy profiles and compare their fidelity to the profiles with that of the Nuker-law fits.  Some of the implications are discussed in Section~\\ref{sec:discuss}, and we conclude with a brief summary in Section~\\ref{sec:summary}.  Finally, several useful mathematical expressions related to our new model are presented in the Appendix. ",
        "conclusions": "\\label{sec:discuss}  We conclude that most, if not all, ``power-law'' ellipticals are probably best understood as having S\\'ersic profiles --- modulo extra components such as nuclear star clusters, nuclear disks, etc.\\ --- into the limits of resolution (or limits imposed by dust).  As discussed in Section~2, this is consistent with an overall trend for elliptical galaxies: low- and intermediate-luminosity ellipticals have pure S\\'ersic profiles (plus optional nuclear disks, clusters, and point sources), and distinct cores appear in high-luminosity systems as deviations from the outer S\\'ersic profile. \\nocite{graham-guzman03}(Graham \\& Guzm\\'an 2003 combine measurements for a large set of elliptical galaxies, including dwarf ellipticals, to make this argument in more detail.)  Moreover, for power-law galaxies, we get excellent fits using a model with fewer parameters, all of which are physically meaningful (i.e., correlate with other galaxy parameters).  These fits work for the \\textit{entire} profile, \\textit{unlike} the Nuker law, yet are as good a fit in the region where the Nuker law is usually used.  The term ``power-law galaxy'' is thus somewhat misleading, since it suggests that the nuclear profile is adequately described by a single power-law, which is probably different from the outer profile.  While this is an appealingly simple description for modeling purposes, our results strongly suggest that this is not accurate.  Instead, elliptical galaxy profiles have logarithmic slopes which continuously decrease as $r \\rightarrow 0$.  Figure~11 of \\citet{lauer95}, which presents representative examples of ``power-law'' profiles, supports this argument: even the galaxy which is closest to a perfect power-law, NGC~1700, shows a systematic deviation from a power-law --- steeper at larger radii, shallower at smaller radii --- as expected for a S\\'ersic profile; see Figure~\\ref{fig:sersic-fits}.  (This is not the case for the central cusps of \\textit{core} galaxies; their Figure~7.)  The ``intermediate'' galaxies reported by \\citet{rest01} and \\citet{ravindranath01} are probably a consequence of Nuker-law fits applied to this overall elliptical-galaxy trend.  Lower-luminosity ``intermediate'' galaxies are most likely S\\'ersic galaxies with low values of $n$ (and hence $<\\gamma>$ in the range 0.3--0.5; see Fig.~\\ref{fig:de-gamma}).  At higher luminosities, core galaxies can appear to have $\\gamma > 0.3$ if the core is not adequately resolved (either due to distance or to inner truncation of the profile by, e.g., dust).  (We \\textit{do} classify two galaxies in our sample as ``possible core'' galaxies, but these are clearly cases of inadequate resolution.)  Although we have not yet attempted to model the complete profiles of \\textit{bulges}, it is reasonable to extend our results to them. \\citet{balcells03} have already done this for a sample of early-type bulges in the near-IR, using NICMOS data in conjunction with ground-based imaging.  They find that the complete bulge profiles, after accounting for the presence of the outer disk, can be well modeled by S\\'ersic profiles, plus optional nuclear components (corresponding to, e.g., nuclear star clusters or point sources). This is in excellent agreement with our hypothesis that the profiles of lower-luminosity ellipticals and bulges are fundamentally S\\'ersic profiles, and promises to resolve a number of ambiguities and ``dichotomies'' reported in the literature.  For example, \\citet{carollo97} and \\citet{seigar02} argue for a dichotomy between $R^{1/4}$ and exponential bulges, with the latter having low $\\gamma$ in contrast to the high $\\gamma$ of $R^{1/4}$ bulges and moderate-luminosity ellipticals.  This is naturally explained if most bulges actually have S\\'ersic profiles (as is well supported by a number of studies) \\textit{and} if these S\\'ersic profiles extend into the nuclear region.  The division between $R^{1/4}$ (S\\'ersic index $n = 4$) and exponential ($n = 1$) bulges is probably an artificial one, given that bulges in reality show a range of values of $n$.  But as \\nocite{paper1}Paper~I shows, bulges with larger $n$ will have higher values of $\\gamma$ than bulges with low $n$.  Thus, ``$R^{1/4}$'' bulges (higher $n$) will exhibit larger values of $\\gamma$ than ``exponential'' (lower $n$) bulges.  Since bulge $n$ decreases along the Hubble sequence, the trend of decreasing $\\gamma$ with Hubble type noted by \\citet[][ their Fig.~3]{seigar02} follows as well.  In retrospect, we can see that most of the early \\textit{HST} studies of galaxy centers, and some of the more recent ones \\citep[e.g.,][]{rest01,ravindranath01}, have focused on relatively high-luminosity systems.  These samples thus included a mix of S\\'ersic galaxies with high $n$ values and genuine core galaxies, making a distinction between core and ``power-law'' galaxies based purely on $\\gamma$ feasible.  More recent studies aimed at low-luminosity systems \\citep[e.g.,][]{carollo98,stiavelli01,seigar02} have since uncovered evidence for the low-$n$--low-$\\gamma$, high-$n$--high-$\\gamma$ trend that pure S\\'ersic profiles generate, and thus show that discriminating core galaxies purely by $\\gamma$ is problematic at best.     \\subsection{Core Identifications and Core Parameters} %[x] \\label{sec:core-params}  We find that most of the previously identified ``core'' galaxies in our sample \\textit{do} have distinct cores with shallow, power-law cusps.  These cores stand out as downward deviations from the outer S\\'ersic profiles.  Fitting with the \\bomba{} model provides a more natural, less ambiguous definition for ``true'' cores, without the possibility of misclassifying low-$n$ S\\'ersic profiles as cores.  We are also able to re-classify one of the ``intermediate'' galaxies (NGC~5557) of \\citet{rest01} as a core galaxy.  The two ambiguous galaxies --- NGC~3613 and NGC~5077 --- are simply cases where the apparent break radius is very close to the inner limits of the data. For NGC~3613, this is because the apparent core is close to the resolution limit (in fact, $r_{b}$ from the \\bomba{} fits is $< 0\\farcs16$ and thus smaller than the suggested resolution-based limit of \\nocite{faber97}Faber et al.\\ 1997).  For NGC~5077, on the other hand, \\citet{rest01} clipped their data at $r = 0\\farcs1$ because of an apparent nuclear excess at smaller radii.  A future fit including data at smaller radii and using an extra nuclear component to account for this excess, may help determine if NGC~5077 truly possesses a core.  While our overall agreement with the core/non-core classifications of \\citet{lauer95} and \\citet{rest01} is quite good for the galaxies we analyze, \\textit{we find that Nuker-law fits systematically overestimate the size of the cores}: our break radii are $\\sim 1.5$--4.5 times smaller in size than the break radii from the published Nuker-law fits.  Consequently, $\\mu_{b}$ values are brighter as well.  We also find consistently higher values of $\\gamma$, though the difference is not as dramatic (see Table~\\ref{tab:core-params} and Figure~\\ref{fig:core-core}).  This is in excellent agreement with the arguments of \\nocite{paper1,paper3}Papers~I and III: \\textit{all} Nuker-law parameters are sensitive to the radial size of the region where the fit is made.  All parameters of the Nuker model, including $\\gamma$ and $r_{b}$, must be adjusted in order to fit both the core \\textit{and} the (non--power-law) part of the profile outside, with its intrinsic (S\\'ersic) curvature.  Table~\\ref{tab:core-params} shows that, on average, the \\bomba{} values of $\\gamma$ match the \\textit{observed} core slope $\\gammap$ (as determined by \\citet{rest01}) better than the Nuker-law values do.  The currently favored theory for core formation is the ejection of core stars by 3-body encounters with a decaying black hole binary formed following a merger of two galaxies with central supermassive black holes.  Various calculations \\citep{ebisuzaki91,quinlan97,milosavljevic01} have estimated the stellar mass ejected during this process (\\mej), and generally find it to be $\\sim \\mbh$, where \\mbh{} is the mass of the resulting central black hole formed by the (assumed) coalescence of the binary. However, attempts to test these predictions by estimating \\mej{} from observed cores and comparing it with various estimates of \\mbh{} consistently produce values of $\\mej > \\mbh$.  \\citet{faber97} found $\\mej = 3.5$--6.4 \\mbh; using more accurate estimates of \\mbh, \\citet{milosavljevic01} found $\\mej \\approx 1$--20 \\mbh. \\citet{ravindranath02} used the prescription for \\mej{} of Milosavljevi\\'{c} \\& Merritt and a much larger data set; they found $\\mej \\approx 2$--20 \\mbh{} at the low-mass end ($\\mbh \\sim 10^{8} M_{\\sun}$), while at the high-mass end ($\\mbh \\sim 10^{9} M_{\\sun}$) $\\mej \\approx 6$--25 \\mbh.  Even considering only the galaxies with \\textit{measured} \\mbh, $\\mej/\\mbh \\approx$ 4--13.  Milosavljevi\\'{c} \\& Merritt pointed out that the \\textit{total} ejected mass should increase with the number of mergers, but the observed ratios still seem high, particularly at the low-mass end, where there have presumably been fewer mergers.  All of the studies cited above used parameters from Nuker-law fits to estimate \\mej.  Since the estimated \\mej{} scales with $r_{b}$ --- in the parameterization introduced by \\citet{milosavljevic01} and used by Ravindranath et al.\\ (2002), $\\mej \\propto r_{b}$ --- overestimating $r_{b}$ will naturally overestimate \\mej.  Thus at least some of the discrepancy between observed and predicted $\\mej/\\mbh$ is probably due to the tendency of Nuker-law fits to overestimate $r_{b}$, as we have found.  Assuming that the core radii from \\bomba{} fits are typically $\\sim 2$--4 times smaller than the Nuker-law values, as is the case for our sample, $\\mej/\\mbh$ values should go down by comparable factors, which would put them in better agreement with the theoretical predictions.   One of our core galaxies (NGC~4291) was noted by \\citet{ravindranath01} for possibly having an isothermal core (with $\\gamma = 0$), on the basis of their Nuker-law fits to a NICMOS image. The Nuker-law fit in \\citet{rest01} to the WFPC2 profile also has $\\gamma = 0.0$, which might seem to strengthen the case for an isothermal core.  However, we find $\\gamma = 0.14$ from our \\bomba{} fit, which agrees very well with $\\gammap = 0.13$ determined by Rest et al.  So the core of NGC~4291 is probably \\textit{not} isothermal.  In Figure~\\ref{fig:core-global} we show the relation between the core and the global properties of the galaxies in our sample.  We also indicate the upper limits on possible core radii for the S\\'ersic galaxies, based on the radii of the innermost valid data.  For those galaxies where a clear core has been measured, we find that the relation between the break radius and the effective radius is approximately given by $r_b = 0.014 r_e$.  This is a factor of two smaller than the relation found by \\citet{faber97}, consistent with our finding that fitting with the Nuker law tends to overestimate core sizes.  There is a suggestion of a weak trend of $r_{b}$ increasing with galaxy luminosity, which would be in agreement with what Faber et al.\\ found \\citep[see also][]{laine03}, but for our sample this ``trend'' is anchored by only two points, so it is dubious.  Unfortunately, the narrow magnitude range spanned by the core galaxies in our sample ($\\lesssim 1.5$ mag) precludes a proper test of the magnitude-$r_{b}$ relation reported Faber et al., which is based on galaxies spanning $\\gtrsim 3$ mag (and the composite trend in Fig.~9 of Laine et al.\\ spans almost 5 magnitudes).  There is no clear magnitude-related trend in the \\textit{ratio} of our $r_{b}$ measurements to the Nuker-law measurements, which suggests that the magnitude-$r_{b}$ trend may be unaffected by changes in $r_{b}$, except possibly in the scatter. However, a proper evaluation of how the magnitude-$r_{b}$ relation is affected by better measurements of $r_{b}$ must await \\bomba{} fits to a larger sample of core galaxies.  There is \\textit{no} evidence for a relationship between $n$ and $r_{b}$; this may be partly due to large uncertainties in $n$ \\nocite{caon93}(Caon et al.\\ 1993 found typical errors of $\\sim 25$\\% when fitting S\\'ersic profiles).  Finally, we find \\textit{no} clear correlation between $\\gamma$ and the global properties of the core galaxies analyzed.  This is agreement with what previous studies have found for core galaxies (e.g., \\nocite{rest01}Rest et al.\\ 2001, Figure~7; \\nocite{ravindranath01}Ravindranath et al.\\ 2001, Figure~3; \\nocite{laine03}Laine et al.\\ 2003, Figure~6; and the core galaxes in Figure~\\ref{fig:gamma-core} of this paper).     \\subsection{Hidden Cores and the Core-Galaxy Fraction} %[x]  An interesting point is to consider how well-resolved the underlying profiles of the various galaxies actually are.  In several cases, \\citet{byun96} and \\citet{rest01} excluded points at small radii from their fits, usually due to the presence of significant nuclear dust or a distinct nuclear component (e.g., a nuclear point source).  Thus, not all of the profiles take full advantage of \\textit{HST} resolution.  While the nuclear components may include cases of nuclear star clusters, which make discussions of the underlying stellar profile ambiguous, the presence of dust means that some ``power-law'' (i.e., S\\'ersic-profile) galaxies could have hidden cores.  If we divide the sample into two groups --- galaxies where the innermost valid data point is at $r < 15$ pc (spatially well resolved centers); and galaxies where the innermost valid point is at $r > 15$ pc (less well-resolved centers) --- we find that the \\textit{less} resolved galaxies are almost all\\footnote{The exceptions are NGC~4168 (core) and NGC~5077 (possible core).} well fit using just the S\\'ersic model.  This suggests that at least some of the S\\'ersic galaxies could have ``hidden'' cores.  This is not a new argument, obviously, as many authors have pointed out that ``power-law'' galaxies could include unresolved cores --- but it is interesting to consider how \\textit{few} of the S\\'ersic galaxies in our sample can really be declared free of \\textit{HST}-resolvable cores.  Of the 21 galaxies, seven clearly have cores, two have possible cores (NGC~3613 and NGC~5077, see Section~\\ref{sec:core-params}), and only five (NGC~4478, NGC~5576, NGC~5796, NGC~5831, and NGC~5845) are clearly free of significant ($r_{b} > 5$ pc) cores.  So in the limited range of absolute magnitude spanned by our full sample ($-18.3 \\gtrsim M_{B} \\gtrsim -21.4$), 33\\% of the galaxies have unambiguous, \\textit{HST}-resolved cores; but this is clearly a lower limit.  The core fraction rises to 43\\% if we include the two possible cases, and in principle could be as high as 76\\%.  It is also interesting to note that we can see in the absolute magnitudes a hint of the well-known dichotomy between core and non-core galaxies \\citep[see, e.g., the discussion in ][]{rest01}, even in our limited sample.  This can be seen in Figure~\\ref{fig:core-global}, where the five \\textit{fully resolved} S\\'ersic galaxies tend to be fainter than the core galaxies; a Kolmogorov-Smirnov test gives a 95\\% probability that the two groups of galaxies come from different parent luminosity distributions.       We have successfully fit the complete surface-brightness profiles of 19 out of 21 elliptical galaxies, from the \\textit{HST}-resolved central regions ($r \\sim 0\\farcs02$) out to $\\sim$ twice the half--light radius, using either: a) a pure S\\'ersic profile; or b) a ``\\bomba{}'' model consisting of an outer S\\'ersic profile joined to an inner power-law core.  The former fits correspond to so-called ``power-law'' galaxies, which are perhaps better described as ``S\\'ersic galaxies,'' and the latter correspond to core galaxies.   The combined use of these two models lets us address the following questions: \\begin{enumerate}  \\item \\textit{How can we relate the central, HST-resolved part of the galaxies' surface-brightness profiles to the outer regions?} We show that most power-law ellipticals are well described at all radii by the simple S\\'ersic law (modulo any nuclear disks, etc.).  On the other hand, core galaxies are extremely well fit with the \\bomba{} model.  We find little need for a significant transition region between the outer (S\\'ersic) part of the \\bomba{} profile and the (power-law) core; any such transition region is small compared to the size of the core.  \\item \\textit{Is there a dichotomy in nuclear profiles between low- and high-luminosity bulges and ellipticals?} Some recent \\textit{HST} studies have suggested that the apparent trend seen in intermediate- and high-luminosity bulges and ellipticals --- cores with shallow logarithmic slopes in high-luminosity systems, steeper nuclear slopes in lower-luminosity (``power-law'') systems --- breaks down at lower luminosities, because fainter bulges and dwarf ellipticals have shallow nuclear slopes.  We show that the power-law galaxies in our sample have S\\'ersic profiles that extend into the limits of \\textit{HST} resolution, with $n \\sim 4$--6; this naturally explains the steep nuclear slopes previously reported.  When combined with the well-known correlation between $n$ and luminosity, we can see that \\citep[as argued by][]{graham-guzman03} the general trend is most likely one of pure S\\'ersic profiles (plus possible extra components such as nuclear star clusters and disks), extending from low-luminosity systems with low-$n$ S\\'ersic profiles --- and thus shallow nuclear slopes --- to high-luminosity systems with high-$n$ profiles and steeper nuclear slopes.  Only the high-luminosity \\textit{core} galaxies break the trend, due to the existence of the cores themselves.  \\item \\textit{How can we unambiguously identify cores in galaxy profiles?} As we demonstrate, the traditional definition of cores using parameters from Nuker-law fits to galaxy profiles ($r_b \\geq 0\\farcs16$ and $\\gamma < 0.3$) leads to the real possibility of misclassifying galaxies with sufficiently shallow slopes (for example, exponential profiles) as core galaxies.  \\textit{We define core galaxies as those possessing a well-resolved downward deviation from the inward extrapolation of the outer (S\\'ersic) profile}.  This definition recovers previous core definitions for the high-luminosity ellipticals in our sample, but is immune to the danger of identifying exponential-like profiles as having cores.  \\item \\textit{How can we more accurately determine the structural properties of cores?} As demonstrated in Paper~I, the Nuker law requires a broad, smooth transition (low values of $\\alpha$) between its two power-law regimes in order to fit the inner profiles of core and power-law galaxies, because this is the only way to reproduce the observed curvature of actual galaxy profiles.  We find that this causes the core-size measurements (i.e., the break radius) to be overestimated by factors of 1.5--4.5 in comparison to the values derived by using the \\bomba{} model, which directly accounts for the intrinsic curvature of galaxy profiles.  We also find that the logarithmic slope $\\gamma$ of the observed core is more accurately recovered with the \\bomba{} model.  Using the smaller values we find, especially for $r_{b}$, should bring estimates of the ejected stellar mass due to core formation more in line with theoretical predictions.  \\end{enumerate}"
    },
    "0403/astro-ph0403251_arXiv.txt": {
        "abstract": "A soft component of thermal emission is very commonly observed from the surfaces of quiescent, accreting neutron stars.  We searched with {\\it Chandra} for such a surface component of emission from the dynamical black-hole candidate XTE~J1118+480 (=\\ J1118), which has a primary mass $M_{1} \\approx$ 8 \\msun.  None was found, as one would expect if the compact X-ray source is a bona fide black hole that possesses an event horizon.  The spectrum of J1118 is well-fitted by a simple power-law model that implies an unabsorbed luminosity of $L_{\\rm x} \\approx 3.5 \\times 10^{30}$ \\lum~(0.3--7 keV). In our search for a thermal component, we fitted our {\\it Chandra} data to a power-law model (with slope and $N_{\\rm H}$ fixed) plus a series of nine hydrogen-atmosphere models with radii ranging from 9/8 to 2.8 Schwarzschild radii.  For the more compact models, we included the important effect of self-irradiation of the atmosphere.  Because of the remarkably low column density to J1118, $N_{\\rm H} \\approx 1.2 \\times 10^{20}$~cm$^{-2}$, we obtained very strong limits on a hypothetical thermal source: $k\\Tinf <$ 0.011 keV and $L_{\\infty,\\rm th} < 9.4 \\times 10^{30}$ \\lum~ (99\\% confidence level).  In analogy with neutron stars, there are two possible sources of thermal radiation from a hypothetical surface of J1118: deep crustal heating and accretion.  The former mechanism predicts a thermal luminosity that exceeds the above luminosity limit by a factor of $\\gtrsim 25$, which implies that either one must resort to contrived models or, as we favor, J1118 is a true black hole with an event horizon.  In addition to neutron stars, we also consider emission from several exotic models of compact stars that have been proposed as alternatives to black holes.  As we have shown previously, accreting black holes in quiescent X-ray binaries are very much fainter than neutron stars. One potential explanation for this difference is the larger and hence cooler surface of an 8 \\msun~compact object that might be masked by the ISM.  However, our upper limit on the {\\it total} luminosity of J1118 of $1.3 \\times 10^{31}$ \\lum~is far below the luminosities observed for neutron stars.  This result strengthens our long-held position that black holes are faint relative to neutron stars because they possess an event horizon. ",
        "introduction": "In principle it is possible to detect the radiation emitted from the surface of any ordinary astronomical body such as a planet or a star of any kind.  On the other hand, it is quite impossible to detect any radiation from an event horizon, which is the immaterial surface of infinite redshift that defines a black hole.  This is unfortunate because demonstrating the reality of the event horizon is a problem central to physics and astrophysics.  Nevertheless, despite the complete absence of any emitted radiation, it is possible to marshal strong circumstantial evidence for the reality of the event horizon. One fruitful approach is based on comparing low mass X-ray binaries (LMXBs) that contain black hole primaries with very similar LMXBs that contain neutron star primaries.  In the quiescent state of these systems (McClintock \\& Remillard 2004), the lack of a stellar surface leads to predictable consequences, such as the faintness of black holes relative to neutron stars (Narayan, Garcia, \\& McClintock 1997b; Garcia et al. 2001; Narayan, Garcia \\& McClintock 2002), and also the lack of a thermal component of emission from black holes, which is commonly present in neutron stars (this work).  Similarly, in the outburst state, the presence of a surface in the neutron star (NS) systems gives rise to some distinctive phenomena that are absent in the black hole (BH) systems: (1) type I thermonuclear bursts (Narayan \\& Heyl 2002); (2) high-frequency ($\\sim1$~kHz) timing noise (Sunyaev \\& Revnivtsev 2000); and (3) a distinctive spectral component from a boundary layer at the stellar surface (Done \\& Gierlinski 2003).  In quiescence, almost all accreting neutron stars (e.g., Cen X--4, Aql X--1, and KS 1731--260) display a soft (kT $\\sim$ 0.1 keV) thermal component of emission (see \\S4.2).  The source of the thermal energy is uncertain; for example, it may be energy liberated by the impact of accreting matter (Narayan et al. 1997b, 2002) or crustal energy from the star's interior (Brown, Bildsten, \\& Rutledge 1998; Rutledge et al. 2002b).  In any case, however, the observed X-ray luminosities, temperatures and distances of these NSs indicate that one is observing thermal emission from a source of radius $\\sim10$~km, which is plainly the stellar surface.  On the other hand, no quiescent thermal component of emission has been reported for any of the 15 BH LMXBs (McClintock \\& Remillard 2004), which is the expected result if they possess an event horizon.  These BH spectra are well-represented by a simple power law with photon index $1.5 < \\Gamma < 2.1$ (McClintock \\& Remillard 2004).  However, a soft component of emission might have escaped detection for several reasons.  For example, the quiescent BH LMXBs are fainter, making a soft component more difficult to observe.  Also, compared to the surface of a NS, a hypothetical material surface surrounding a $\\sim10$~\\msun~compact object would be larger and therefore have a correspondingly lower surface temperature.  Herein we search for a thermal component of emission in the quiescent spectrum of XTE J1118+480 (hereafter J1118), a BH LMXB with an extraordinarily high Galactic latitude ($b = 62$\\deg) and correspondingly low interstellar absorption: $N_{\\rm H}~\\approx 1.2~\\times~10^{20}$~cm$^{-2}$ (see \\S2).  For this nominal column depth, the transmission of the interstellar medium (ISM) is 70\\% for the softest X-rays (0.3 keV) that we consider.  Thus, J1118 provides a unique opportunity to search for a soft thermal component of X-ray emission.  For the purposes of this study, we adopt a mass for the BH primary of $M_1 = 8$ \\msun~(McClintock et al. 2001a; Wagner et al. 2001; Orosz et al. 2004).  The quiescent X-ray luminosity is $\\approx3.5 \\times 10^{30}$ \\lum~ (0.3--7 keV; D = 1.8 kpc), which is $10^{-8.5}$ of the Eddington luminosity (McClintock et al. 2003). Both the luminosity and the photon spectral index, $\\Gamma~= 2.02~\\pm~0.16$, are typical for a quiescent BH LMXB with a short orbital period, $P_{\\rm orb} = 4.1$~hr (McClintock et al. 2003).  All of the BH and NS LMXBs that we consider herein are X-ray novae (a.k.a. soft X-ray transients) that undergo bright outbursts lasting several months, which are followed by years or decades of quiescence. During its outburst maximum, J1118 was exceptionally underluminous in the 2--12 keV band, $L_{\\rm x} \\approx 3 \\times 10^{35}$ \\lum, compared to the other BH LMXBs in outburst, $L_{\\rm x} \\sim 10^{38}$ \\lum (McClintock \\& Remillard 2004).  In this work, we determine a strong upper limit on any soft thermal component in the {\\it Chandra} X-ray spectrum of J1118.  Using our upper limit on this emission component, and assuming that J1118 possesses a hypothetical material surface, we set stringent temperature and luminosity upper limits on thermal emission from this surface for a wide range of assumed surface radii; we compare these limits to the observed temperatures and luminosities of quiescent NS LMXBs.  We conclude that the absence of a soft thermal component of emission in the spectrum of J1118 rules strongly against the presence of a material surface and hence argues for the existence of an event horizon.  This work is organized as follows.  In \\S2 we examine the central question of the column density to J1118, and in \\S3 we discuss the observations, data analysis and model-fitting techniques.  The development and computation of stellar atmospheric models appropriate to a compact and massive star are presented in Appendix A; the models include the effects of self-irradiation for stars so compact as to lie within their own photon spheres.  In \\S4, upper limits on the temperature and luminosity of a thermal component from J1118 are summarized and discussed, and these results are compared to the thermal spectra observed for neutron stars.  In \\S5, we interpret the absence of thermal emission form J1118 in terms of two conventional sources of thermal emission from NSs; in addition we consider emission from exotic models of massive, compact stars that have been proposed as alternatives to BHs.  Our conclusions are summarized in \\S6.  J1118 and the other BHs referred to throughout this work are among the 18 dynamically-confirmed BHs; for a review of the properties of these massive, compact X-ray sources, see McClintock \\& Remillard (2004). ",
        "conclusions": "We have examined the possibility that the dynamical black-hole candidate J1118 possesses a material surface rather than an event horizon.  Either accretion onto such a surface or deep crustal heating would be expected to produce a quiescent thermal component of emission like those commonly observed for neutron stars.  We have fitted our Chandra spectrum of J1118 to a model consisting of a fixed power-law component plus an atmospheric thermal component with variable temperature, $\\Tinf$.  The spectral fits were repeated for a series of nine atmospheric models with radii ranging from the minimum allowable, $9/8\\, \\Rs$, to a maximum of $2.8\\, \\Rs$.  For the most compact of these models, which lie within their own photon spheres, the self-irradiation of the atmospheres was taken into account.  No emission in excess of a simple power-law component was detected in J1118, and very strong upper limits were set on the presence of a thermal source: $k\\Tinf <$ 0.011 keV and $L_{\\infty,\\rm th} < 9.4 \\times 10^{30}$ \\lum~(99\\% confidence level).  If one assumes that the hypothetical crust of J1118 is composed of normal nuclear matter, then this stringent limit on a thermal component of luminosity is hard to reconcile with the theory of deep crustal heating and the observed fluence of J1118 during its outburst in 2000: The predicted quiescent luminosity exceeds the above limit on $L_{\\infty,\\rm th}$ by a factor of $\\gtrsim 25$.  Possibly a contrived model of deep crustal heating and/or an extreme model of neutrino cooling of the core could explain this difference.  On the other hand, if J1118 possesses a material surface and accretion powers the thermal emission seen from NSs, then one expects J1118 to have a luminosity at least as great as that of an average NS, whereas its {\\it total} luminosity in Eddington-scaled units is about 100 times less than the luminosity of a typical NS and fully 10 times less than the luminosity of even SAX J1808.4-3658 (Fig. 4).  The above limit on thermal emission, in combination with the observed power-law emission, yields a very tight limit on the {\\it total} quiescent luminosity of $1.3 \\times 10^{31}$ \\lum, which is far below the luminosities observed for NSs.  Because of the high transparency of the ISM, our results rule out the possibility that the total luminosity of J1118 could be augmented significantly by any ultrasoft component of emission.  Thus J1118 -- and by inference the other dynamical BH candidates -- are truly faint relative to NSs (Figs. \\ref{fig:4} and \\ref{fig:5}).  In summary, a sensitive search has failed to detect any thermal emission from a hypothetical surface surrounding J1118, although NSs very commonly show such surface emission due to either deep crustal heating or accretion.  Our sensitivity to a thermal component of emission from J1118 is much greater than the emission predicted by the theory of deep crustal heating, assuming that J1118 has a material surface analogous to that of NSs.  Likewise, there is no evidence that accretion is occurring in quiescence onto the surface of J1118, which is the mechanism often invoked to explain the far greater thermal luminosities of NSs.  The simplest explanation for the absence of any thermal emission is that J1118 lacks a material surface and possesses an event horizon.  Finally, our limits on thermal emission from J1118 rule out the possibility that there is a heretofore unseen and appreciable soft component of luminosity.  This result implies that the dynamical BH candidates are truly faint relative to NSs and underscores our original argument that these compact objects have event horizons and are therefore genuine black holes (Narayan et al. 1997b).  As discussed in \\S 5.3, however, we cannot at this time rule out certain very exotic alternatives.   \\appendix"
    },
    "0403/astro-ph0403298_arXiv.txt": {
        "abstract": "The masses of central massive black holes in BL Lac objects are estimated from their host galaxy absolute magnitude $M_{\\rm R}$ at $R$-band by using the empirical relation between absolute magnitude of host galaxy $M_{\\rm R}$ and black hole mass $M_{\\rm bh}$. Only a small fraction of BL Lac objects exhibit weak broad-line emission, and we derive the sizes of the broad-line regions (BLRs) in these BL Lac objects from the widths of their broad emission lines on the assumption of the clouds being virilized in BLRs. It is found that the sizes of the BLRs in these sources are usually 2-3 orders of magnitude larger than that expected by the empirical correlation $R_{\\rm BLR}-\\lambda L_{\\lambda}$(3000\\AA) defined by a sample of Seyfert galaxies and quasars \\citep{mj02}. We discuss a variety of possibilities and suggest it may probably be attributed to anisotropic motion of the BLR clouds in these BL Lac objects. If the BLR geometry of these sources is disk-like, the viewing angles between the axis and the line of sight are in the range of $\\sim 2^\\circ-12^\\circ$, which is consistent with the unification schemes. ",
        "introduction": "There is a tight correlation between the black hole mass $M_{\\rm bh}$ and the stellar dispersion velocity $\\sigma$ \\citep{fm00,g00}.  This tight correlation $M_{\\rm bh}-\\sigma$ is widely used to estimate the central black hole masses of active galactic nuclei (AGNs). Unfortunately, the stellar dispersion velocity $\\sigma$ is available only for a small fraction of AGNs. \\citet{md02} derived a very tight correlation between host galaxy absolute magnitude $M_R$ at $R$-band and black hole mass $M_{\\rm bh}$. \\citet{cao03} used this relation $M_{\\rm bh}-M_R$ to estimate the central black hole masses of 29 BL Lac objects. The central black hole masses of three sources in their sample have also been measured from the stellar dispersion velocity \\citep{f02,b03}, which agree well with the black hole masses estimated from the host galaxy absolute magnitude $M_R$. \\citet{o02} also found that the black hole masses estimated from the host galaxy luminosity are quite reliable for radio galaxies.  The BLR sizes $R_{\\rm BLR}$ of broad-line ${\\rm H}\\beta$ are measured by \\citet{k00} for a sample of quasars and Seyfert galaxies using reverberation mapping method. They found a tight correlation between the BLR size $R_{\\rm BLR}$ and optical continuum luminosity $\\lambda L_{\\lambda}$. Using the width of the broad emission line and measured  BLR size $R_{\\rm BLR}$, they estimated the central black hole masses of the sources in their sample assuming the clouds in BLRs to be virilized. For sources at high redshifts, the emission of line ${\\rm H}\\beta$ is usually unavailable. Instead, the width of the line Mg\\,{\\sc ii} can be used to estimate the central black hole masses \\citep{mj02}. Mg\\,{\\sc ii} is a low-ionization line as ${\\rm H}\\beta$, so that Mg\\,{\\sc ii} is expected to be produced in the same region as ${\\rm H}\\beta$, which is supported by the tight correlation between the FWHM of Mg\\,{\\sc ii} and ${\\rm H}\\beta$ ($V_{\\rm FWHM}$({\\rm Mg\\,{\\rm\\sc ii}})$\\sim V_{\\rm FWHM}({\\rm H}\\beta$)) found by \\citet{mj02}. It is therefore reasonable to expect that Mg\\,{\\sc ii} and ${\\rm H}\\beta$ are produced in the same region. Using the same sample and the BLR sizes measured by \\citet{k00}, \\citet{mj02} obtained a correlation between $R_{\\rm BLR}$ and the monochromatic continuum luminosity at 3000 {\\AA}, which is useful for estimate of black hole masses of the sources at high redshifts with only Mg\\,{\\sc ii} emission line profiles.  For most AGNs, their BLR sizes have not been measured directly by reverberation mapping method and the empirical relation $R_{\\rm BLR}-\\lambda L_{\\lambda}$ is used to derive $R_{\\rm BLR}$. Combining the line width, the central black hole masses can be estimated by assuming the motion of clouds in BLRs to be virilized. The estimated black hole masses depend sensitively on the velocity of the clouds in BLRs ($\\propto V_{\\rm BLR}^2$). The broad-line width is mainly governed by the component of the cloud velocity $V_{\\rm BLR}$ projected to the line of sight. If the motion of BLR clouds is anisotropic (e.g., disk-like BLR geometry), the estimate of black hole mass becomes complicated (e.g., Jarvis \\& McLure, 2002). \\citet{md01} argued that BLRs in some AGNs have disk-like geometry. If the disk-like BLR geometry is indeed present, the observed broad-line width depends sensitively on the orientation of disk axis and the orientation o is therefore crucial in the estimate of black hole mass from its broad-line width. In the unification schemes, the jets of BL Lac objects are supposed to be inclined at small angles with respect to the line of sight (see Urry \\& Padovani, 1995 for a review), so the broad-line profiles will be significantly narrowed if disk-like BLRs are present perpendicular to the jets, which can be used to test the geometry of BLRs in BL Lac objects. In this paper, we use the observed broad emission line widths and black hole masses derived from host galaxy luminosity to test the geometry of BLRs in BL Lac objects if the black hole masses can be estimated by independent method. The cosmological parameters $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\Lambda}=0.7$, and $H_0=70~ {\\rm km~s^{-1}~Mpc^{-1}}$ have been adopted in this paper. ",
        "conclusions": "We plot the relation between the ionizing luminosity $\\lambda L_{\\lambda}(3000)$ at 3000 {\\AA} and  BLR size $R_{\\rm BLR}$ in Fig. \\ref{fig1}. The BLR size $R_{\\rm BLR}$ is derived from the width of broad line Mg\\,{\\rm\\sc ii} and black hole mass $M_{\\rm bh}$ on the assumption of isotropic motion of clouds in the BLR, i.e., $f=\\sqrt{3}/2$ is adopted in Eq. (\\ref{rblrmbh}). We can also derive the BLR size $R_{\\rm BLR}^{\\rm emp}$ using the empirical relation (\\ref{rblr3000}). If the motion of BLR clouds is indeed isotropic, one may expect similar BLR sizes derived by these two different methods. However, it is found that the sizes of BLRs $R_{\\rm BLR}$ in all these sources are $\\sim$2-3 orders of magnitude larger than $R_{\\rm BLR}^{\\rm emp}$ expected by relation (\\ref{rblr3000}).  We note that only upper limits on the masses of the black holes in 9 of all 16 sources. The BLR sizes may be over-estimated for these 9 sources, because of the BLR sizes $R_{\\rm BLR}$  being derived from the line widths and black hole masses (see Eq. \\ref{rblrmbh}). The black hole masses estimated from the host galaxy luminosity for these 16 sources are in the range of $\\sim 10^{8.6-10.4} M_{\\odot}$ (see Table 1).  The deviations of the BLR sizes $R_{\\rm BLR}$ from $R_{\\rm BLR}^{\\rm emp}$ expected by relation (\\ref{rblr3000}) cannot be solely attributed to the overestimate of black hole masses for those nine sources with upper limits on galaxy luminosity, unless the black hole masses have been overestimated by 2$-$3 orders of magnitude, i.e., the realistic black hole masses should be in the range of $\\sim 10^{6-8}$ $M_\\odot$ for these sources, which seems impossible. It will be more difficult to attribute such deviations to the overestimate of black hole masses for those 7 sources with well measured host galaxy luminosity.  The black hole mass of the source $1807+698$ has been measured from its stellar dispersion velocity $\\sigma$ \\citep{f02,b03}, which is consistent with our estimate of the black hole mass $10^{8.88} M_{\\odot}$. For this source, its BLR size $R_{\\rm BLR}$ derived from relation (\\ref{rblrmbh}) is about three orders of magnitude higher than $R_{\\rm BLR}^{\\rm emp}$ predicted by relation (\\ref{rblr3000}) between $R_{\\rm BLR}^{\\rm emp}$ and $\\lambda L_{\\lambda}(3000)$.  \\figurenum{1} \\centerline{\\includegraphics[angle=0,width=10.0cm]{a0312f1.ps}} \\figcaption{\\footnotesize  The relation between the BLR size $R_{\\rm BLR}$ and the ionizing luminosity $\\lambda L_{\\lambda}(3000)$ at 3000 {\\AA}  derived from narrow-line luminosity $L_{[{\\rm O}_{\\rm II}]}$. The line represents the correlation $R_{\\rm BLR}^{\\rm emp}-\\lambda L_{\\lambda}(3000)$ defined by Seyfert 1 galaxies and quasars  \\citep{mj02}. The source names are labelled in the plot.\\label{fig1}} \\centerline{}     The ionizing luminosity at 3000 {\\AA} is derived from the luminosity of narrow-line [O\\,{\\sc ii}] adopting EW$_{\\rm ion}$ =10 {\\AA} \\citep{w99}, which is in general consistent with the optical spectroscopic observations on the radio quasars selected from the Molonglo Quasar Sample (MQS), of which the equivalent widths of [O\\,{\\sc ii}] in the source frame are in the range from less than 1 {\\AA} to more than 100 {\\AA} with an average of 14.7 {\\AA} \\citep{b99}.  However, using a single value EW$_{\\rm ion}$ =10 {\\AA} is still a rough estimate, and may induce uncertainties on the estimates of ionizing luminosity. In Fig. \\ref{fig1}, we find that the deviations of $R_{\\rm BLR}$ from $R_{\\rm BLR}^{\\rm emp}$ will not be solved on the isotropic BLR geometry assumption even if a rather small  EW$_{\\rm ion}$=1 {\\AA} is adopted, i.e., the ionizing luminosity is an order of magnitude higher than the present values. In Table 1, we have listed the measured equivalent widths of [O\\,{\\sc ii}] in the source frame for the sources in our present sample, which are in 0.2$-$9.2 {\\AA}. Considering the optical continuum emission contributed by the beamed emission from the jets in these BL Lac objects, their equivalent widths (EW$_{\\rm ion}$) corresponding to the ionizing optical continuum emission should be larger than the measured values listed in Table 1. This implies that the uncertainties on the estimates of the ionizing optical continuum luminosity would not change the conclusion on this point.  If photo-ionization of the gases in narrow-line region by radiative shocks driven by the radio source is important (e.g., Inskip et al. 2002), the central ionizing luminosity should be lower than the present values, which would lead to even larger deviations.   \\citet{wz03} found that the BLR sizes of dwarf AGNs are systematically larger than the prediction of $R_{\\rm BLR}-\\lambda L_{\\lambda}$ correlation defined by Seyfert galaxies and quasars \\citep{k00}. They suggested that the flat ionizing spectra are in these dwarf AGNs as predicted by advection dominated accretion flow (ADAF) models, and the BLRs in those dwarf AGNs have lower ionization or/and lower density than those in Seyfert 1 galaxies and quasars of \\citet{k00}'s sample. The sources in our present sample have brighter ionizing luminosity than those dwarf AGNs. As most BL Lac objects do not exhibit any line emission, all these 16 BL Lac objects have measured broad-line profiles and narrow-line emission have relatively high ionizing luminosity amongst all BL Lac objects \\citep{cao02a}. There is a critical accretion rate $\\dot{m}_{\\rm crit}$, and an ADAF can exist only if accreting at a rate $\\dot{m}<\\dot{m}_{\\rm crit}$(e.g., Mahadevan, 1997), which leads to  an upper limit on optical continuum emission from an ADAF for a given black hole mass \\citep{cao02a}. \\citet{cao03} calculated optical spectra of ADAF$+$SD(standard disk) systems and compared them with the observed spectra, which suggests that only ADAFs themselves are unable to produce such bright ionizing optical continuum emission and standard thin disks should be present at least in the outer regions of the disks for these BL Lac objects. The ionizing photos are therefore mainly from the standard thin disk regions  in these BL Lac objects, unlike the dwarf AGNs considered by \\citet{wz03}. The luminosity $\\lambda L_{\\lambda}(3000)$ of the sources in the sample used to derive the correlation $R_{\\rm BLR}-\\lambda L_{\\lambda}(3000)$ is in the range of $\\sim 10^{34-39}$ W \\citep{mj02}.  The ionizing luminosity of the BL Lac objects in present sample at 3000 {\\AA} are in the range of $\\sim 10^{35.8-38}$ W, which is in the similar range as their sample. The $R_{\\rm BLR}-\\lambda L_{\\lambda}(3000)$ relation derived from a sample of quasars and Seyfert galaxies should be valid for these BL Lac objects, unless the physics of BLRs in these BL Lac objects is significantly different from that in the sources of the sample considered by \\citet{mj02}.  We derive the BLR sizes of these BL Lac objects assuming the motion of clouds to be isotropic, which may not be the case in these sources. Our estimate of  BLR sizes may strongly be overestimated due to anisotropic cloud motion, if the velocity component projected to the line of sight is only a small fraction of its real velocity. A most likely candidate for such anisotropic motion of clouds is the clouds orbiting in a disk-like BLR (e.g., McLure \\& Dunlop, 2001). For the clouds orbiting in a disk-like BLR, the correction factor $f$ in Eq. (\\ref{rblrmbh}) is $1/(2\\sin i)$, where  $i$ is the angle of the axis inclined to the line of sight \\citep{md01}. If this is the case, we can estimate the inclination angle $i$ of these BL Lac objects assuming they indeed to obey the correlation $R_{\\rm BLR}-\\lambda L_{\\lambda}(3000)$ suggested by \\citet{mj02}, i.e., the value of $f$ is estimated by letting $R_{\\rm BLR}=R_{\\rm BLR}^{\\rm emp}$. The derived results are listed in Table 1. We find that the inclination angles are around $\\sim 2^\\circ-12^\\circ$ for these BL Lac objects.  There is evidence that the velocity field of BLR is better described by a combination of a random isotropic component, with  characteristic velocity $V_{\\rm r}$, and a component only in the plane of the disk, with characteristic velocity $V_{\\rm p}$ \\citep{wb86}. In this case, the observed FWHM will be given by \\be V_{\\rm FWHM}=2(V_{\\rm r}^2+V_{\\rm p}^2\\sin^2 i)^{1/2} \\ee\\citep{md01}, so $f=0.5[(V_{\\rm r}/V_{\\rm p})^2+\\sin^2 i]^{-1/2}$. If the random isotropic component is important, i.e., $V_{\\rm r}$ is comparable with $V_{\\rm p}$, then the term $(V_{\\rm r}/V_{\\rm p})^2$ cannot be neglected and the derived inclined angle of the disk axis will be less than that listed in Table 2. This is in general consistent with the unification schemes that the jets of BL Lac objects are inclined at small angles to the line of sight."
    },
    "0403/astro-ph0403584_arXiv.txt": {
        "abstract": "{ GRB 020410 is by far the longest $\\gamma$-ray burst (with a duration of about 1600~s) to have been followed up from the X-ray through the radio regime.  Afterglow emission was detected in X-rays and at optical wavelengths whereas no emission was detected at 8 GHz brighter than 120 $\\mu$Jy.  The decaying X-ray afterglow, back extrapolated to 11 hours after the burst, had a flux of $7.9\\times10^{-12}$ \\fu\\ (2--10 keV); the brightest detected so far. No direct redshift determination is available yet for this GRB, but according to the empirical relationship between the peak energy in the $\\nu F_\\nu$ spectrum and the isotropic energy output, $z$ is constrained in the range 0.9--1.5. The reconstructed optical afterglow light curve implies at least two breaks in the simple power-law decay.  This may be related to emergence of a SN, or refreshment of the external shock by a variation in the circumstellar medium. By comparing the backward extrapolation of the 2--10 keV afterglow decay, it is shown that the long duration of the prompt emission is not related to an early onset of afterglow emission, but must be related to prolonged activity of the ``central engine''.} ",
        "introduction": "Gamma-ray Bursts (GRB) show great diversity with regard to both their durations and spectral properties. GRBs last from a fraction of a second to thousands of seconds, as established by the BATSE survey (e.g. \\cite{Paciesas99}). Prompt X-ray counterparts of GRBs, detected by Ginga, \\sax, HETE-2 have a very wide distribution of intensities and durations.  Tails and precursors of X-ray counterparts were also observed by WATCH/Granat (\\cite{ct94}).  Those events characterized by an X-rays-to-$\\gamma$-rays (2--10/40--700 kev) fluence ratio larger than $\\sim0.5$ are classified as X--ray rich (e.g. \\cite{Feroci01}). Moreover, transient X--ray sources with characteristics similar to those of GRB counterparts, although with no simultaneous GRB detection (so called ``X-ray flashes''; \\cite{Heise01, moc03}) were detected by the \\sax\\ Wide Field Cameras (WFC) and, subsequently, by the HETE--2/FREGATE instrument. Recently (\\cite{Zand03}), have reported the detection of 4 long, faint X-ray transients during sky surveys with the \\sax-WFC. Three of these are confirmed GRBs, because they coincide with BATSE detections. They show durations ranging from 540 s to 2550 s and are characterized by a mildly soft spectrum.  The very different ratios of $\\gamma$-ray vs X-ray peak fluxes or fluences point either to different viewing angles of the relativistic jets in which GRBs are formed (e.g. \\cite{Granot02, Yamazaki02}) or to a different amount of baryon contamination of the fireball (e.g. \\cite{Dermer99, Huang02}). Furthermore, the existence of a class of GRBs with long X-ray durations is important for the investigation of the connection between the prompt and afterglow components and may also suggest a high redshift origin.  However, redshift constraints imposed on XRF 020903, 030723 and, possibly, 031203 do not support the high redshift scenario (\\cite{Soderberg03, Prochaska03, Fynbo04}).  \\grb, first detected in X-rays only by the \\sax-WFC (\\cite{Gandolfi02}), stands out for its long duration, more than 1500 s in the 2--28 keV band (see Sect. \\ref{wfc_an}), and for the relative weakness of its $\\gamma$-ray signal, detected by Konus-Wind in an offline analysis.  Based on its X-ray-to-$\\gamma$-ray fluence ratio (see Sect. 3), \\grb\\ lies in the soft tail of genuine GRBs and marginally qualifies as an ``X-ray rich'' GRB (see \\cite{Heise01}). Upon detection of the GRB we started an X-ray and optical search and monitoring campaign of its afterglow.  We present here the results of our study of the prompt and afterglow emission. ",
        "conclusions": "\\grb\\ is by far the longest GRB event for which X-ray afterglow emission and an optical counterpart (though weak) have been discovered.  The long duration of \\grb\\ can be considered ``tail of the distribution'' rather than a peculiar case, though peculiar circumstances are required.  As the other very long GRBs detected by the WFC (\\cite{Zand03}), this event shows a high ratio between X-ray and $\\gamma$-ray fluence, although it cannot be classified as an X-ray rich GRB. In addition, apart the high X-ray content, X-ray rich GRBs are also characterized by the lack of optical afterglow (not true for the X-ray flashe 030723 and possibly 020903), which is not the case for \\grb.  We note, however, that the two things could not be linked as the rapidly fading OT of the ``not-quite'' X-ray rich GRB 021211 has demonstrated (\\cite{crew03}). Spectroscopically, \\grb\\ does not show peculiarities.  This in addition to (a) the measured upper limit on the intrinsic absorption, (b) the ``smooth\" increase of the spectral index, (c) the late time X-ray ``afterglow\" decay that does not connect to the late time prompt emission, suggest the prompt event is all due to internal shock rather than being a superimposition of internal and external shocks, or at least the latter is negligible.  In the fireball model, a connection between the long duration of the ``prompt'' event and the slow decay of the afterglow could be accounted for e.g.  by the superimposition of several (external) shocks produced by each of the main peak of the internal shock, which are all relatively long when compared to normal GRBs, causing a continuous refreshment of the external shock (e.g. \\cite{bj02}).  Unlike what is observed for several GRBs (e.g. \\cite{costa99}) the backward extrapolation of the afterglow fading law of \\grb\\ is inconsistent with the flux measured during the last part of the prompt emission (see Fig.~\\ref{f4_xdecay}). This may be linked to the extremely long duration of the event and prevents us to derive an indication of the afterglow emission onset time.  We can estimate an upper limit to the afterglow 2--10 keV fluence if, by following \\cite{Frontera00}, we assume that the afterglow emission starts at 63\\% of the duration of the GRB and thus we integrate the fading law between 973 s and $1\\times10^6$ s.  The result is $1.96\\times10^{-6}$ erg cm$^{-2}$, corresponding to about 34\\% of the fluence measured in the prompt event in the same energy range and to about 9\\% of the prompt fluence in 40--700 keV. These values are well within the observed range of normal GRBs (\\cite{Frontera00}).  Alternatively, it is possible to identify the onset of the external shock at $t\\sim500$~s when the spectrum of the prompt emission becomes consistent with the late time MECS spectra. In the simple case in which the fireball is homogeneous and thin, the GRB variability should be suppressed and the lightcurve be described as a power-law initially rising as $t^2$ and then smoothly turning over to a decay slope which depends on the spectral range and dynamics of the fireball (\\cite{SaPi99}). In fact the lightcurve of \\grb\\ is highly variable after the spectral transition, showing a prominent emission episode at $t~1500$~s (P4 in Fig.~\\ref{f1_promptlc}). This behavior can be understood if the inner engine does not turn off at the end of the gamma-ray phase, but releases a sizable amount of energy at $t\\sim1500$~s. This late emission, however, should be inefficient in the production of $\\gamma$-rays or, in terms of the internal-external shock scenario, it should avoid the internal shock phase. The time $t\\sim1500$~s is not the deceleration time of the fireball, but the delay with which the inner engine released the fireball component that re-energized the external shock to produce the P4 rebrightening. The cause of the lack of $\\gamma$-ray emission associated with the delayed energy release is not clear and, lacking WFC data for the P4 episode, it is difficult to constrain observationally; though the slight count excess in the Konus soft band could be an hint. The delayed energy release may however be associated to the recycling of the energy wasted while the relativistic jet propagates into the host star (\\cite{Mesz01, Ram02}). In that case, the acceleration of the delayed fireball takes place at the star surface, and is therefore characterized by a variability timescale many orders of magnitude larger than that of the jet, effectively preventing the occurrence of internal shocks.  We also note that among the other GRBs afterglows for which the extrapolation of the decay law to the prompt emission is inconsistent with the observed flux, GRB 990704 (the X-ray richest event observed by \\sax) is the only analogous case. The afterglow X-ray flux decay of XRF 031203 also shows an extrapolation below the ``probable'' prompt flux (\\cite{Watson04}). GRB 990510, 010222, 010214 show an extrapolation \\emph{above} the prompt emission, which is explained with a break a few hours after the onset (\\cite{Pian01, Zand01, Guidorzi03})  The peak width dependence as function of the energy was tested for P1 and P3 (see Fig.~\\ref{f1_promptlc}). To this aim we produced rebinned light curves with bin size between 1 and 8 s. Their FWHM were obtained using Gaussian fits; a 20\\% systematic error on their estimate was added. By using a law ${\\rm FWHM}=kE^\\alpha$ (expected by the synchrotron model, Fenimore et al. (1995)) for the two peaks we obtain $\\alpha=-0.48\\pm0.20$ for the P1 and $\\alpha=-0.44\\pm0.12$ for P3 (errors are 90\\% confidence level). These results are consistent with the results from the BATSE GRBs (\\cite{Fenimore95}) as well as for GRB 960720 (\\cite{Piro98}) and 990704 (\\cite{Feroci01}).  Due to the paucity of data and the complexity of the optical light curve, it is not possible to constrain the fireball and environment properties completely. From X-ray spectroscopy we infer the electron distribution slope $p=2.1\\pm0.25$ under the assumption that X-rays are above the synchrotron cooling frequency. Due to the large uncertainty, the X-ray decay slope of $\\delta_{\\rm X}=0.81\\pm0.07$ can be accounted for both in an ISM and wind environment. It is tantalizing to note, however, that the early time optical slope seems to be larger than the X-ray one. This would fit in a wind environment scenario, consistent with the possible detection of a supernova bump at late time (\\cite{Levan04}). In this case one would expect $\\delta_{\\rm X}=(3p-2)/4=1.1\\pm0.2$ and $\\delta_{\\rm O}=(3p-1)/4=1.3\\pm0.2$, fully consistent with the X-ray slope and the optical lower limit. Even this interpretation bears however some degree of uncertainty. \\grb\\ has a flux of 10.5 $\\mu$Jy in R and $7.9\\times 10^{-12}$ \\fu\\ in X-rays which falls outside the distribution found by \\cite{depas02}) (see their Fig.~5) and would classify it as a dark GRB (\\cite{laz02}). Even assuming that the cooling frequency lies exactly at the edge of the \\sax\\ band, the synchrotron spectrum would over-predict optical emission by a factor $\\sim5$. There are two possible interpretations for this. One possibility is that the X-ray emission is boosted by an IC component, like in the case of GRB 000926 (\\cite{har01}). This would require a moderately dense environment, either uniform or stratified. Alternatively, the optical emission may be extincted by a sizable amount of dust in the host galaxy, with $A_{\\rm V}\\sim2$. This would correspond, for a Galactic mixture, to a column density $\\nh\\sim3\\times10^{21}$~cm$^{-2}$, consistent with the upper limit derived from X-ray spectroscopy. The lack of constraints on the optical spectrum prevents us to reach a definite conclusion. The optical spectrum should be bluer in the case of IC emission and red in the case of dust obscuration.  Assuming that the emission line in the MECS spectra is real and due to fluorescence of H-like iron (rest energy of 6.97 keV), then the change in line position can be explained by a variable iron recombination edge showing its maximum in the second half of ToO~1 (or later, but before ToO~2). In fact if we derive the redshift from the line position in ToO~2, we obtain $z\\simeq1.7$ which leads to a recombination edge of $\\sim3.4$ keV. Also the ratio between the iron recombination edge rest energy, 9.28 keV, and 6.97 keV is $\\simeq 1.3$, like for the ratio of the ToO~1b over ToO~2 line energies. Again, our statistics does not allow us to perform a simultaneous fit for a Gaussian and a recombination edge line. However this hypothesis appears in agreement with the data.  As no direct $z$ measurement exists for \\grb, we calculated the peak energy $E_{\\rm p}$ in the $\\nu F_\\nu$ spectrum and the isotropic energy $E_{\\rm rad}$ for a grid of $z$ values. We then compared the results with the relation reported by Amati et al. (2002). We find that the relation is satisfied (with a discrepancy level $<20\\%$) for $0.9 < z < 1.5$ and $1.1\\times10^{53}< E_{\\rm rad} < 3.0\\times10^{53}$ erg. This range of $z$ would exclude the value of 0.5 obtainable by assuming a 1998bw-like SN re-bump (\\cite{Levan04}) and is marginally in agreement with the value derived above. Even assuming that our flux estimate for the missing part of the X-ray light curve must be increased by an extra 20\\% (which is unlikely), the lower limit for $z$ becomes 0.6. Besides this, we note that $z\\simeq0.5$ together with the reported magnitude of ${\\rm V}\\simeq 28.7$ for the host galaxy (\\cite{Levan04}) would place it in the very low end of the galaxy luminosity function (${\\rm M_{\\rm V}} = -14.3$), which is unusual for GRB hosts; this independently of considering the X-ray spectrum derived $\\nh\\simlt3\\times10^{21}$~cm$^{-2}$ as being ``local'' or ``global''."
    },
    "0403/astro-ph0403067_arXiv.txt": {
        "abstract": "We summarize some recent results from our observational campaign to study the central regions of spiral galaxies of late Hubble type (Scd - Sm). These disk-dominated, bulgeless galaxies have an apparently uneventful merger history. The evolution of their nuclei thus holds important constraints on the mechanisms that are responsible for bulge formation and nuclear activity in spiral galaxies. We discuss the structural properties, stellar populations, and dynamical masses of the compact, luminous star cluster that is found in the nuclei of most late-type spiral galaxies. Although preliminary, our results strongly indicate that many galaxies of our sample experience repeated periods of nuclear star formation. While the exact mechanism that leads to the required high gas densities in the galaxy nucleus remains unclear, results from our recent CO survey of late-type spirals demonstrate that in most cases, the central kpc contains enough molecular gas to support repetitive nuclear starbursts. ",
        "introduction": "In most formation scenarios for spiral galaxies, the central bulge is the ``trashbin of violent relaxation'' where a dynamically hot stellar component has formed either through external potential perturbations such as early mergers of proto-galaxies (e.g. Carlberg 1992), or perhaps via internal effects such as violent bar instabilities (e.g. Norman, Sellwood, \\& Hasan 1996). The latest-type spirals, then, must have lived very sheltered and uneventful lives, since their central ``trashbin'' is virtually empty, i.e. they are devoid of prominent starburst events, have no discernible stellar bulges, and rarely show signs of nuclear activity. These galaxies often have gently rising rotation curves (e.g. Matthews \\& Gallagher 1997) that indicate a nearly homogeneous mass distribution on scales $\\sim 1\\kpc$. On these scales, gravity therefore hardly provides a vector pointing at the center, and it is not obvious that the nucleus of these galaxies is well-defined and a unique environment.  \\begin{figure}[!ht]  \\vspace*{4cm} \\centerline{see attached file fig1a.gif} \\plotone{fig1b.ps} \\caption{Top: HST/WFPC2 F814W (I-band) images of three representative nuclear star clusters. Shown is the PC chip with a field of view of $\\approx 35\\arcsec \\times 35\\arcsec$. The bar in the top left of each panel denotes a spatial scale of 1 kpc, the north-east orientation is indicated by the compass arrow. Bottom: I-band surface brightness profiles, measured from elliptical isophote fits to the images above. Note the clear transition between the underlying disk and the NC at radii around $0.2\\arcsec$. Also shown are analytical fits to the disk and cluster profiles that yield the cluster photometry (for details, see B\\\"oker et al. 2002.)} \\label{fig:images} \\end{figure} Surprisingly, the photocenter of many late-type spiral galaxies is nevertheless occupied by a compact, luminous stellar cluster (Phillips et al. 1996; Carollo et al. 1998; Matthews et al. 1999; B\\\"oker et al. 2002, hereafter Paper~I). Figure~\\ref{fig:images} shows HST/WFPC2 I-band images of three representative examples of such nuclear star clusters (NCs), together with their surface brightness profiles as measured from elliptical isophote fits to the WFPC2 data. The typical luminosities of NCs are in the range $10^6$ - $10^7\\,\\lsun$ (Paper~I). Nuclear clusters are therefore much brighter than average stellar clusters in the disks of nearby spiral galaxies (e.g. Larsen 2002), and comparable to young ``super star clusters'' in luminous merger pairs (Whitmore et al. 1999) or circumnuclear starforming rings in spiral galaxies (e.g. Maoz et al. 2001).  Knowledge of the stellar populations and masses of NCs is essential in order to constrain the mechanism(s) that lead to their formation. The age(s) of the stellar population(s) can be determined from spectral synthesis methods applied to medium-resolution spectra of NCs. Measuring the stellar masses of NCs requires both accurate knowledge of the velocity dispersion (from high-resolution spectra) and the cluster light distribution (from HST imaging). This kind of analysis has been successfully applied to NCs (B\\\"oker et al. 1999) as well as young clusters in the disks of nearby galaxies (Smith \\& Gallagher 2001; Mengel et al. 2002, McCrady et al. 2003). In this paper, we summarize our ongoing study of the properties and formation mechanism(s) of NCs in late-type spirals. ",
        "conclusions": "In ellipticals and ``classical'' spirals, phenomena that indicate nuclear activity such as AGN and massive black holes have received much attention over the past decades. Only recently, however, has it become evident that even in bulgeless, ``pure'' disk galaxies which are generally devoid of any obvious signs for nuclear activity, the galaxy center is a ``special'' location in that it is occupied by a massive, compact, and often young stellar cluster. These nuclear star clusters (NCs) have sizes and luminosities that are very similar to those of young super star clusters (SSCs) found in the disks of starburst galaxies. However, NCs most likely have a different, more complex, formation history, as can be inferred from population synthesis age-dating of their spectra as well as estimates of their masses from stellar dynamical modelling.  Unless one is willing to believe that we live in a special time, the fact that many nuclear star clusters are young suggests that they experience repetitive ``rejuvenation'', most likely due to infall of molecular gas into the central few pc and associated star formation. This is not implausible, because bulgeless spirals are in many ways normal spirals. In particular, their molecular gas content follows the same scaling relation with galaxy luminosity as earlier-type spirals, and typical molecular gas masses in their central kpc are of the order $10^7\\,\\msun$.  While the ``duty cycle'' for nuclear starbursts still needs to be reliably established - which is one of the goals of our project - the possibility of repetitive nuclear cluster formation in late-type spirals has interesting consequences for galaxy evolution scenarios, both with respect to morphological classification (i.e. Hubble type) and nuclear activity (i.e. black hole growth)."
    },
    "0403/astro-ph0403317_arXiv.txt": {
        "abstract": "We introduce a model for pulsars in which non-radial oscillations of high spherical degree (\\el) aligned to the magnetic axis of a spinning neutron star reproduce the morphological features of pulsar beams.   In our model, rotation of the pulsar carries a pattern of pulsation nodes underneath our sightline, reproducing the longitude stationary structure seen in average pulse profiles, while the associated time-like oscillations reproduce ``drifting subpulses''---features that change their longitude between successive pulsar spins.  We will show that the presence of nodal lines can account for observed $180^\\circ$ phase jumps in drifting subpulses and their otherwise poor phase stability, even if the time-like oscillations are strictly periodic. Our model can also account for the ``mode changes'' and ``nulls'' observed in some pulsars as quasiperiodic changes between pulsation modes of different (\\el) or radial overtone ($n$), analogous to pulsation mode changes observed in oscillating white dwarf stars. We will discuss other definitive and testable requirements of our model and show that they are qualitatively supported by existing data. While reserving judgment until the completion of quantitative tests, we are inspired enough by the existing observational support for our model to speculate about the excitation mechanism of the non-radial pulsations, the physics we can learn from them, and their relationship to the period evolution of pulsars. ",
        "introduction": "\\label{intro}  Upon the discovery of radio pulsations from pulsars by \\citet{hew68}, \\citet{rud68}~ proposed that the pulses arose from non-radial oscillations of a neutron star.  This idea was quickly displaced by a rotational model \\citep{gol69}, but \\citet{dra68}~ again raised the possibility of pulsations when they measured individual pulse sequences for two pulsars and found within them narrow subpulses that moved to successively earlier times within the main pulse.  Because this drift represented the presence of a ``second periodicity'' incommensurate with the spin period, it was natural to propose a time-like oscillation of the star.  Subsequent measurements, however, revealed complex subpulse patterns that did not conform to a pulsation model in any obvious way. Moreover, the persistence of unique subpulse shapes from pulse to pulse, along with problems of phase stability we will address in later sections, led Drake to conclude that the drifting subpulses were incompatible with the pulsation hypothesis \\citep[see][]{sta70,hew70}. Ultimately, pulsations were abandoned in favor of purely geometric models, although they reappeared from time-to-time in the theoretical literature \\citep[notably][]{hans80,vanh80,mcd88first,car86,fin90,rei92,str93}. Most recently \\citet{dun98}~ invoked toroidal modes to account for oscillations of soft gamma repeaters, but other than the work of \\citet{str92}~ and \\citet{scvfirst92}, there has been no determined attempt to account for the properties of classical pulsars with models involving non-radial pulsations.  Instead, most current models, though not all \\citep[cf.][]{lyn88,han01}, incorporate a circulating pattern of sub-beams, whose motion about the magnetic pole produces the drifting subpulses.  In these models, pulsar emission comes from accelerated particles that originate near the pulsar magnetic pole and travel along curved paths in the star's magnetic field \\citep[see][]{rad69,kom70}.  The radiation is confined to a narrow beam by the dipole magnetic field geometry \\citep{gj69} and relativistic beaming along the direction of particle motion, which is roughly parallel to the magnetic axis, not perpendicular as in the models of \\citet{gol69}, \\citet{smi70} and \\citet{zhe71}.  The observed brightness of pulsar beams effectively demands that the radiation is coherent, but the question of how it is produced is not settled \\citep{jes01,les98,mel95}.  Early studies of pulsar single pulses and average pulse shapes \\citep[][and others]{tay75first, lyn71}~ led to the addition of more elaborate emission structures within the model pulsar beam. These features sweep past our sightline and recreate the variety of pulse shapes we observe.  \\citet{bac76}~ described a target-shaped emission pattern (a central core surrounded by an annulus) that can reproduce a wide variety of pulse morphologies depending on whether our sightline crosses the center of the pattern, yielding a three component pulse, or crosses only the annulus, resulting in a one or two component pulse.  \\citet{ost77} added a second annulus and rotating features to reproduce pulses with more than three components and drifting subpulses.  In 1975, \\citet{rud75} supplied a physical basis for the model by suggesting that the emission arises from localized discharges or sparks near the polar cap. These are arranged in annular patterns, and rotate naturally due to the crossed components of the magnetic and electric fields.  In addition to the fixed and drifting substructure, models must account for observations of two kinds of discrete events observed in some pulsars; ``mode changes'', which abruptly alter the character of the substructure, and pulse ``nulling'', during which the pulse emission drops below detectable levels for one or more spin periods of the pulsar \\citep{bac70a,bac70b,bar82}.   In the  \\citet{rud75} model, mode changes and nulling result from a collapse or reorganization of the fixed and moving spark structures, after which they must reappear with the same features they had previously.  Several reviewers have summarized observational and theoretical progress in the study of pulsar beams.   The most ambitious is Rankin \\citep{ran83a,ran83b,ran86,ran90,rad90,ran93,mit02}, who has both reviewed and synthesized the observations into an empirical model incorporating polarization and spectral behavior.  \\citet{man95} gives a somewhat different view of the beam geometry. Most recently, \\citet{gra03} has published a succinct review that includes both ``normal'' and millisecond pulsars.  Against this backdrop, as a student project, we conducted a re-analysis of archival data on PSR0943+10 to look for evidence of non-radial pulsations, which, according to theory, might have periods ranging from milliseconds to seconds \\citep{mcd88}.   Our analysis, which will appear in a subsequent paper, convinced us that time-like oscillations with a period of $31.8$ msec are a viable alternative to the rotating carousel of emission beams proposed by \\citet{des01}, but we could find no compelling reason other than aesthetics to~{\\it prefer} a pulsational model.  In search of a definitive test, we reviewed the extensive observational literature on pulsars, and found intriguing evidence for non-radial pulsations as a universal mechanism for drifting and stationary subpulses. Moreover, we found that the original reason for abandoning pulsational models does not apply to non-radial pulsations of high azimuthal degree (\\el) in which our sightline crosses pulsation nodal lines.  The presence of nodal lines increases the variety and subtlety of expected subpulse behavior.  The purpose of this paper is to introduce a model in which high \\el~ pulsations aligned to the pulsar magnetic pole take the place of the fixed and moving structures of the circulating spark model, but other details of the geometry remain unchanged.  In this paper we will explore only the phenomenological consequences of this substitution, and compare them qualitatively to published observations.  We will not discuss in any detail problems in the physics of pulsed radio emission or polarization mechanisms.  In \\S\\ref{morph}, we will present the basic features of our model, and explore its observational properties, some of which are not immediately obvious.  Our main purpose is to lay the groundwork for future application of the model to radio measurements of individual and average pulse profiles.   In \\S\\ref{comp} we will examine qualitative evidence in favor of our model, reserving quantitative comparisons for subsequent papers.  The strongest evidence we will present comes from published measurements that show subpulse phase behavior difficult or impossible to explain using the circulating spark model, but demanded by high \\el~ pulsations. We will also discuss analogies between pulsar behavior and that of known pulsating stars, specifically the rapidly oscillating peculiar A stars (roAp) and the pulsating white dwarf stars. This will demonstrate that there are precedents for the model behavior we propose.  In \\S\\ref{disc}, we will speculate about theoretical aspects of our model, such as the pulsation driving mechanism, and we will introduce the notion of ``horizontal mode trapping'', which can account for the high \\el~ character of the proposed modes and relate them to the observed period evolution of pulsar beam widths.  We will end by highlighting the potential for neutron star seismology, which can yield direct measurements of interesting physical quantities like the buoyancy of neutron star surface oceans. ",
        "conclusions": "\\label{conc}  Whether or not the foregoing discussion has revealed anything about pulsars, it has certainly demonstrated \\citet{cle83} maxim that we can get  ``wholesale returns of conjecture out of such a trifling investment of fact.''  Nevertheless, the fact remains: pulsar beams show subpulse phase reversals at the longitude-stationary boundaries separating individual pulse components. We have shown that these changes are comprehensible in the context of an oblique pulsator model incorporating non-radial pulsations of high degree \\el.  The important features of our model are: non-radial oscillations aligned to and symmetric about the pulsar {\\it magnetic} axis; surface displacements that follow a spherical harmonic distribution; radio emission that follows the displacements but is never negative; pulsation modes of sufficiently high \\el~ that nodal lines often appear in the pulse window; and pulsation frequencies that remain coherent over many pulsar spin periods. Variations on this basic model might include multiple pulsation modes with non-zero azimuthal orders, pulsations that are distorted, in reality or in appearance, by non-dipole fields, and modes that interact either through mode coupling or a non-linear emission mechanism.  Our model qualitatively reproduces the mean shapes of pulsar beams and the radio frequency dependent behavior of subpulses with a minimum of free parameters. In the most basic form of the model these are $\\alpha$, $\\beta$, \\Ptime, \\el, and \\Pone.   Our model also dictates specific requirements that can be tested quantitatively using new or archival data. We have embarked on a program to conduct such tests and we encourage others to do likewise. If the model survives these tests, then we will have the opportunity to measure fundamental properties of matter in a domain not accessible to laboratory experiments. The first challenge will be to determine the site of the pulsations, and then to connect measured eigenfrequencies with the eigenmodes of a structural model.  Given the number of modes in the pulsation spectrum at large \\el, this may be a daunting task, but even rough identification will provide limits on the thermal, electrical, and mechanical properties of constituents of a neutron star, the densest objects accessible to direct observational scrutiny."
    },
    "0403/astro-ph0403121_arXiv.txt": {
        "abstract": "We analyze high resolution spectra of a multi--cloud weak [defined as $W_r({\\MgII}) < 0.3$~{\\AA}] absorbing system along the line of sight to PG~$1634+706$.  This system gives rise to a partial Lyman limit break and absorption in {\\MgII}, {\\SiII}, {\\CII}, {\\SiIII}, {\\SiIV}, {\\CIV}, and {\\OVI}.  The lower ionization transitions arise in two kinematic subsystems with a separation of $\\simeq 150$~{\\kms}.  Each subsystem is resolved into several narrow components, having Doppler widths of $3$--$10$~{\\kms}.  For both subsystems, the {\\OVI} absorption arises in a separate higher ionization phase, in regions dominated by bulk motions in the range of $30$--$40$~{\\kms}.  The two {\\OVI} absorption profiles are kinematically offset by $\\simeq 50$~{\\kms} with respect to each of the two lower ionization subsystem.  In the stronger subsystem, the {\\SiIII} absorption is strong with a distinctive, smooth profile shape and may partially arise in shock heated gas. Moreover, the kinematic substructure of {\\SiIV} traces that of the lower ionization {\\MgII}, but may be offset by $\\simeq 3$~{\\kms}.  Based upon photoionization models, constrained by the partial Lyman limit break, we infer a low metallicity of $\\sim 0.03$ solar for the low ionization gas in both subsystems.  The broader {\\OVI} phases have a somewhat higher metallicity, and they are consistent with photoionization; the profiles are not broad enough to imply production of {\\OVI} through collisional ionization.  Various models, including outer disks, dwarf galaxies, and superwinds, are discussed to account for the phase structure, metallicity, and kinematics of this absorption system.  We favor an interpretation in which the two subsystems are produced by condensed clouds far out in the opposite extremes of a multi--layer dwarf galaxy superwind. ",
        "introduction": "\\label{sec:intro}   Quasar absorption line systems, as traced by their {\\MgII}, provide a unique way to study our universe.  Current classifications schemes for systems with detected {\\MgII} file them into three main categories: those that are characterized by their strong {\\MgII} absorption, those that reveal weak, narrow, single--cloud {\\MgII} profiles, and those that exhibit weak, multiple cloud {\\MgII} absorption.  In this paper we analyze in detail the absorption profiles from a particular multiple cloud weak {\\MgII} absorber at $z=1.04$ along the line of sight toward PG~$1634+706$.  The longer--term goal of a collection of similar studies will be to understand the relationships between the different classes of absorption systems, and the connections with the different types of galaxies and structures at various redshifts.  Although there are no distinct divisions between the three categories of {\\MgII} absorption systems, by their properties they appear to be related to three different types of gaseous structures at $z\\sim1$.  Strong {\\MgII} absorbers [those with $W_r(2796)>0.3$~{\\AA}] show Lyman limit breaks and contain multiple clouds spread over tens to hundreds of kilometers per second \\citep{archiveI}.  A large majority of these absorbers are known to be associated with luminous galaxies ($>0.05 L^*$, where $L^*$ is the Schechter luminosity), within an impact parameter of $40h^{-1}$~kpc of the quasar \\citep{bb91,bergeron92,lebrun93,sdp94,steidel95,3c336}. All of the strong {\\MgII} absorbers ($dN/dz = 0.91 \\pm 0.1$ at $<z>=0.9$ \\citep{ss92}) can be accounted for by regions of this size around the known population of luminous galaxies.  The strong {\\MgII} absorber spectral profiles, in chemical transitions of low and high ionization states, generally require multiple phases of gas, i.e. regions of differing densities that are spatially distinct (e.g. \\citep{ding1634,ding1206}).  The kinematics of the low ionization gas is generally consistent with what one would expect from the combined disks and halos of galaxies of a variety of morphological types \\citep{kinmod,steidel02}.  In contrast to the strong {\\MgII} absorbers, the single--component systems with weaker {\\MgII} (rest frame equivalent widths $W(2796) <0.3$~{\\AA}) are typically not known to be directly associated with luminous galaxies (i.e. they are not within $50$--$100$~kpc of $>0.05 L^*$ galaxies) \\citep{weak1}.  These single--component weak {\\MgII} systems have a significant absorption cross-section, with $dN/dz = 1.10 \\pm0.06$ for $0.02< W< 0.3$~{\\AA} at $0.4<z<1.0$ \\citep{weak1}, similar to the value for all strong {\\MgII} absorbers.  The distribution of the number of Voigt profile components used to fit {\\MgII} absorption systems is roughly Gaussian with a median of seven components.  However, there is a strong excess of single cloud components relative to the rest of the distribution (see Fig.~2 in \\citet{weak2} and Table 7 of \\citet{cv01}). This implies that the single--component weak {\\MgII} absorbers are primarily produced by a different class of object than the strong {\\MgII} absorbers.  Consistent with this interpretation, most of the single--component weak {\\MgII} absorbers do not produce Lyman limit breaks \\citep{archiveI}.  Their {\\MgII} profiles are generally quite narrow, with Doppler parameters of $b\\sim 2$ -- $5$~{\\kms}.  Their {\\CIV} profiles are broader and require a separate, higher ionization, phase of gas.  Their low ionization phases are inferred to have metallicities greater than a tenth solar and to have sizes less than $\\sim 10$--$100$~pc \\citep{weak2,weak1634}.  The general properties of the single--cloud weak {\\MgII} absorbers suggest an origin in some type of faded early extragalactic star cluster or in metal--rich fragments in cold dark matter mini--halos ''failed galaxies'' \\citep{weak2}.  The third category of {\\MgII} systems is characterized by multiple, weak components spread over tens of kilometers per second. About one third of the weak {\\MgII} absorbers are in this category \\citep{weak1}. By definition, these are ``weak {\\MgII} absorbers'', since the total rest frame equivalent width of {\\MgII}~2796 is $<0.3$~{\\AA}.  Unlike the the strong {\\MgII} absorbers, little is directly known about the environments of the multiple--component weak {\\MgII} systems.  However, the multiple--component weak category appears consistent with an extension of the distribution of the numbers of components fit to strong {\\MgII} absorbers (see Fig.~2 in \\citet{weak2}).  The weakness and kinematics of the multiple--component weak {\\MgII} profiles might suggest a line of sight through the outer regions of a galaxy where the gas could be more diffuse and less dense.  It is also possible that a dwarf galaxy, or even a low--metallicity giant galaxy would give rise to such weak {\\MgII} absorption.  In this paper, we study a multiple--component, weak {\\MgII} [rest frame equivalent width, $W_r(2796) =0.097 \\pm 0.008$~{\\AA}] absorber at $z\\sim1$. This absorber was modeled previously by \\citet{anticip}, using low--resolution FOS/HST data along with the high--resolution {\\MgII} profiles from HIRES/Keck.  They determined that the {\\MgII} clouds have a relatively low metallicity of $0.03$ solar, based on a comparison to a partial Lyman limit break.  The {\\CIV}, detected in the FOS spectrum, was clearly offset from the low ionization clouds.  This required an additional high ionization cloud producing absorption $\\sim 200$~{\\kms} to the blue of the detected {\\MgII} clouds.  \\citet{anticip} predicted that this cloud should produce observable {\\CII}, {\\SiIII}, and {\\SiIV} in the high resolution STIS spectrum.  However, with just the low--resolution UV spectrum, they were unable to determine if the {\\CIV} detected at the same velocity as {\\MgII} could arise in the same phase of gas with the {\\MgII}, or if a separate phase would be needed.  Also, the ionization conditions and metallicities of the high ionization clouds could not be well--constrained.  STIS/HST spectra of PG~$1634+706$ are now available, with high resolution coverage of {\\Lya} and the Lyman series, {\\SiII}, {\\SiIII}, {\\CII}, {\\SiIV}, and {\\OVI}.  With these new data, we aim to test the models of \\citet{anticip}, and to reach more detailed conclusions about the phase structure, metallicity, ionization conditions and kinematic properties of the $z=1.04$ system toward PG~$1634+706$.  In \\S~\\ref{sec:data} we describe the spectra of quasar PG~$1634+706$ used in our analysis, obtained with STIS/HST, FOS/HST and HIRES/Keck.  We then present the data for the $z=1.04$ system along this line of sight in \\S~\\ref{sec:system} and make a qualitative comparison of the individual chemical transitions.  Our analysis method, considering photoionization and collisional ionization models, is outlined in \\S~\\ref{sec:method}.  In \\S~\\ref{sec:results} we present results from modeling of this system and we summarize the conclusions in \\S~\\ref{sec:summary}. In \\S~\\ref{sec:discussion}, we discuss possible scenarios for the origin of the $z=1.04$ system and the general implications for the nature of multiple--component, weak {\\MgII} absorbers. ",
        "conclusions": "\\label{sec:summary}  In the previous section, we gave a detailed description of how we constrained physical conditions for the multi-cloud, weak {\\MgII} absorber at $z=1.04$ along the line of sight to PG~$1634+706$. Here, we summarize those results, present an overview of a consistent model for the absorber, and compare to the results of \\citet{anticip}, which were based on lower resolution spectra:  \\begin{enumerate}  \\item{  Two subsystems, A and B, detected in low ionization transitions, are separated by $\\sim150$~{\\kms}.  Assuming, for simplicity, that each cloud in the subsystem has the same metallicity, then the observed partial Lyman limit break is consistent with a metallicity of $\\log Z\\sim-1.4$.  Models with higher metallicities and dust depletion are not viable because carbon and silicon are depleted less than magnesium, and they are already overproduced by models with a solar abundance pattern. Subsystem A has clouds with an ionization parameter of $\\log U\\sim -2.3$, however, the {\\MgII} could either be produced in those clouds (Model 1) or in additional lower ionization clouds, $\\log U\\sim -3.1$ (Model 2).  The ionization parameter of the clouds in subsystem B are slightly lower, $\\log U\\sim -2.7$.}  \\item{ The issue of whether subsystem A, centered at $v=0$~{\\kms}, has such an additional low ionization phase hinges partially on the question of whether the {\\MgII} and {\\SiIV} profiles are kinematically aligned. Figure~\\ref{fig:offset} shows that, although the shape of the {\\SiIV}~$\\lambda 1394$ absorption profile is similar to the {\\MgII} profile in subsystem A, it appears to be offset to the red by a few kilometers per second. An offset of the same magnitude was observed for one of the clouds in the $z\\sim1.9$ absorber along the line of sight to the HDF-South quasar, where {\\CIV}, {\\SiIII}, and {\\SiIV} are offset from {\\CII} and {\\SiII} by $\\sim4$~{\\kms} \\citep{DOdorico}. \\citet{DOdorico} interpret this offset as the consequence of an {\\HII} region flow.  A similar offset between {\\SiII} versus {\\CII} and {\\SiIII} may exist in the $z=0.00530$ absorber toward 3C~273 \\citep{tripp3c273}.  The interpretation of the $z\\sim1.04$ system is more complex because several clouds would be offset, but a similar explanation is possible.}  \\item{ The cloud sizes may also provide insight into whether a one (Model 1) or two (Model 2) phase model of subsystem A is more plausible.  If the {\\MgII} and {\\SiIV} arise in the same phase, the four components of subsystem A have sizes ranging from $6$--$24$~kpc.  It is difficult to conceive of a physical picture in which four such large structures would lie along the line of sight, within $50$~{\\kms}.  For a two phase model (with the offset {\\SiIV} components), the cloud sizes of the components responsible for the {\\SiIV} are reduced to several kpcs, with smaller ($100$--$300$~pc) {\\MgII} clouds also contributing to the absorption. This does seem more plausible in relationship to basic features of galactic interstellar gas.}  \\item{ Although both the offset {\\SiIV} clouds and the clouds centered on {\\MgII} contribute to the {\\SiIII} absorption profile, they do not account for its smooth, broad shape.  This may indicate that much of the {\\SiIII} is produced by collisional ionization at a temperature of $\\log T\\sim4.6$.}  \\item{ Two broad {\\OVI} absorption profiles are observed, one $50$~{\\kms} to the red of subsystem A and the other $50$~{\\kms} to the blue of subsystem B.  Assuming the {\\OVI} and {\\Lya} arise in the same phase, the measured Doppler parameter of {\\Lya} associated with the {\\OVI} places an upper limit on the temperature such that the {\\OVI} could not be produced by collisional ionization.  Both the {\\OVI} and the {\\Lya}, which could not be fully produced by the subsystems, can arise in the same photoionized phase.  The low resolution {\\CIV} profile, observed by FOS/HST, also requires a contribution from the blue broad component.  These highly--ionized clouds ($-1.5 < \\log U < -0.5$) have metallicities larger than the low--ionization subsystems. The model cloud sizes range from 10's to 100's of kpc, depending on the metallicity, with larger metallicities leading to smaller sizes. }  \\end{enumerate}  \\subsection{Comparison with Previous Predictions} \\label{sec:compareprev}  \\citet{anticip} studied the $z=1.04$ system before the STIS/HST spectrum was released, using the low resolution FOS/HST spectra in conjunction with the same HIRES/Keck spectra as were analyzed here. Now, using the higher resolution STIS/HST UV data, we confirm most of the basic results of that paper, and we are able to resolve some ambiguities that were presented there. The metallicity constraint is confirmed, along with the presence of an additional high ionization cloud $\\sim 200$~{\\kms} to the blue (corresponding to the blue broad component in the present models).  The metallicity agreement is not surprising, since in both cases the constraint came from HIRES/Keck spectra covering {\\MgII} in conjunction with the Lyman limit break whose appearance is not sensitive to resolution. The presence of an additional high ionization component is convincingly confirmed by the presence of {\\OVI} absorption at a consistent velocity.  It seems likely that the {\\CIV} seen in the FOS spectrum and the {\\OVI} seen in the STIS spectrum would arise in the same phase.  Because of the limitations of the low resolution data, the \\citet{anticip} work was unable to distinguish between two scenarios for the origin of the {\\CIV} in the redward part of the profile.  Based on those data, it could have arisen either in the same clouds with the {\\MgII} or in a separate phase.  Now, because the {\\SiIV} profile has narrow components resolved in the high resolution STIS data, we can see that the {\\CIV} arises either in the {\\MgII} clouds (Model 1) or in the separate {\\SiIV} clouds (Model 2.)  It does not arise in the red broad component which is responsible for producing {\\OVI} and the red wing of {\\Lya}.  The metallicities and ionization conditions of the blue broad component were poorly constrained by \\citet{anticip}, because only {\\Lya} and {\\CIV} were available, and at low resolution. They found that a model with $\\log Z = -1.5$ and $\\log U = -2.1$ could be consistent with the data, but higher metallicities were also compatible, depending on the $b$ parameter assumed for the unresolved broad component.  They indicated that low ionization transitions might be observed at this same velocity. Now, we find that the detection of strong {\\OVI} indicates a higher ionization parameter and a higher metallicity.  Also, although low ionization transitions are detected in the STIS spectrum, their absorption profiles are offset by $\\sim50$~{\\kms} from the blue broad component. The previous work therefore had the general idea that there was absorption blueward of the main low ionization system, but was able to discern little about the details.    From the derived physical constraints of the $z=1.04$ system along the line of sight to PG~$1634+706$, we determined certain basic properties that this absorber must have.  There are two separate subsystems, implying that there are two distinct structures.  These subsystems have lower metallicity, lower ionization regions, and kinematically distinct, higher metallicity and higher ionization ``halos''. Although these conclusions do not indicate a unique solution, they suggest scenarios and types of structures that might be responsible.   {\\it The subsystems of low-ionization material in the $z=1.04$ system have a low metallicity ($\\log Z\\sim -1.4$)---} Although giant galaxies would be expected to have lower metallicities at higher redshifts than they do at present, a value as low as $\\log Z\\sim -1.4$ would be more extreme than our expectations.  Almost all of the stars in the disk of the Milky Way formed well before redshift one (coincident with the peak of the cosmic star formation), and nearly all of them have metallicities greater than $0.1$~solar, with values as large as solar being common \\citep{wyse}. It is consistent that the metallicities of the low ionization clouds composing strong {\\MgII} absorbers at $z\\sim1$ appear to have higher metallicities, from 1/10th solar up to the solar value \\citep{ding1634,ding1206}.  Low metallicities at $z\\sim1$ are seen in damped {\\Lya} absorbers \\citep{pettini}, and this is attributed to their hosts having a mix of morphologies, commonly including dwarfs and low surface brightness galaxies \\citep{rao,bowendla}.  In Local Group dwarf galaxies at present, metallicities range from $\\log Z=-1.0$ to $-2.0$ \\citep{mateo}.  The outer disks of spiral galaxies could also have low metallicities consistent with our derived value for the $z=1.04$ system.  The iron radial metallicity gradient in the Milky Way disk is $\\sim -0.06$~dex~kpc$^{-1}$, out to a radius of $\\sim 16$~kpc \\citep{metgrad}.  We would expect weak {\\MgII} absorption (as opposed to strong) to arise only at galactocentric radii greater than $30$--$40$~kpc. An extrapolation of the observed metallicity gradient would therefore produce a low metallicity for radii at which weak {\\MgII} absorption might be expected.   {\\it The broad, higher ionization components in the $z=1.04$ system have a higher metallicity than the low ionization subsystems---} The subsystem absorption is produced in higher density regions than the broad higher ionization components.  Also, the individual subsystem clouds are narrower and smaller (sizes of $\\sim 1$~kpc) than the high ionization clouds (sizes of tens to hundreds of kpcs).  This suggests that the subsystems arise in a smaller region internal to a galaxy or structure, more likely to be close to the stellar sources of metals.  If so, at face value, it is puzzling why the more diffuse, highly ionized gas would have a higher metallicity than the concentrated regions.  One possibility would be a scenario in which relatively undiluted supernova ejecta are ejected from the central regions of a galaxy into a diffuse medium.  This is seen to some extent in the ejecta of nearby galactic winds, observed with Chandra, although the metallicities of the wind material in these local systems are only somewhat larger than those of the interstellar material from which they originate \\citep{martin1569,martinproc}.  Another clue may result from the abundance pattern of the blue broad component. If we tune its ionization parameter and metallicity so as to fit the {\\CIV} and {\\OVI} absorption (along with {\\Lya}), the {\\NV} absorption is severely overproduced.  A reduction of the abundance of nitrogen relative to carbon and oxygen would bring the model in accord with the data.  Such a nitrogen deficiency (about an order of magnitude) is characteristic of the dwarf galaxies in the Local Group \\citep{mateo}. This pattern is consistent with primary production of nitrogen in a low metallicity environment \\citep{henry}.  {\\it Each of the two subsystems of the $z=1.04$ absorber has several narrow low ionization components spread over tens of {\\kms}, which may be related to a broad, high ionization component offset by $50$~{\\kms}---} The kinematic spread of a subsystem is reminiscent of that we would expect for a line of sight through a spiral disk with some increase due to the effects of interactions (eg. contribution from warps, tidal debris, or fountain material). Strong {\\MgII} absorbers often show subsystems with kinematic spreads of tens of {\\kms}, though they have larger equivalent width and they also often have ``satellite clouds'' at larger velocities which contribute to a larger kinematic spread for the overall system \\citep{cv01}.  In the case of the strong {\\MgII} absorbers, the kinematics of the strongest subsystem have been demonstrated to be consistent with the expectations of thick disk kinematics \\citep{kinmod}. A kinematic spread of tens of {\\kms} is also similar to the rotation velocities of the larger Local Group dwarf irregulars and to the velocity dispersions of Local Group dwarf spheroidals \\citep{mateo}.  Significant offsets between the low and high ionization gas are not uncommon in the nearby universe.  For example, offsets are typically $\\sim 30$~{\\kms} between the {\\OVI} absorption detected in the Large Magellanic cloud and its low ionization absorption, which is associated with its interstellar medium \\citep{howk}.  The LMC {\\OVI} is consistent with an origin in coronal gas, being similar in strength and velocity dispersion to that associated with the Milky Way \\citep{howk,savageo6}. The observed kinematics for each of the $z=1.04$ subsystems could be due to a low ionization contribution from clumps in the disk or internal region of a dwarf, which would be surrounded by a broad {\\OVI} phase corona or halo.  The measured {\\OVI} Doppler parameter of $b\\sim20$--$30$~{\\kms} is consistent with a dwarf galaxy halo dispersion \\citep{mateo}.  The offset of $\\sim50$~{\\kms} of the broad {\\OVI} component from the low--ionization gas could be due to a kinematic offset between central/disk and halo components of the dwarfs.  Another possibility for the kinematic difference between the low and high ionization components in the $z=1.04$ absorber is a layer structure of the cone surrounding a breakthrough superbubble or wind. In the case of a starbursting dwarf, hot, X--ray emitting gas is ejected at high velocity from the central region of the galaxy (see, eg. \\citet{read, ptak, dahlem}).  The edges of the cone surrounding the X-ray gas are cooler and are likely to have condensations that would give rise to lower ionization absorption.  In fact, strong {\\NaI} absorption is detected in absorption against the majority of nearby starbursts \\citep{heckman00}. We would expect ``transition layers'' in between the hot and cooler regions. In these regions, moderate to higher ionization absorption could occur.  The broad {\\OVI} absorption observed in the $z=1.04$ system could trace such a transition layer in the interior.  As the line of sight passes through the outer layers, one might expect to see a group of clouds producing lower-ionization absorption that would be moving at a speed lower than that of the interior of the wind.  {\\it The low ionization absorption in the two $z=1.04$ absorber subsystems is quite weak.  Each of them would be classified as a multiple cloud, weak {\\MgII} absorber---} Because of the weak, low ionization absorption, statistically, we would infer that neither of the two subsystems would be produced by a line of sight through the inner $30$--$40$~kpc of a giant galaxy.  Unfortunately, in the case of PG~$1634+706$ we have no direct information from imaging of the field (as described in \\S~\\ref{sec:data}).  We can gain insight by considering the nature of the host galaxies or structures for nearby weak {\\MgII} absorbers, for which imaging can reach down reasonably far on the dwarf galaxy luminosity function. \\citet{rosenberg} have proposed that the weak low ionization absorber at $z=0.00530$ toward 3C~273 is produced in a thin shell around a $M_B = -14.5$ ``post--starburst'' galaxy that has been found at an impact parameter of $71 h_{70}^{-1}$~kpc.  This is consistent with \\citet{tripp3c273}'s speculations about that same system, and with the theoretical ideas of \\citet{theuns} concerning the formation of condensations in dwarf galaxy winds.  \\citet{rosenberg} have proposed that this is likely to be a common origin of weak, low ionization absorption systems.  There is another weak, low ionization absorber found at $z\\sim 0.005$ along the RXJ~$1230.8+0115$ sightline, which is adjacent to the 3C~273 sightline, at a separation of $\\sim 350h_{70}^{-1}$~kpc \\citep{rosenberg}.  However, in that case, the galaxy responsible has not yet been identified.  There is a candidate host dwarf galaxy known for the multiple--cloud, low ionization absorber at $z=0.051$ toward PG~$1211+143$, but its star formation history is unknown (2003, J. Stocke, private communication)  {\\it The $z=1.04$ absorber has a possible offset between the {\\MgII} and {\\SiIV} in subsystem A, and the {\\SiIII} in that same subsystem implies the presence of a collisionally ionized layer at intermediate temperature---}  Offsets between low to intermediate ionization transitions have been seen in other absorption systems, as was noted in \\S~\\ref{sec:summary}.  The interpretation could be a layered structure, with different kinematics for the layers responsible for the different ionization stages.  In the case of subsystem A, the offsets in multiple components would require a series of such structures with a common orientation.  The presence of a collisionally ionized layer at an intermediate temperature ($T\\sim100,000$~K) has also been suggested for other $z\\sim1$ absorption systems \\citep{ding1634,ding1206}.  This may also be related to shocks and is consistent with a complex, layered structure.  {\\it Two relatively strong, broad {\\OVI} absorption features are present in the $z=1.04$ absorber, each with a $50$~{\\kms} offset from one of the subsystems---}  \\citet{tripp1} have shown that most of the baryons in the universe are trapped in a warm/hot phase traced by {\\OVI}.  This has been interpreted as intragroup gas, however, the kinematic spread of the {\\OVI} in some of the absorbers is not as large as would be expected for a poor galaxy group \\citep{tripp2,zabludoff}.  It is also possible that such {\\OVI} arises in the gaseous halo of a single galaxy \\citep{tripp2}. Some of the stronger {\\OVI} systems have associated weak lower ionization absorption \\citep{savage02}.  It follows that there is some overlap between the stronger {\\OVI} absorbers seen by \\citet{tripp1} and multiple cloud weak {\\MgII} absorbers such as the one studied here.  More generally, the kinematic spreads of the broad {\\OVI} components in the $z=1.04$ absorber are consistent with the velocity dispersions of the larger Local Group dwarf galaxies \\citep{mateo}.  They are also consistent with the kinematics of the coronal gas around the Milky Way and the LMC \\citep{howk,savageo6}, as mentioned above. The {\\OVI} would also arise naturally in a wind model, somewhere between the central X--ray emitting gas and the cone in which the low ionization absorption would arise.  {\\it The two subsystems of the $z=1.04$ absorber are separated by $\\sim 150$~{\\kms}, and they appear to have fairly similar physical conditions.}  So far, we have discussed possible origins of the individual subsystems in the $z=1.04$ absorber (and their possibly associated high ionization broad components), but it is important to consider how two such systems could arise.  The two subsystems are kinematically distinct, i.e. they are well separated in velocity space. There are many scenarios that could give rise to absorption components separated by $\\sim 150$~{\\kms}, but most can be categorized into one of two themes.  In the first theme, there are two separate structures along the line of sight that exist within the same group or halo. In the second theme, a superbubble or wind is responsible for producing two kinematic structures that arise in the same object.  To evaluate the idea that two separate structures are responsible for the two subsystems, we must consider the likelihood of intercepting two structures along the same line of sight. The velocity separation of $150$~{\\kms} between subsystems A and B is characteristic of galaxy group velocity dispersions \\citep{zabludoff}. The responsible structures might be outer disks of galaxies or dwarf galaxies, as discussed above.  There are many ways of considering this problem.  For example, if there are typically ten satellite galaxies clustered within $200$~kpc around a giant galaxy at $z\\sim 1$, and if each satellite has a cross section for weak low ionization absorption of $1200$~kpc$^2$, there would be a $\\sim 10~\\%$ chance of a line of sight intercepting one if it intercepted another.  In this scheme, we would expect that most multiple--cloud weak {\\MgII} absorbers would not have two separate subsystems, i.e. we would expect that typically only one satellite would be intercepted.  This seems consistent with the data since, in a sample of $\\sim 10$ multiple cloud, weak {\\MgII} absorbers, only in this one case are there two separate subsystems \\citep{weak1}.  If a superwind or bubble is responsible for the $z=1.04$ absorber, we would expect that each of the two low-ionization subsystems and its associated {\\OVI} phase would represent an opposite side of a wind.  The two sides could plausibly be separated by $v\\sim150$~{\\kms} \\citep{heckman00}.  There would have to be cold condensations with a fairly large covering factor on both sides of the wind in order to see two low ionization subsystems.  One could easily imagine other cases with asymmetric blowouts or stochastic structure such that only one side was detected or such that one side was only detected in high ionization transitions.  The very strongest {\\MgII} absorbers (which are common only at $z>1$), could also have an origin in superwinds \\citep{bond}. These absorbers have a characteristic saturated ``double--trough'' absorption profile in {\\MgII}, which breaks up into multiple clouds in other weaker low-ionization transitions.  As with our two subsystems, the saturated troughs are typically separated by $\\sim100$--$200$~{\\kms}. Strong low-ionization absorption could arise from clouds close to the point of origin of the wind or from early evolutionary stages of the superwind phenomenon.  As we reach larger distances or later stages when low--ionization superwind clouds may have fragmented or dispersed, we might expect weaker absorption profiles like the ones we are seeing in the $z=1.04$ absorber toward PG~$1634+706$.  If more multiple cloud, weak {\\MgII} absorbers with two subsystems are observed it should be possible to test between an origin in two structures and an origin in a wind or outflow.  If winds are responsible then we would expect to see broad, high ionization absorption to the blue of the blueward subsystem and to the red of the redward subsystem. If two separate structures were responsible, the offset low and high ionization components would be due to kinematic differences between two components in such structures.  In such a case, the broad, high ionization absorption should arise at random to the red or to the blue of each of the subsystems, and sometimes would even be superimposed.  {\\it There are about two thirds as many multiple cloud weak {\\MgII} absorbers as there are strong {\\MgII} absorbers---} There must be a set of structures responsible for multiple cloud weak {\\MgII} absorption with a cross section comparable to the set of $\\sim30$~kpc regions around all $L^*$ galaxies. If these are outer disks around all $>0.05L^*$ galaxies, an annulus with thickness of $\\sim 10$~kpc would be required.  \\citet{bond} found that it was plausible for the expected number of superwinds at redshift $z\\sim1.5$ to account for the observed number of saturated ``double -- trough'' strong {\\MgII} profiles.  Similarly, we speculate whether the outer winds of a plausible number of bursting dwarfs could account for the observed number of multiple-cloud weak {\\MgII} absorbers.  There are about 2/3 as many multiple cloud weak {\\MgII} absorbers are there are strong {\\MgII} absorbers \\citep{weak1,weak2}.  Strong {\\MgII} absorption arises from regions around $>0.05 L^*$ galaxies with sizes of $\\sim30h^{-1}$~kpc \\citep{bb91,bergeron92,lebrun93,sdp94,steidel95,3c336}.  If the regions around starbursting dwarfs that give rise to multi-cloud weak {\\MgII} absorption have sizes of $10h^{-1}$~kpc, then a number density $6$ times that of $\\sim L^*$ galaxies could account for all these multi-cloud weak {\\MgII} absorbers.   Realistically, it seems likely that there would be contributions to multiple cloud weak {\\MgII} absorption from a variety of types of structures and processes. Conversely, all types of structures that contain gas must all make some contribution to the absorber population, contributing either strong or weak {\\MgII} absorption and/or contributing to {\\CIV} or {\\OVI} absorption. It is fairly well understood that $L^*$ galaxies produce most all of the observed strong {\\MgII} absorption at $z\\sim1$.  But it is not well understood what type of absorption systems quiescent dwarf galaxies or dwarf galaxies with winds would produce.  Above, we have considered the various properties of the $z=1.04$ absorber and discussed what types of structures and processes might be consistent. Clearly, there is no unique interpretation.  However, a model in which this absorber is related to dwarf galaxy winds is consistent with most of the listed properties.  It can explain the phase structure of the absorber, the relative metallicities of the phases, and the kinematics of the two subsystems relative to the broad, higher ionization components.  This model is also appealing in that it accounts for some of the absorption cross section presented by starbursting dwarfs, which were common at redshift one.  Many more multiple cloud, weak {\\MgII} absorbers must be studied in order to evaluate the relative contributions of dwarf galaxy winds, of quiescent dwarf galaxies, of outer galaxy disks, and of other phenomena.  In conjunction with such a statistical study, it will be essential to find analogs in the local universe so that the host galaxies and responsible processes can be directly identified."
    },
    "0403/hep-th0403270_arXiv.txt": {
        "abstract": "We consider warped compactifications in $(4+d)$-dimensional theories, with four dimensional de Sitter $dS_4$ vacua (with Hubble parameter H) and with a compact internal space. After introducing a gauge-invariant formalism for the generic metric perturbations of these backgrounds, we focus on modes which are scalar with respect to $dS_4$.  The physical eigenmasses of these modes acquire a large universal tachyonic contribution $-12d/(d+2) H^2$, independently of the stabilization mechanism for the compact space, in addition to the usual KK masses, which instead encode the effects of the stabilization. General arguments, as well as specific examples, lead us to conjecture that, for sufficiently large dS curvature, the compactified geometry becomes gravitationally unstable due to the tachyonic growth of the scalar perturbations.  This mean that for any stabilization mechanism the curvature of the dS geometry cannot exceed some critical value.  We relate this effect to the anisotropy of the bulk geometry and suggest the end points of the instability. Of relevance for inflationary cosmology, the perturbations of the bulk metric inevitably induce a new modulus field, which describes the conformal fluctuations of the 4 dimensional metric.  If this mode is light during inflation, the induced conformal fluctuations will be amplified with a scale free spectrum and with an amplitude which is disentangled from the standard result of slow-roll inflation. The conformal 4d metric fluctuations give rise to a very generic realization of the mechanism of modulated cosmological fluctuations, related to spatial variation of couplings during (p)reheating after inflation. ",
        "introduction": "De Sitter or quasi-de Sitter 4d geometries describe the present day acceleration of the Universe as well as the inflationary expansion at very early times.  This stimulates significant interest towards the construction of 4 dimensional dS geometry in the context of fundamental string/M theory, which are formulated in 10/11 dimension, with a compact internal space \\cite{Mal,ds1,KKLT}.  Similarly, compactification to 4d dS geometry takes place in phenomenological braneworld models, where the inner space is periodic with orbifold branes at the edges \\cite{ds2}. Although most of the activity in this area has been enhanced by very recent progress (both on the observational and theoretical side) the issue of dimensional reduction to the outer cosmological space--time was popular since the 1980s, either in high dimensional supergravity theories or on phenomenological grounds. For example, see the collection of references on KK cosmology given in \\cite{ss}.  The geometries we discuss here have $d$ spatial dimensions wrapped on a compact manifold ${\\cal M}\\,$, in addition to the standard $(3+1)$ space--time.  Many of the mentioned examples are covered by the $(4+d)$ dimensional geometry with the metric \\begin{eqnarray}\\label{blocks} d s^2 = e^{2\\hat A(y)}\\left(-dt^2+e^{2Ht} d \\vec x^2\\right)+e^{-2\\hat A(y)}{\\hat g}_{a b} \\left( y \\right) d y^ad y^b \\,, \\end{eqnarray} where the outer space is a dS 4d metric with a Hubble parameter $H$, while ${\\hat g}_{a b}$ is the metric of the compact inner space with $d$ coordinates $y^a$. For generality, we also include the warp factor $\\hat A(y)$. In the following we use greek letters to describe the $(3+1)$ outer space--time coordinates ($\\mu=0,...,3$) while roman letters span the inner compact space coordinates only. Capitalized roman letters span all coordinates. With the re-definition of the inner space metric ${ g}_{a b} \\equiv e^{-4\\hat A(y)}{\\hat g}_{a b}$ and the warp factor $A(y) \\equiv e^{\\hat A(y)}$, we can re-write (\\ref{blocks}) in a conformally-factorized form \\begin{equation}\\label{conf} d s^2=A(y)^2 \\left(d s_4^2 +{ g}_{a b} d y^a \\, d y^b \\right) \\ , \\end{equation} where $ds_4^2$ is de Sitter metric. This form of the metric is more convenient for developing the formalism of metric perturbations around a warped $dS_4 \\times {\\cal M}\\,$ background, which we will present below.  >From the string theory/supergravity perspective, recent studies have concentrated on the compatibility of high $(4+d)$ dimensional geometries ($d=6$) with a 4d de Sitter geometry as the outer space--time, and on the stabilization of the internal space to $ dS_4$ \\cite{KKLT}. The bulk geometry, which is usually treated in the supergravity limit, requires a careful study of the $(4+d)$ dimensional (bulk) Einstein equations.  The progress in finding solutions is related to the identification of various possible sources for the bulk stress energy tensor $T^{A}_{B}$, including supergravity lagrangian fields, branes with fluxes, etc. In phenomenological braneworld models, the $(4+1)$ or $(4+2)$ dimensional bulk geometry is stabilized, for instance, by means of bulk scalar fields as in the Goldberger-Wise model~\\cite{GW}, or more generally through the Casimir effect \\cite{Cas,pp}.  Very often, however, the stabilization is studied at the level of a 4 dimensional effective theory, where the inner space geometry emerges in terms of moduli fields. If the effective potential of the moduli has a minimum, the moduli are considered to be stabilized. An example will be the stabilization mechanism with quantum field theory effects which emerge from the properties of ${\\cal M}\\,$, e.g. the Casimir effect from compact inner dimensions. It is a reasonable assumption that in principle there shall be a bulk prototype  $\\langle T^{A}_{B}\\rangle$ (may be as a complicated functional of the metric and the topology) which generates corresponding terms of the 4d effective potential. The version of  string theory dimensional reduction of~\\cite{KKLT} is realized with the participation of the instanton effects (which are manifested in 4dim superpotential of $N=1$ supergravity). An interesting question is whether it is possible to think about a bulk prototype  $T^{A}_{B}$  of these effects.   While it is reasonable to assume that the low energy 4d effective action is Einstein gravity plus moduli fields (or Brans-Dike gravity plus moduli), this picture lacks the full higher dimensional solution of the Einstein equations with proper sources. In particular, the four dimensional description is inadequate for studying the high energy regime (high dS curvature) where the stabilization can  break down. For example, in the 5 dimensional braneworld models, the issue of stabilization due to the bulk scalar field can be formulated fully in terms of the stability of the 5 dimensional warped geometry against scalar metric perturbations. From this study one can determine the modes which enter in the exact low energy 4 dimensional description, namely the radion and the other KK modes of the system \\cite{radion}, which cannot be found from an heuristic 4 dimensional effective potential.  We are not concerned here with the details of the stabilization mechanism.  Our main goal is to determine the effects of the dS expansion on the stability of these geometries.  As we will show, compactifications to a 4dim dS space--time are more difficult to achieve than compactification to a 4dim flat space--time. This is due to a tachyonic contribution to the square of the mass of scalar metric perturbations arising from the de Sitter curvature.  In what follows we will discuss the properties of $(4+d)$ dimensional classical Einstein equations assuming a bulk $T^{A}_{B}$ as a source, but without specifying it.  The gravitational instability which we will discuss comes from the gravity sector, so that the exact form of $T^{A}_{B}$ will not be crucial. Of course it will be attractive to check how the effect works for each particular model of $T^{A}_{B}$.  However, we will try to argue that the instability effect is generic.  Previous studies of the gravitational sector ({\\it i.e.} metric perturbations) of the compactification concentrated on the dimensional reduction to 4dim flat space--time \\cite{ss}.  The instability effect which we find emerges when the outer space is curved instead of flat, and it is proportional to the curvature. Notice that test scalar fields propagating in $dS_4 \\times {\\cal M}$ do not exhibit tachyonic masses under KK projection~\\cite{KS}. Thus, there is a significant difference in the mass spectrum between the test fields and the self-consistent treatment of the metric fluctuations around this background.  The plan of this paper is as follows. In Section ~\\ref{sec:metric} we introduce a generic formalism of metric fluctuations around the $dS_4 \\times {\\cal M}$ background and define scalar, vector and tensor fluctuations acording to their transformation properties relative to $dS_4$. In Section~\\ref{sec:lin} we introduce the linearized equations that govern the evolution of scalar perturbations of the metric and show how the tachyonic instability of the scalar modes arises from the gravitational sector. In Sections~\\ref{sec:ex1} and ~\\ref{sec:ex2} we show two explicit example where the instability is manifested. In Section~\\ref{sec:mod} we discuss how the instability may be relevant for the generation of the primordial perturbations, providing a very natural and model-independent realization of the idea of modulated perturbations~\\cite{modulated,dgz}. Conclusions and a brief discussion can be found in Section~\\ref{sec:end}. ",
        "conclusions": "~\\label{sec:end}  A typical denominator in many extensions of the Standard Model is the presece of extra-dimensions. In the simplest possibilities, the extra space is a compact and static manifold ${\\cal M}\\,$. A stabilization mechanism is typically required to explain why ${\\cal M}$ should remain static, while the $3$ noncompact spatial dimensions are undergoing cosmological expansion. In particular, significant activity has been recently made trying to reconcile this picture with a de Sitter (or quasi de Sitter) geometry for the noncompact coordinates. Indeed, observations are telling us that the expansion of the universe was accelerating at very early times, and it is accelerating also at present.  The present note is focused on the effects of the inflationary expanison on the stability of ${\\cal M}\\,$. We have argued that the dS curvature (in other words, a nonvanishing expansion rate $H$) has typically the effect to destabilize the internal space, and that any given stabilization mechanism can be effective only up to a certain curvature. We have aimed to discuss this effect in the most general way possible, enlighting the consequences of making the only assumption that the geometry is of the $dS \\times {\\cal M}_4$ type, with an arbitrary compact space of $d$ dimsensions, and possibly with the presence of a warp factor. Such a set-up is also relevant for string-theory, in the supergravity limit.  To discuss the stability of the system, we have studied the most general  set of perturbations of the $dS_4 \\times {\\cal M}_4$ geometry. It is very convenient to classify the perturbations into irreducible representations of the $dS_4$ symmetry group. The big advantage of doing so, is that modes belonging to different representations are not coupled at the linear level. The perturbations can be devided into scalar, vecor, or tensor modes with respect to the $dS_4$ isometries. After a general classification, we have focused on the scalar modes, since they are ususally the most relevant one for inflationary geometries, and since they encode the effects of the instability we want to discuss. The linearized Einstein equations for the perturbations show that the physical mass squares $m^2$ of these modes acquire a negative contribution due to the $dS$ expansion, \\begin{equation} m_n^2 \\left( H \\right) = - \\frac{12 \\, H^2}{1+d/2} \\,\\,. \\label{tacconc} \\end{equation} This is true for arbitrary ${\\cal M}\\,$, $d\\,$, and for any possible underlying stabilization mechanism.  The presence of this tern is a signal for a possible gravitational instability of the system. Indeed, if the whole $m_n^2$ turns out to be tachyonic, the $dS \\times {\\cal M}$ is unstable. Clearly, verifying whether (and at which $H$) the instability takes place is very model dependent issue which should be verified case by case (namely, for any given source $T^A_B \\,$). Indeed, other contributions to $m^2$ also dpend on $H$ (although in most cases only indirectly, due to the fact that these terms depend on other background quantities, and that these quantities are related to $H$ through the background Einstein equations). However, the contribution~(\\ref{tacconc}) is {\\it generally} present, and it has to be taken always into account. Moreover, the study of several different examples, as well as general arguments, lead us to conjecture that the term~(\\ref{tacconc}) is the dominant one at large $H\\,$, and that no stabilization mechanism can be effective up to arbitrarly large dS curvature.  In this work, we have presented explicit and exact calculations in two specific examples. The first of them is codimension one braneworld configurations with bulk scalar field(s). As shown in~\\cite{FK}, in this case it is possible to derive an explicit upper bound $m^2 \\leq - 4 \\, H^2 + \\mu_0 \\left( A \\right)\\,$ (where $A$ is the warp factor, see eq.~(\\ref{mass3})) for the mass of the lightest eigenmode. This bound allows to determine at which $H$ the system becomes unstable. The second example is when ${\\cal M}$ is a $d-$sphere, and when the only source term isw a cosmological constant. This system is well known to be unstable, and we have shown that the instability is precisely due to the tachyonic nature of the scalar modes which we have identified in the general calculation.  There are also general arguments in support of the instability at high $H\\,$. The first is related to causality. If, as in the standard $4$ dimensional case, the Hubble length $H^{-1}$ has the meaning of a casual horizon, we should expect that any stabilization mechanism cannot be at work when $H$ becomes much greater than the inverse size of ${\\cal M}$ (since different ``edges'' of ${\\cal M}$ would then be casually disconnected). Such a behavior can be found in~\\cite{lpps}, where the cosmology of the Randall-Sundrum model~\\cite{RS1} was studied under the assumption of radion stabilization. It was shown in~\\cite{lpps} that radion stabilization cannot be imposed if the physical energy on the hidden brane is greater than $\\sim {\\rm TeV^2} \\, M_p^2 \\,$. The appearence of this intermediate scale is somewhat surprising, and indeed it was left unexplained in~\\cite{lpps}. However, it can be verified that it is precisely at this energy density that $H^{-1}$ becomes smaller than the distance between the two branes.~\\footnote{We thank Lorenzo Sorbo for this observation.}  A second argument can be inferred from the $dS_4 \\times S_d$ example. We noted that the instability has two possible end-points. One is characterized by a shrinking internal space, with asymptotic Kasner solution (a similar behavior was already noted in~(\\cite{branecode}) in the case of brane collisions). The second one is instead exact $dS_{4+d}\\,$, with all the compact and noncompact dimensions expanding with the same asymprtotic rate. This suggests that the instability can be attributed to the tendency of an inflationary expansion to homogenize and isotropize the {\\it whole} geometry, in this case by ``dragging'' also the compact space into a de--Sitter solution. This behavios, which is a well known feature of inflationary expansions, is encountered also in many Bianchi models.    In the final part of the work, we have also commented on the possible role of the instability in the generation of primordial perturbations. The scalar perturbations appear as a conformal mode from the four dimensional point of view (for example, for an observer living on a $3$-brane, or after integrating out the internal space ${\\cal M}\\,$). This mode is a clear signature of the extra dimesnions, since it is absent in the standard $4$ dimensional case. If, due to the tachyonic contribution, this mode light ($0 < m < H$) during the de Sitter stage, it will be amplified with a scale invariant spectrum and with an amplitude which is disentangled from the standard result of slow-roll inflation. One can think of a number of ways how the conformal perturbation could be responsible for the adiabatic mode in the later FRW evolution. For instance, consider a mixture of fields conformally and nonconformally coupled to the noncompact geometry. Only the nonconformally coupled fields will be sensitive to the conformal perturbations, so that we expect that the final mode perturbation will be proportional to the amount of these fields in the mixture. In the text we have instead focused on a much simpler possibility, related to the mechanism of modulated perturbations~\\cite{modulated,dgz}. The conformal factor can be usually rescaled away by a rescaling of the energy scales in the $4$ dimensional theory (this is exactly what occurs in the Randall-Sundrum model~\\cite{RS1}). Hence, its fluctuations can be interpreted as fluctuations masses of particles and rates of processes in the $4$ dimensional theory. In oarticular, this may be the origin of the fluctuations of the decay rate of the inflaton, which is at the basis of the mechanism of modulated perturbations."
    },
    "0403/astro-ph0403647_arXiv.txt": {
        "abstract": "We describe  measurements of the mirror vignetting in the XMM-Newton Observatory made in-orbit, using observations of SNR G21.5-09 and SNR 3C58 with the EPIC imaging cameras. The instrument features that complicate these measurements are briefly described. We show the spatial and energy dependences of measured vignetting, outlining assumptions made in deriving the eventual agreement between simulation and measurement. Alternate methods to confirm these are described, including an assessment of source elongation with off-axis angle, the surface brightness distribution of the diffuse X-ray background, and the consistency of Coma cluster emission at different position angles. A synthesis of these measurements leads to a change in the  XMM calibration data base, for the optical axis of two of the three telescopes, by in excess of 1 arcminute. This has a small but measureable effect on the assumed spectral responses of the cameras for on-axis targets. ",
        "introduction": "\\label{sect:intro}  % XMM-Newton  \\cite{Jansen} comprises 3 co-aligned telescopes, each with effective area at 1.5keV of $\\sim$1500cm$^{2}$, and Full Width Half Maximum (FWHM) angular resolution of $\\sim$5 arcseconds. The 3 telescopes each have a focal plane CCD imaging spectrometer camera provided by the EPIC consortium. Two also have a reflection grating array, which splits off half the light, to provide simultaneous high resolution dispersive spectra. These two telescopes are equipped with EPIC MOS cameras \\cite{mos}, which are conventional CMOS CCD-based images enhanced for X-ray sensitivity. The third employs the EPIC PN camera \\cite{pn} which is based on a pn-junction multi-linear readout CCD.  The EPIC cameras offer a field of view (FOV) of $\\sim$30 arcminute diameter, and an energy resolution of typically 100~eV (FWHM) in the range $\\sim$0.2--10~keV. The two MOS telescopes are equipped with a Reflection Grating instrument \\cite{denHerder} that has its own dedicated readout camera.  The in-orbit calibration of the XMM-Newton mirrors has been reported elsewhere \\cite{Aschenbach}, with special reference to the on-axis angular resolution (Point Spread Function, PSF). A second important calibration data set that is critical for analyzing spectroscopic information is the energy-dependent effective area \\cite{Aschenbach2002}. Both these features are under constant review as a result of improving knowledge of the instrumentation, and the requirements imposed by new science investigations. In this work we concentrate on different aspects of mirror performance that must be calibrated in the context of other scientific drivers which include, for example, cluster radial brightness distribution for determining gas mass, exposure maps and counts-to-flux conversions in population studies and diffuse background normalization measurements etc.. The reduction in effective area with radial distance from the field of view centre, or vignetting, must be accurately determined to support these investigations.  To highlight the effect visually, Figures~\\ref{seren_image} and \\ref{seren_image2} show the excess flux per source detected in the 1XMM catalogue of EPIC serendipitous source detections (plotted in units of sigma). The images are displayed in the EPIC camera detector coordinates, and flux determinations assume the nominal vignetting correction centred on the reference pixel of the detector co-ordinate system (DETX, DETY in the nomenclature of the XMM data analysis system). These figures show that some low level discrepancy in the spatial variation in effective area calibration must be present.  \\begin{figure} \\caption{Image in the MOS-2 detector plane, of the mean difference in the total-band (0.2-12 keV) flux seen by MOS-2 and MOS-1 expressed in Sigma. Bright pixels indicate an excess of flux in MOS-2 and dark pixels an excess in MOS-1. } \\label{seren_image} \\begin{tabular}{c} \\includegraphics[bb = 20 100 573 740, clip,scale=0.65]{m2flux2_box.ps} \\end{tabular} \\end{figure}  \\begin{figure} \\caption{ The equivalent image for the Epic-pn detector plane, of the mean difference in the total-band (0.2-12 keV) flux seen by Epic-pn and MOS-1 expressed in Sigma. Bright pixels indicate an excess of flux in pn and dark pixels an excess in MOS-1. }  \\label{seren_image2} \\begin{tabular}{c} \\includegraphics[bb = 36 170 478 680, clip, scale=0.65]{pnflux2_box.ps} \\end{tabular} \\end{figure}  For XMM-Newton, direct measurement on the ground of the X-ray vignetting function was prevented because nearly all X-ray beam measurements were performed in a non-parallel beam. The installation of an X-ray stray-light baffle in front of the mirrors, and the Reflection Grating Array (RGA) stack at the mirror exit plane \\cite{denHerder}, introduced potential complications that were only measured in long wavelength, visible light at an EUV parallel beam facility \\cite{Tock}.  Although the measured {\\em geometric} vignetting factor at longer wavelengths was comparable with predictions, it was necessary to use in-orbit data to confirm the X-ray energy dependence, and check that the geometric factor was maintained through the spacecraft assembly, integration, verification and launch campaigns. ",
        "conclusions": "The energy dependent vignetting calibration can be well matched to pre-launch predictions, but only on an assumption that the telescope optical axis is not perfectly aligned with the telescope boresight. This is not unexpected following difficulties on-ground of maintaining and/or measuring the telescope axis to better than 10's arcseconds.  We note that the assumed telescope axis misalignment implies that ``on-axis'' targets at the common boresight location are actually at a slightly different vignetting value per telescope. We speculate that this partly accounts for some of the observed flux discrepancies between the MOS and pn cameras (Figure~\\ref{vigcal}). After reviewing these data sets, it was decided that the XMM-Newton calibration database would be updated in 2004 to account for new reference axes to be centred at PN (DETX,DETY = 1240,400 ), MOS1 (DETX,DETY = 100,-200 ) and MOS2  (DETX,DETY = 500,-1250 ).  \\begin{figure}[ht] \\begin{center} \\begin{tabular}{c} \\includegraphics[scale=0.5]{vigcal.ps} \\end{tabular} \\end{center} \\caption{ Ratio of uncorrected and vignetted effective areas for an on-axis target. The flux differences for typical 0.5-2keV band will be $\\sim$ a few \\% for typical objects, and a small difference in recovered spectral slope will be caused by the change with energy  }\\label{vigcal} \\end{figure}"
    },
    "0403/astro-ph0403192_arXiv.txt": {
        "abstract": "We present observations of a remarkable submillimetre-selected galaxy, SMM\\,J16359+6612.  This distant galaxy lies behind the core of a massive cluster of galaxies, A\\,2218, and is gravitationally lensed by the foreground cluster into three discrete images which were identified in deep submillimetre maps of the cluster core at both 450 and 850$\\mu$m.  Subsequent follow-up using deep optical and near-infrared images identify a faint counterpart to each of the three images, with similar red optical--near-infrared colours and {\\it HST} morphologies. By exploiting a detailed mass model for the cluster lens we estimate that the combined images of this galaxy are magnified by a factor of $\\sim$45, implying that this galaxy would have un-lensed magnitudes $K_s=22.9$ and $I=26.1$, and an un-lensed 850-$\\mu$um flux density of only 0.8\\,mJy.  Moreover, the highly constrained lens model predicted the redshift of SMM\\,J16359+6612 to be $z=2.6\\pm0.4$. We confirm this estimate using deep optical and near-infrared Keck spectroscopy, measuring a redshift of $z=2.516$.  SMM\\,J16359+6612 is the faintest submm-selected galaxy so far identified with a precise redshift. Thanks to the large gravitational magnification of this source, we identify three sub-components in this submm galaxy, which are also seen in the NIRSPEC data, arguing for either a strong dust (lane) absorption or a merger. Interestingly, there are two other highly-amplified galaxies at almost identical redshifts in this field (although neither is a strong submm emitter). The three galaxies lie within a $\\sim 100$\\,kpc region on the background sky, suggesting this submm galaxy is located in a dense high-redshift group. ",
        "introduction": "Recent submillimetre (submm) surveys show that the majority of the submm background at wavelengths of $\\sim 850 \\mu$m arises from a population of distant, highly luminous infrared galaxies with 850\\,$\\mu$m flux densities of $\\gs 1$\\,mJy (Blain et al.\\ 1999; Cowie et al.\\ 2002; Smail et al.\\ 2002).  The steep slope of the submm counts indicates that the integrated background is dominated by $\\sim 1$\\,mJy galaxies, below the $\\sim 2$\\,mJy confusion limit of the deepest blank-field 850$\\mu$m surveys carried out to date (Hughes et al.\\ 1998), and well below the 3--8\\,mJy flux density limits (3--4$\\sigma$) typically achieved in wider-field 850\\,$\\mu$m surveys (e.g.\\ Eales et al.\\ 1999; Scott et al.\\ 2002; Webb et al.\\ 2003).  Submm galaxies in the 1-mJy regime can, however, only be studied at present by exploiting a natural telescope -- a massive gravitational lens formed by a rich clusters of galaxies -- which provide the opportunity of overcoming both the confusion limit (by spatially magnifying an area of the sky behind the lens) and the sensitivity limit (by gravitational magnification of the sources within this area).  This strategy has successfully been used by Smail, Ivison \\& Blain (1997), Chapman et al.\\ (2002), Cowie et al.\\ (2002) and Smail et al.\\ (2002) to detect submm sources amplified by factors typically of 1.5--4$\\times$.  Unfortunately, the intrinsically faintest sources are usually only detected at modest significance, complicating the identification and analysis of their counterparts in other wavebands.  To investigate the properties of the mJy-population in detail we require a high-signal-to-noise detection of a sub-mJy submm galaxy. Such systems can arise in rare configurations where the source is located within a small region on the background sky defined by the caustic of the lens, which then forms multiple, highly magnified images of the background galaxy (see also Borys et al.\\ 2004).  This situation not only provides a tool for investigating the properties of intrinsically faint galaxies in detail, but also provides a unique opportunity to measure the redshift of the background galaxy from the geometry of the images combined with an accurate model of the cluster lens, using a ray-tracing technique to triangulate the three-dimensional position of the galaxy in the volume behind the lens. This technique, which was demonstrated successfully by Kneib et al.\\ (1996) and confirmed by Ebbels et al.\\ (1998), circumvents the need for bright optical or infrared counterparts to a submm galaxy prior to measuring its redshift. We may thus avoid a bias in redshift measurements towards unrepresentative, low redshift or less obscured systems (Smail et al.\\ 2002; Webb et al.\\ 2003).  This paper describes the discovery of a multiply-imaged submm galaxy, SMM\\,J16359+6612, which is gravitationally lensed by the core of the rich cluster A\\,2218, and appears as three distinct sources in our 850\\,$\\mu$m discovery map.  We identify faint counterparts to all three sources in deep optical and near-infrared (NIR) imaging, consistent with their identification as three images of a single background galaxy. We estimate its redshift using a detailed mass model for the cluster lens. We then confirm the accuracy of the redshift with spectroscopic observations and finally discuss the intrinsic properties of this system. Throughout we will assume an $\\Omega=0.3$, $\\Lambda=0.7$ cosmology with $H_0=70$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$. At a redshift of $z=2.516$ the angular scale is thus 8.06 kpc/arcsec. ",
        "conclusions": "\\label{sec.discussion}  The occurrence of a submm galaxy falling within the caustic lines of a massive foreground gravitational lens provides unique constraints on the properties of this faint submm galaxy.  Using the cluster mass model we estimate that the background galaxy is gravitationally amplified by a factor of $\\sim$45 (integrated across all three images), indicating that the intrinsic 850\\,$\\mu$m flux density of this galaxy would be 0.8\\,mJy in the absence of gravitational magnification, while in the optical and NIR the galaxy would have magnitudes $I_{814}=26.1$ and $K_s=22.9$.  The galaxy therefore represents a serendipitously positioned example of the submm galaxy population at flux levels of $\\sim 1$\\,mJy, the population which produces the bulk of the submm background (Blain et al.\\ 1999).  This provides a unique opportunity to compare the properties of this low-luminosity submm galaxy with those of more luminous, submm galaxies studied in brighter, blank-field surveys (e.g.\\ Ivison et al.\\ 2002; Chapman et al.\\ 2003).  A second, multiply-imaged submm galaxy has been identified by Borys et al.\\ (2004), giving two high-magnification examples from the 24 cluster lenses at $z>0.1$ mapped with sufficient sensitivity using SCUBA (Chapman et al.\\ 2002; Smail et al.\\ 2002; Knudsen et al.\\ in prep.).  The surface density of $\\geq 1$\\,mJy submm galaxies is around 3\\,arcmin$^{-2}$ (Smail et al.\\ 2002; Cowie et al.\\ 2002), suggesting that roughly ten clusters need to be surveyed to detect a strongly-lensed mJy-flux submm galaxy.  This is consistent with the expectation based on models for the cluster lenses of uniform $\\sim$800\\,km\\,s$^{-1}$ potential wells at $z\\sim 0.2$, each with a critical curve encircling 0.05\\,arcmin$^{-2}$, and background submm galaxies at $z\\sim 2.4$ (see Kraiberg et al 2004 for a thorough discussion).   In our assumed cosmology the far-infrared luminosity of SMM\\,J16359+6612 is $1.0\\times10^{12}$\\,L$_\\odot$, roughly 1.5$\\times$ fainter than Arp\\,220, close to the border-line between luminous infrared galaxies and ultraluminous infrared galaxies (ULIRGs). The star formation rate derived from the far-infrared luminosity is about 500\\,M$_\\odot$\\,yr$^{-1}$.  Figure~3 compares the restframe (lens-corrected) SED of the SMM\\,J16359+6612 to the SEDs of Arp\\,220, the well-studied $z=2.8$ submm galaxy SMM\\,J02399$-$0136 and the $z=1.44$ Extremely Red Object and submm galaxy HR10 (Dey et al.\\ 1999). The L$_{FIR}$/L$_{opt}$ ratio is higher than that seen for Arp\\,220, suggesting it is more obscured, and similar to the high-redshift, but much more luminous, SMM\\,J02399$-$0136 (Ivison et al.\\ 1998).  \\smallskip \\centerline{\\psfig{file=a2218_multi_sed_col5.ps,width=3.4in}} {\\small\\addtolength{\\baselineskip}{-1pt} \\noindent{\\sc Fig.~4.} --- The observed Spectral Energy Distribution (SED) of the multiply-imaged submm galaxy SMMJ\\,16359+6612 (solid points with error bars) compared to Arp\\,220 (solid line, composite spectrum from Anantharamaiah et al.\\ 2000, Klaas et al.\\ 1997, Lisenfeld et al.\\ 1996 and Surace et al.\\ 2000) which is 1.5$\\times$ more luminous in the far-infrared, as well as the well-studied submm galaxy SMMJ\\,02399$-$0136 (open diamonds, at $z=2.8$, Ivison et al.\\ 1998) which is 10$\\times$ more luminous in the far-infrared, and the Extremely Red object HR10 (open triangles, at $z=1.44$, Dey et al.\\ 1999) which is 7$\\times$ more luminous in the far-infrared. All SEDs are redshifted to match that of SMM\\,J16359+6612 and scaled so they have a similar far-infrared luminosity.  Notice the relatively large variation of L$_{FIR}$/L$_{opt}$ between the different galaxies.\\\\ \\label{fig.sed} }  At $z=2.515$ the spatial resolution of the {\\it HST} images, corresponds to $\\sim 0.1$\\,kpc, taking into account the gravitational magnification.  Thus the colour gradient within the background galaxy is on 1--2\\,kpc scales -- similar to the obscured region in nuclear starbursts in ultraluminous infrared galaxies at the present-day. A constraint on the size of the submm emission region is obtained by subtracting point sources from the 850 and 450\\,$\\mu$m maps. This procedure leaves no detectable residual of extended emission, indicating that to our measurement accuracy we are dealing with point sources. The tightest constraint is obtained at 450\\,$\\mu$m where we find that the intrinsic extent of the emission is less than $4''$, which corresponds to $\\sim 8$\\,kpc. There is thus no indication for a highly extended obscured starburst in this relatively low-luminosity submm galaxy.  The precise redshift we have measured for SMM\\,J16359+6612 enables us to identify that this galaxy lies at an identical redshift to that of \\#273 and the \\#384/\\#468 multiple image system.  Using the lens model we estimate that SMM\\,J16359+6612 lies between the \\#384/\\#468 and \\#273 galaxies in the source plane, and all three galaxies are less than 130\\,kpc apart. If these galaxies were not magnified by the cluster lens, all three galaxies would appear within a radius of 8$''$.  Moreover, we can place a limit of $\\ls 100$\\,km\\,s$^{-1}$ on the possible velocity offset between these three galaxies. They are thus all part of a single group and it is likely that they are interacting, which may explain the activity we detect in the submm waveband.  The strong clustering of the submm galaxies with UV-bright populations highlights the opportunity for measuring the distances to submm galaxies from the redshifts of less-obscured companions, as well as the possible confusion which may arise when trying to relate UV- and submm-selected populations, in the absence of precise positions from radio counterparts.  SMM\\,J16359+6612 has much redder colour in $(V-I)$ and $(I-K)$ than the two other nearby UV-selected galaxies, confirming the dustier nature of this galaxy.  If we compare the unlensed star formation rate (SFR) for this three systems, we find based on their UV continuum that: SMM\\,J16359+6612 has 6 M$_\\odot$\\,yr$^{-1}$ (although certainly underestimated due to dust extinction) \\#468: 14 M$_\\odot$\\,yr$^{-1}$ and \\#273: 4 M$_\\odot$\\,yr$^{-1}$.  For SMM\\,J16359+6612 we can compare the three estimates of its amplification-corrected star formation rate: 6 M$_\\odot$\\,yr$^{-1}$ from its dust-corrected UV luminosity, 11\\,M$_\\odot$\\,yr$^{-1}$ from the H$\\alpha$ flux (uncorrected for extinction or aperture losses) and 500\\,M$_\\odot$\\,yr$^{-1}$ based on the far-infrared emission.  Applying the median extinction correction derived for mid-IR selected luminous infrared galaxies at $z\\sim 0.7$ by Flores et al.\\ (2003), A$_{{\\rm H}\\alpha}\\sim 2.1$, and a modest correction for slit losses would increase the H$\\alpha$-derived star formation rate by a factor of $\\sim 10\\times$.  These results suggest the vast majority of the young stars in this galaxy are obscured by dust and are undetectable in the restframe far-UV.  Although it is 10$\\times$ less luminous than the typical submm-selected galaxy studied in blank-field surveys, this system shares the same extreme levels of obscuration seen in the more luminous galaxies, rather than the more modest dust obscuration inferred for the somewhat less luminous UV-selected galaxies at these redshifts. We also note that the H$\\alpha$ line width and velocity structure, if they reflect the dynamics of the galaxy, suggest that this system is more massive for its UV luminosity than typical UV-selected galaxies at this epoch (Erb et al.\\ 2003).  In summary, we have identified a multiply-imaged submm galaxy seen through the core of the rich cluster A\\,2218.  The cluster lens amplifies the background galaxy by a factor of $\\sim 45\\times$, providing a high signal-to-noise view of an example of the sub-mJy submm population which provides the bulk of the extragalactic background in this waveband.  We estimate a redshift for this galaxy from our highly-constrained lens models for the cluster and confirm this using optical and near-infrared spectroscopy from Keck.  The redshift of the submm galaxy is $z=2.516$, placing it at the same redshift as two other strongly-lensed UV-bright galaxies in this field.  Our Keck spectroscopy suggests that the emission from the submm galaxy is dominated by star formation and in contrast to the typical UV spectral properties of more luminous submm galaxies, this galaxy shows Ly-$\\alpha$ absorption, rather than emission (c.f.\\ Chapman et al.\\ 2003).  The star formation rates we derive from the UV continuum, H$\\alpha$ and far-infared emission show that most of the star formation activity is obscured and we suggest it is likely to be located in the component $\\gamma$, given its extreme $(I-K)$ colour. Because of its faint submm flux this galaxy is likely to be a good example of the type of galaxy that makes most contribution to the star formation history, thus deserving a detailled study with submm/mm interferometer to study the dynamics and mass of this obscured region. The presence of three highly-amplified $z=2.515$ galaxies in our survey field indicates that the submm galaxy resides in a compact group, interactions within which may help explaining the triggering of the obscured starburst we detect."
    },
    "0403/astro-ph0403471_arXiv.txt": {
        "abstract": "A plausible model is proposed for the enhancement of the abundance of molecular species in bipolar outflow sources.  In this model, levels of HCO$^+$ enhancement are considered based on previous chemical calculations, that are assumed to result from shock-induced desorption and photoprocessing of dust grain ice mantles in the boundary layer between the outflow jet and the surrounding envelope. A radiative transfer simulation that incorporates chemical variations within the flow shows that the proposed abundance enhancements in the boundary layer are capable of reproducing the observed characteristics of the outflow seen in HCO$^+$ emission in the star forming core L1527. The radiative transfer simulation also shows that the emission lines from the enhanced molecular species that trace the boundary layer of the outflow exhibit complex line profiles indicating that detailed spatial maps of the line profiles are essential in any attempt to identify the kinematics of potential infall/outflow sources.  This study is one of the first applications of a full three dimensional radiative transfer code which incorporates chemical variations within the source. ",
        "introduction": "The complex morphologies of molecular outflows have been identified in numerous studies, mostly based on high resolution, interferometric CO and optical surveys (eg. Arce \\& Goodman, 2001, 2002; Lee et al. 2000, 2002). These surveys show that the molecular distributions are generally composed of a large scale, poorly collimated low velocity outflow that essentially traces the interaction between a surrounding molecular cloud envelope and a collimated high velocity jet.  These morphologies are hard to interpret unambiguously but often resemble limb-brightened shells or sheaths surrounding cavities.  Observations at single-dish resolution suggest that the abundance of HCO$^+$ in star forming cores with bipolar outflows may be enhanced by a factor of up to 100--1000$\\times$ relative to dark-cloud values - for example, a fractional abundance of $4\\times 10^{-8}$ is implied in the case of the Class 0 source B335 (Rawlings, Taylor and Williams 2000).  This, and similar observations, provided the basis for the models of Rawlings~et~al.\\@~(2000) in which the enhancement of HCO$^+$ originates in the turbulent mixing layer which is the interface between the high velocity outflow and the quiescent or infalling core gas which it is steadily eroding. The molecular enrichment is driven by the desorption of molecular-rich ice mantles, followed by photochemical processing by shock-generated radiation fields.  Thus, CO and H$_2$O are released following mantle evaporation.  The CO is then photodissociated and the C that is produced is photoionized by the shock-induced radiation field that is generated in the interface:  \\[ {\\rm CO} + {\\rm h}\\nu \\longrightarrow {\\rm C} + {\\rm O} \\] \\[ {\\rm C} + {\\rm h}\\nu \\longrightarrow {\\rm C}^+ + {\\rm e}^- \\] The subsequent reactions of C$^+$ with H$_2$O lead to an enhancement of the HCO$^+$ abundance: \\[ {\\rm C}^+ + {\\rm H}_2{\\rm O} \\longrightarrow {\\rm HCO}^+ + {\\rm H} \\] This abundance enhancement is only temporary, and progresses as a wave that fans out from the interface into the core. Such behaviour is both qualitatively and quantitatively consistent with the scenario proposed by Velusamy and Langer (1998) for B5 IRS1 on the basis of their $^{12}$CO~(2-1) and C$^{18}$O~(2-1) observations; the data clearly show a parabolic outflow cavity that appears to be steadily widening, so that the opening angle is growing at a rate of 0.006$^\\circ$yr$^{-1}$. An extension of the chemistry to a larger species/reaction set by Viti, Natarajan and Williams (2002) confirmed this result and concluded that other molecular tracers, such as H$_2$CS, SO, SO$_2$ and CH$_3$OH could be similarly affected.  At higher resolutions, aperture synthesis observations are capable of identifying the morphology of the interaction between the jets and their surroundings.  Hogerheijde et al. (1998) made high resolution HCO$^+$ J=1$\\to$0 observations using the Owens Valley Millimetre Array (OVRO) towards a number of low mass YSOs. One of the most interesting results from that survey was the detection of compact emission associated with the walls of an outflow cavity in L1527 (Hogerheijde et al. 1998, figure 1). The $^{12}{\\rm CO} 3-2$ outflow lobes are well-developed, with lengths of $\\sim 0.12~\\rm pc$ and a large opening angle ($\\sim 90^{\\circ}$). The single dish observations of Hogerheijde et al (1997) show that the HCO$^+$ emission is confined to a smaller scale, $\\sim 0.015$ to 0.022~pc, in the center of the flow. The interferometric observations with higher angular resolution show that much of this HCO$^+$ emission originates in the boundary of the lobes of the CO outflow rather than in the quiescent core of the cloud surrounding the outflow.  The enhanced HCO$^+$ emission exhibits a cross-shaped morphology which is, of course, undetectable at single-dish telescope resolution (e.g. see Hogerheijde et al. 1997, figure 3). The data is also consistent with an HCO$^+$ enhancement by a factor of about 10$\\times$ in the outflow sheaths.  The region of apparent HCO$^+$ emission excess has a limited extent, $20-30\\arcsec\\sim 3000-4500~{\\rm AU}$, at an assumed distance of 150~pc, a factor of 10 smaller than the CO emission. If the HCO$^+$ were enhanced over the length of the outflow, then more extended HCO$^+$ emission would be expected in the single-dish observations even if the enhanced region were at a scale too large to be detected by the interferometer (the OVRO will not recover any emission larger than $\\sim 30\\arcsec$). The small extent of the HCO$^+$ emission implies that the enhancement is temporary.  For example, if the HCO$^+$ molecules created in the boundary layer of the outflow close in to the star are not destroyed, then the molecules ought to be transported down the outflow and ought to be detected in the single dish observations. The extent of the HCO$^+$ provides a constraint on the timescale for decay. The position-velocity maps for this source (Hogerheijde et al. 1998, figure 7) suggest that the outflow velocity in the emission-enhanced mixing layer is only $\\sim 5~{\\rm km~s^{-1}}$.  This implies a dynamical age for the HCO$^+$ emitting gas of just $\\sim$3100 yrs, assuming the outflow is close to the plane of the sky.  We therefore require a mechanism to restrict the HCO$^+$ enhancement to this region with a timescale of just a few thousand years.  As stated above, this source has an outflow with a wide opening angle ($\\sim 90^{\\circ}$).  Note that the $^{12}\\rm CO$ emission is much more extensive and has a larger dynamical age than the HCO$^+$ emission excess gas. The orientation of the outflow axis is close to the plane of the sky for this source which implies that the outflow and core emission will be not be clearly separable.  It should also be noted that the dense core may be comparable to this size which could truncate the HCO$^+$ emission.  However, there is still some ambiguity as to the interpretation of the morphology; whether it genuinely traces the abundance, temperature or density enhancements in the outflow interface, or whether it is a simply an excitation effect; the elevation of the HCO$^+$ excitation temperature in the boundary layer perhaps deriving from the outflow cavity acting as a low opacity pathway for photons from the star-disk boundary layer (Spaans et al. 1995).  However, we must also recognise that it would be most unlikely for the HCO$^+$ abundance in the mixing layer to be the same as in the surrounding core.  Previous models of boundary layers, including both analytical studies of turbulent interfaces (e.g. Rawlings and Hartquist 1997) and numerical hydrodynamical calculations (Lim, Rawlings and Williams 2001) have all shown strong HCO$^+$ abundance enhancements within the interface. Moreover, these studies have not included the effects of gas-grain interactions which are likely to further enhance the HCO$^+$ abundance.  In this paper we do not attempt to model the chemical processes within the interface in any detail. Rather we consider the possible levels of molecular enhancement that are consistent with the observations.  This is a timely precursor to the more detailed chemical/radiative transfer calculations that will be required as more higher resolution observational facilities become available. These facilities will be capable of resolving the dynamics and internal structures of sources on milliarcsecond scales.  We have considered two different scenarios that can result in molecular enhancement (and specifically, of HCO$^+$) in the outflow sheath which forms the interface between the bipolar outflow jet and the molecular core envelope: \\begin{itemize} \\item Molecular enrichment through entrainment of dense circumstellar (disk) material into the outflow, and \\item Shock-induced mantle sublimation and photochemical enhancements (as in Rawlings~et~al.\\@~2000), dependent on the precise geometry of the outflow. \\end{itemize} ",
        "conclusions": "In this study we have concentrated on the implications of the radiative transfer modelling for our understanding of the chemical structure of bipolar outflow jet/boundary layers. We have not attempted to present or test a dynamical model of any sophistication. This is in contrast with some previous studies (e.g. Hogerheijde, 1998; Lee et al., 2000; Arce and Goodman, 2002) which have considered the hydrodynamic activity in the jet, boundary layer and surrounding cloud in some detail.  Instead, although we do not constrain the hydrodynamics of the flow, our study makes some important conclusions concerning the origin and level of anomalous chemical activity within boundary layers, and how - using a state-of-the-art radiative transfer code - high resolution observations can be used to diagnose that activity.  In his study, Hogerheijde (1998) utilized an axisymmetric non-LTE Monte-Carlo to study a boundary layer whose thickness is 1/5 of the total width of the (evacuated) jet cavity. His model of the boundary layer was based on a plausible Couette flow in which the density, flow velocity and temperature vary across the boundary layer in simple monotonic fashions, maintaining pressure balance at all points (Stahler, 1994). Whilst somewhat more sophisticated than our uniform density flow, the model had problems in reproducing the observed morphology of L1527 - and in particular the density contrast between the limb-brightened edges and the body of the outflow cavity.  Indeed, reasonable fits to the observed morphology - the line intensities and spatial contrasts in L1527 - were only obtained if the HCO$^+$ abundance in the boundary layer is {\\em not} enhanced over that in the envelope, but the fractional ionization in the boundary layer is very high [X(e$^-$)$\\sim 10^{-4}$]. \\begin{figure*} \\psfig{file=MD740f5.eps,width=450pt,bbllx=0pt,bblly=0pt,bburx=581pt,bbury=544pt} \\caption{ Line profiles of HCO$^+~(1-0)$ for the best model. For clarity, a subset of the 96x96 array of line profiles is displayed: only the central regions are shown and only every fourth profile, in both X and Y directions, is displayed. The viewing angle is $0\\deg$. The tiny peaks of emission are from the cold envelope gas} \\label{plot5} \\end{figure*} \\begin{figure*} \\psfig{file=MD740f6.eps,width=450pt,bbllx=0pt,bblly=0pt,bburx=616pt,bbury=542pt} \\caption{Line profiles of HCO$^+~(1-0)$ for the best model. For clarity, a subset of the 96x96 array of line profiles is displayed: only the central regions are shown and only every fourth profile, in both X and Y directions, is displayed. The viewing angle is $60\\deg$. The tiny peaks of emission are from the cold envelope gas} \\label{plot6} \\end{figure*}  We do not reproduce this result, and in particular we do not find that by raising the HCO$^+$ abundance in the boundary layer that the cavity `fills' due to increased opacity. What we do find is that the density and contrast are more sensitive to the adopted HCO$^+$ abundance enhancement and the thickness of the boundary layer, whilst the detailed morphology is dependent on the assumed outflow shape.  In any case we do not find that such an extreme level of ionization is necessary. Indeed, it is not clear what the source of such a high level of ionization would be. More importantly, it must be remembered that the dominant loss route of HCO$^+$ in dark clouds is dissociative recombination, so it would seem reasonable to expect a significant {\\em suppression} of HCO$^+$ if X(e$^-$)$\\sim 10^{-5}-10^{-4}$.  Our study is more restricted to a chemical analysis of the interface and thus adopts a much simpler dynamical model.  Our modelling has attempted to reproduce the cross-shaped emission seen in L1527 (and other sources) with the appropriate intensity contrasts in the outflow lobes. The cause of the brightness in the cross-arms is essentially due to limb-brightening effects. Clearly, if the boundary layer is too thick then significant emission will be apparent along all lines of sight which pass through the boundary layer, and the morphology would tend to look more like a `bow tie' rather than a hollow cross.  Our {\\em tanh} geometry is somewhat arbitrary, but was chosen because it matches well the observed morphology. The level of HCO$^+$ enhancement and our plausible explanation of its correlation with the shape of the outflow then yields patterns of intensity that closely resemble the high resolution maps; in particular the strong enhancements (and the location of the peaks of emission) in the limb-brightened wings, a natural explanation for the limited extent of the emission and the lack of emission close to the origin.  The essential findings of this paper therefore are: \\begin{enumerate} \\item A strong enhancement of the HCO$^+$ abundance is required in the boundary layer between the outflow jet and the surrounding molecular core in order to explain the observations of certain outflow sources associated with star-forming cores \\item We have proposed a plausible, if simple, mechanism by which this enhancement can occur with the observed morphologies and contrasts; shock liberation and photoprocessing of molecular material stored in icy mantles. The degree of shock activity is closely related to the morphology of the source \\item We have shown how asymmetric, double-peaked line profiles can be generated with strong spatial variations in the relative strength of the red and blue wings. \\end{enumerate}"
    },
    "0403/astro-ph0403701_arXiv.txt": {
        "abstract": "We have obtained moderate resolution ($\\sim 6$ km s$^{-1}$) spectroscopy of several hundred M giant candidates selected from Two Micron All Sky Survey photometry.  Radial velocities are presented for stars mainly in the southern Galactic hemisphere, and the primary targets have Galactic positions consistent with association to the tidal tail system of the Sagittarius (Sgr) dwarf galaxy.  M giant stars selected from the apparent trailing debris arm of Sgr have velocities showing a clear trend with orbital longitude, as expected from models of the orbit and destruction of Sgr.  A minimum 8 kpc width of the trailing stream about the Sgr orbital midplane is implied by verified radial velocity members. The coldness of this stream ($\\sigma_v \\sim 10$ km s$^{-1}$) provides upper limits on the combined contributions of stream heating by a lumpy Galactic halo and the intrinsic dispersion of released stars, which is a function of the Sgr core mass. We find that the Sgr trailing arm is consistent with a Galactic halo that contains one dominant, LMC-like lump, however some lumpier halos are not ruled out. An upper limit to the total mass-to-light ratio of the Sgr core is 21 in solar units.  Evidence for other velocity structures is found among the more distant ($>13$ kpc) M giants.  A second structure that roughly mimics expectations for wrapped, leading Sgr arm debris crosses the trailing arm in the Southern Hemisphere; however, this may also be an unrelated tidal feature.  Among the bright, nearby ($< 13$ kpc) M giants toward the South Galactic Pole are a number with large velocities that identify them as halo stars; these too may too trace halo substructure, perhaps part of the Sgr leading arm near the Sun.  The positions and velocities of Southern Hemisphere M giants are compared with those of Southern Hemisphere globular clusters potentially stripped from the Sgr system.  Support for association of the globular clusters Pal 2 and Pal 12 with Sgr debris is found based on positional and radial velocity matches.  Our discussion includes description of a masked-filtered cross-correlation methodology that achieves better than 1/20 of a resolution element velocities in moderate resolution spectra.  The improved velocity resolution achieved allows tighter constraints to be placed on the coldness of the Sgr stream than previously established. ",
        "introduction": "The extensive length of the tidal tails of the disrupting Sagittarius (Sgr) dSph system has recently been demonstrated in all-sky views of this system provided by the Two Micron All Sky Survey (2MASS) database (Majewski et al. 2003, ``Paper I\" hereafter). The relatively metal-rich stellar populations in the Sgr system means that M giant stars are prevalent in the Sgr debris stream. These stars are easily identified to distances of more than 60 kpc using 2MASS $JHK_s$ photometry, and primary leading and trailing tidal arms from Sgr are evident in all-sky M giant maps. However, as described in Paper I, while Sgr debris dominates the M giant population in the high halo, some ambiguities in the precise length and placement of the Sgr arms remain due to contamination by other M giant populations, particularly near the Galactic plane, and residual photometric errors. Checks on the membership of M giants to the streams using radial velocities can help delineate the extended morphology of the Sgr tidal tails.  Stellar velocities are also useful tracers of mass. As well as revealing clues to the structure and integrity of the Sgr core itself, the placement and motions of its expansive tidal arms can provide important constraints on the mass of the Milky Way and the shape of its potential (e.g., Murai \\& Fujimoto 1980, Lin \\& Lynden-Bell 1982, Kuhn 1993, Johnston et al. 1999, Murali \\& Dubinski 1999, Ibata et al. 2001, Law, Johnston \\& Majewski 2004a,b). In addition, in principle, the velocity dispersion of tidal debris should provide a sensitive probe of the lumpiness of the halo (Moore et al. 1999, Johnston, Spergel \\& Haydn 2002, Ibata et al. 2002, Mayer et al. 2002).  The present study represents a first effort to accumulate velocity data on 2MASS-selected M giants in the apparent Sgr debris trails. It includes a new radial velocity cross-correlation methodology that achieves better than 1/20 of a resolution element discrimination; this analysis approach, in combination with the fact that M giants are intrinsically bright and accessible with modest telescope apertures, places new, tighter constraints on the coldness of the Sgr stream using observations obtained with only a 1-m aperture telescope.  We present results from a pilot radial velocity survey of 284 M giants from Paper I using the Swope 1-m telescope at Las Campanas Observatory.  Most of the stars are in the Southern Galactic Hemisphere and lie within 5 kpc of the best-fit Sgr orbital plane (Paper I). Forty percent of the stars were selected to be in the distinct, Sgr Southern Arc (trailing tail), and, as we show below, most of these stars have velocities consistent with that association.  The remaining stars are primarily very bright, nearby M giants toward the South Galactic Pole, which were selected for study to probe the possibility that Sgr tidal debris may be quite close to the Sun.  We show that this latter sample likely contains an admixture of M giants from the Milky Way's Intermediate Population II/thick disk as well as stars from other halo substructure near the Sun --- perhaps nearby Sgr stars. Future papers will detail observations obtained at other telescopes subsequent to the data collected and analyzed here. ",
        "conclusions": "Radial velocities have been obtained for nearly three hundred M giant candidates identified in the 2MASS survey to be near the Sgr orbital plane.  The survey data here concentrates mainly on stars in the Southern Hemisphere.  The radial velocity trend of the Sgr trailing arm is clear and distinct and provides an important constraint on models of the Sgr debris stream (see Law et al. 2004a,b).  The velocity dispersion of the trailing debris arm limits both the degree of lumpiness of the MW halo as well as the mass of the Sgr bound core. Because these two contributors to the velocity dispersion play off one another, it may be difficult simultaneously to have both a large $M/L$ for the Sgr core {\\it and} a lumpy halo:  At present we limit the $(M/L)_{V,tot}$ of the Sgr core to 21 in solar units (see Law et al. 2004a,b), and the heating of the Sgr stream to be consistent with expectations for debris encounters with only the LMC.  More stringent constraints on both may be possible after dynamical study of the Sgr tidal boundary to determine the zero-age velocity dispersion of Sgr debris.  At least one other debris stream may be present in the distant M giant sample, and its velocity trend is consistent with model (Law et al. 2004a,b) predictions for wrapped, leading arm material in the Southern Hemisphere, but it may also be an unrelated tidal stream.  The globular cluster Pal 12 may be associated with the stars in this feature, as previously postulated by Dinescu et al. (2000), but we note that the radial velocity of this cluster lies within 1.5-2$\\sigma$ of the mean for Sgr trailing arm M giants at this same location in the Galaxy. Based on its radial velocity and positional match with Sgr M giants, we find compelling evidence that the globular cluster Pal 2 is associated with Sgr trailing arm debris.  Among the closer ($< 13$ kpc) M giant stars observed are more than half a dozen with the very high positive radial velocities expected for Sgr debris in the neighborhood of the Sun, and we point to the presence of many high velocity ($|v_{GSR}| > 150$ km s$^{-1}$), high latitude M giants that are difficult to account for as other than halo substructure relatively close to the Sun.  The data presented here are better understood within the Sgr debris context by reference to interpretive models.  Such models are given in Law et al. (2004a,b), in which the data presented here provide useful observational constraints.  Moreover, additional velocities obtained in the Northern Hemisphere will help clarify the interpretations offered here. Future papers will present new data that increase by more than five times the present M giant velocity sample.  We thank Carnegie Observatories Director A. Oemler for generous access to the Swope telescope and all of the Las Campanas staff for their diligent and generous support at the telescope. The results presented in this publication make use of data from the Two Micron All Sky Survey (2MASS), which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center (IPAC), funded by the National Aeronautics and Space Administration and the National Science Foundation.  This work was supported by National Science Foundation grant AST-0307851 and a Space Interferometry Mission Key Project grant, NASA/JPL contract 1228235. MFS acknowledges support from NASA/JPL contract 1234021.  MDW was supported in part by NSF grant AST-9988146.  This work was also partially supported by the Celerity Foundation.  We appreciate comments from the referee that have improved the clarity of the paper."
    },
    "0403/astro-ph0403537_arXiv.txt": {
        "abstract": "The 692 GHz para ground-state line of D$_2$H$^+$ has been detected at the Caltech Submillimeter Observatory towards the pre-stellar core 16293E. The derived D$_2$H$^+$ abundance is comparable to that of H$_2$D$^+$, as determined by observations of the 372 GHz line of ortho-H$_2$D$^+$. This is an observational verification of recent theoretical predictions (Roberts, Herbst \\& Millar 2003), developed to explain the large deuteration ratios observed in cold, high-density regions of the interstellar medium associated with low mass pre-stellar cores and protostars. This detection confirms expectations that the multiply deuterated forms of H$_3$$^+$ were missing factors of earlier models. The inclusion of D$_2$H$^+$ and D$_3$$^+$ in the models leads to predictions of higher values of the D/H ratio in the gas phase. ",
        "introduction": "Recently, millimeter and submillimeter spectroscopy of the dense interstellar medium has shown that, in cold dense regions, deuterated molecular species are highly abundant, sometimes more than 10$^{-1}$ of the H version. Amazingly, doubly and triply deuterated species can be observed, e.g. D$_2$CO \\citep{ceccarelli98}, NHD$_2$ \\citep{roueff00}, CHD$_2$OH \\citep{parise02}, D$_2$S \\citep{vastel03}, ND$_3$ \\citep{lis02,vdtak02}, CD$_3$OH \\citep{parise04}. Several models have been developed to account for such high levels of deuteration \\citep{tielens83,roberts00a,roberts00b}. \\citet{phillips02} have pointed out that the deuteration of H$_3$$^+$ will be extended beyond H$_2$D$^+$, to D$_2$H$^+$ and D$_3$$^+$, and that detection of D$_2$H$^+$ might be possible. A calculation taking a high degree of deuteration into account has been carried out by \\citet{roberts03} and \\citet{walmsley04}, confirming the expectation that, in dense depleted regions, the abundance of D$_2$H$^+$ will be similar to that of H$_2$D$^+$, and that D$_3$$^+$ will be abundant.  The key enabling work in the astronomical search for D$_2$H$^+$ is the laboratory measurement of the para ground-state transition (1$_{10}$-1$_{01}$) by \\citet{hirao03}. We report here the first astronomical detection of that transition.\\\\  Chemical reactions go in the direction to minimize energy. The chemical fractionation process favors the production of the heavier more deuterated species, because of the mass dependence of the zero-point vibration energies of the isotopic variants. Gas phase species are expected to be depleted at the centers of cold, dark clouds, since they accrete on the dust grains \\citep[see, e.g.,][]{charnley97}. A series of observations has shown that the abundances of molecules like CO decrease in many pre-stellar cores \\citep{bacmann02}. The removal of these reactive species affects the gas-phase chemistry and particularly the deuterium fractionation within the cloud. Indeed, the removal of species that would normally destroy H$_3$$^+$ \\citep[e.g. CO;][]{roberts00a} means that H$_3$$^+$ is more likely to react with HD and produce H$_2$D$^+$. For example, if [CO/H$_2$]~$\\sim$~5 $\\times$ 10$^{-6}$ \\citep{bacmann02}, this leaves HD at [HD/H$_2$]~$\\sim$~5 $\\times$ 10$^{-5}$ as the most abundant molecule available for reaction with H$_3^+$ and H$_2$D$^+$, and favors the production of high deuterium content molecules:  \\begin{equation} H_3^+ + HD \\longleftrightarrow H_2D^+ + H_2 + \\Delta E_a \\end{equation} \\begin{equation} H_2D^+ + HD \\longleftrightarrow D_2H^+ + H_2 + \\Delta E_b \\end{equation} \\begin{equation} D_2H^+ + HD \\longleftrightarrow D_3^+ + H_2 + \\Delta E_c \\end{equation}  where $\\Delta$E$_a$, $\\Delta$E$_b$ and $\\Delta$E$_c$ are the released energies of the exothermic reactions. Using the zero-point energies computed by \\citet{ramanlal03}, and the energy of the first allowed rotational state of the H$_3^+$ molecule permitted by the Pauli exclusion principle ($\\sim$ 92 K), these values are: $\\Delta$E$_a$ = $\\sim$ 230 K, $\\Delta$E$_b$ = $\\sim$ 180 K and $\\Delta$E$_c$ = $\\sim$ 230 K.  After a long frustrating search \\citep{phillips85,pagani92,vandishoeck92,boreiko93}, and with the advent of new submillimeter receivers, H$_2$D$^+$ was detected toward two young stellar objects, NGC 1333 IRAS 4A (Stark et al. 1999) and IRAS 16293-2422A (Stark et al. 2004), although with relatively low signal strength. The H$_2$D$^+$ search has now been extended to pre-stellar cores, and has been detected with relatively strong emission (Caselli et al. 2003; Caselli et al. 2004, {\\it in preparation}; Vastel et al. 2004, {\\it in preparation}) confirming that H$_2$D$^+$ is dramatically enhanced in a gas depleted of most molecules.  The ammonia and DCO$^+$ emission around the proto-binary system IRAS16293-2422 \\citep{wootten87,mizuno90,lis02} does not only peak on IRAS16293-2422 itself but shows a second peak, about 90$^{\\prime\\prime}$ to the southeast, in a condensation called 16293E. \\citet{lis02} found that CO in this region is depleted by a factor of 7. It is known to be a region where deuterium fractionation is strong and was chosen to be searched for D$_2$H$^+$. ",
        "conclusions": "From the observed line strengths, given in column 3 of Table \\ref{table1}, we estimate the H$_2$D$^+$ and D$_2$H$^+$ column densities (see Table \\ref{table2}) for an excitation temperature T$_{ex}$ of 10 K, assuming a 25\\% calibration uncertainty (3 $\\sigma$). The column density is given by:  \\begin{equation} N_{tot}=\\frac{8\\pi\\nu^3}{c^3}\\frac{Q(T_{ex})}{g_uA_{ul}}\\frac{e^{E_{u}/T_{ex}}}{e^{h\\nu/kT_{ex}}-1}\\int_{}^{}\\tau\\, dv \\end{equation}  where Q(T$_{ex}$) is the partition function, Assuming LTE conditions, we can estimate the optical depth from the observed line intensity:  \\begin{equation} T_{mb} = [J_{\\nu}(T_{ex})-J_{\\nu}(T_{bg})](1-e^{-\\tau}) \\end{equation}  where $J_{\\nu}(T) = (h\\nu/k)/(e^{h\\nu/kT}-1)$ is the radiation temperature of a blackbody at a temperature T, and T$_{bg}$ is the cosmic background temperature of 2.7 K. In the case of the H$_2$D$^+$ transition, g$_u$ = 9, A$_{ul}$ = 1.04 10$^{-4}$ s$^{-1}$, E$_{ul}$ = 17.9 K; in the case of the D$_2$H$^+$ transition, g$_u$ = 9, A$_{ul}$ = 4.55 10$^{-4}$ s$^{-1}$, E$_{ul}$ = 33.2 K. The derived column densities depend on the assumed value of the excitation temperature. Using NH$_3$, \\citet{mizuno90} estimate the gas temperature to be 12 K. Using D$_2$CO line ratios, \\citet{loinard01} obtained a rotational temperature of 8-10 K. Thus we quote, in Table \\ref{table2}, the values obtained for an excitation temperature of 10 K. Figure \\ref{figure4} presents the evolution of the ortho-H$_2$D$^+$ and para-D$_2$H$^+$ column densities as well as the para-D$_2$H$^+$/ortho-H$_2$D$^+$ ratio, as a function of temperature between 9 and 15 K. Figure \\ref{figure4} and Table \\ref{table2} represent the case where the source emission is extended compared to the beam size. If the source emission is comparable to or smaller than the CSO beam size, the para-D$_2$H$^+$/ortho-H$_2$D$^+$ ratio is then increased by a factor of $\\sim$ 1.5 at the average excitation temperature of 10 K. At thermal equilibrium, the ortho to para (respectively para to ortho) concentration ratio for H$_2$D$^+$ (respectively D$_2$H$^+$) is equal to 9 $\\times$ exp(-86.4/T) (respectively 9/6 $\\times$ exp(-50.2/T)), so that at 10 K, this ratio would be $\\sim$ 2 $\\times$ 10$^{-3}$ (respectively $\\sim$ 10$^{-2}$). However, taking into account the limited rates of the spin allowed collisions with H$_2$, it is found that at these low temperatures, the ortho to para H$_2$D$^+$ concentration ratio is close to unity \\citep{gerlich02}. The para to ortho D$_2$H$^+$ ratio is estimated by \\citet{walmsley04} to be about one, for the same conditions. The para-D$_2$H$^+$/ortho-H$_2$D$^+$ ratio presented in Figure \\ref{figure4} should then approximately represent the actual D$_2$H$^+$/H$_2$D$^+$ ratio.  The 1.3 mm dust continuum strength \\citep[see][]{lis02} is $\\sim$ 0.3 Jy in a 11$^{\\prime\\prime}$ beam and $\\sim$ 1.3 Jy in a 20$^{\\prime\\prime}$ beam corresponding to the angular resolution of the H$_2$D$^+$ and D$_2$H$^+$ data. Assuming a dust temperature of 12 K and a mass opacity coefficient of 0.005 cm$^2$~g$^{-1}$ (appropriate for pre-stellar cores), we derive an H$_2$ column density of $\\sim$ 5~10$^{23}$ cm$^{-2}$. We then derive the H$_2$D$^+$ and D$_2$H$^+$ abundances to range between $\\sim$ 10$^{-10}$ (at 10 K) and $\\sim$ 10$^{-11}$ (at 15 K) compatible with abundances found by \\citet{roberts03} for a cloud at 10 K and n(H$_2$)=3 10$^6$ cm$^{-3}$.  The main result of this work is {\\it the detection of D$_2$H$^+$, with an abundance comparable to that of H$_2$D$^+$}. This is a remarkable verification of recent theoretical predictions, aimed at explaining the large deuteration ratios observed in low mass pre-stellar cores and protostars. Two models, \\citet{roberts03} and \\citet{walmsley04}, have recently considered the effect of including all possible deuterated isotopomers of H$_3^+$ in the chemical networks, as suggested by \\citet{phillips02}. \\citet{roberts03} studied the temporal evolution of a cold and dense cloud, and found that at late times ($\\sim$ 10$^4$ yr), when CO is severely depleted in the gas phase (more than a factor of 1000), the D$_2$H$^+$/H$_2$D$^+$ ratio reaches unity. \\citet{walmsley04} studied the evolution of gas, depleted in CO, as a function of gas density and grain size distribution. For densities larger than 10$^6$ cm$^{-3}$, they also found that D$_3$$^+$ can be the most abundant ion and that the D$_2$H$^+$/H$_2$D$^+$ ratio reaches unity (see their figure 2). Both models need an extreme CO depletion to account for  such a ratio. As discussed in the Introduction, the measured CO depletion in 16293E is a factor 7 \\citep{lis02}, rather than the extreme CO depletion needed by the models. However, the CO was measured in a large region (31\") compared with that probed by the present observations: 20$^{\\prime\\prime}$ for the H$_2$D$^+$ and 11$^{\\prime\\prime}$ for the D$_2$H$^+$ observations. Also, regions containing CO will not contain much H$_2$D$^+$ and D$_2$H$^+$ and vice-versa, so the inevitable inhomogeneities in the region inhibit a clear result.  In summary, after some years of inconclusive results for theoretical models in understanding the observed high deuteration ratios of doubly and triply deuterated molecules, the present observation seems to suggest that the basic process is now almost completely understood: the large deuteration is due to extreme CO depletion, and the factor that was previously missing in the models is the multiply deuterated forms of H$_3^+$. This is quite an achievement, and one remaining step will be to verify that the last prediction, a significant abundance of D$_3^+$, is also fulfilled. At present, the lines detected here are the only ones available for H$_2$D$^+$ and D$_2$H$^+$. Knowledge of the ortho to para ratios and abundances of H$_2$D$^+$ and D$_2$H$^+$ would be considerably improved if the ground state transitions of para-H$_2$D$^+$ (at 1370.15 GHz) and ortho-D$_2$H$^+$ (at 1476.60 GHz) were available. These lines could be detected from space telescopes such as Herschel with the Heterodyne Instrument for the Far-Infrared and also possibly from the stratospheric observatory SOFIA. However, D$_3$$^+$, like H$_3$$^+$ has no permanent dipole moment. Therefore this molecule probably can only be detected in absorption in the near-infrared."
    },
    "0403/astro-ph0403270_arXiv.txt": {
        "abstract": "We reported in a previous paper the discovery of large-scale structure of Lyman $\\alpha$ emitters (LAEs) at $z=4.86\\pm 0.03$ with a projected size of $20 h_{70}^{-1}$ Mpc $\\times$ $50 h_{70}^{-1}$ Mpc in narrow-band data of a $25' \\times 45'$ area of the Subaru Deep Field ($\\Omega_0=0.3, \\lambda_0=0.7, H_0 = 70 h_{70}$ km s$^{-1}$ Mpc$^{-1}$). However, the surveyed area, which corresponds to $55 h_{70}^{-1}$ Mpc $\\times 100 h_{70}^{-1}$ Mpc, was not large enough that we can conclude that we are seeing a typical distribution of $z\\simeq 5$ LAEs. In this Letter, we report the results of follow-up imaging of the same sky area using a new narrow-band filter (NB704, $\\lambda_c=7046$\\AA\\ and FWHM$=100$\\AA) to detect LAEs at $z=4.79$, i.e., LAEs lying closer to us by $39 h_{70}^{-1}$ Mpc on average than the $z=4.86$ objects. We detect 51 LAEs at $z=4.79 \\pm 0.04$ down to ${\\rm NB704}=25.7$, and find that their sky distribution is quite different from the $z=4.86$ LAEs'. The clustering of $z=4.79$ LAEs is very weak on any scales and there is no large-scale high-contrast structure. The shape and the amplitude of the angular correlation function are thus largely different between the two samples. These results demonstrate a large cosmic variance in the clustering properties of LAEs on scales of $\\sim 50 h_{70}^{-1}$ Mpc. ",
        "introduction": "\\label{sec:introduction}  Search for Lyman $\\alpha$ emission of galaxies using a narrow-band filter is a powerful tool to detect high-$z$ faint galaxies. Indeed, many observations have successfully detected such Lyman $\\alpha$ emitters (LAEs) from $z\\sim2$ up to $z\\simeq 6.6$ (e.g., Hu, Cowie, \\& McMahon 1998; Pascarelle, Windhorst, \\& Keel 1998; Campos et al. 1999; Hu, McMahon, \\& Cowie 1999; Kudritzki et al. 2000; Rhoads et al. 2000; Stiavelli et al. 2001; Ajiki et al. 2002; Hu et al. 2002; Venemans et al. 2002; Fynbo et al. 2003; Kodaira et al. 2003; Ouchi et al. 2003a; Shimasaku et al. 2003; see also Maier et al. 2003). Narrow-band surveys can also map effectively large-scale distributions of LAEs (e.g., Steidel et al. 2000; Venemans et al. 2002; Hu et al. 2004).  We recently made a survey of LAEs at $z=4.86$ using the narrow-band filter NB711 ($\\lambda_c=7126$\\AA, FWHM=73\\AA) in an area of $25' \\times 45'$, and found a large-scale structure of LAEs of $20h_{70}^{-1}$ Mpc $\\times$ $50h_{70}^{-1}$ Mpc size (Shimasaku et al. 2003). This is the first discovery of large-scale structure in young universes, suggesting that the birth of large-scale structure is very early in the history of the universe and that LAEs are strongly biased against dark matter, since Cold Dark Matter (CDM) models predict that the density fluctuations of dark matter are very small on such large scales. However, the size of the large-scale structure is nearly comparable to the size of the survey region ($55h_{70}^{-1}$ Mpc $\\times$ $100h_{70}^{-1}$ Mpc), and thus we cannot safely conclude that we are seeing a typical distribution of LAEs at $z\\simeq 5$. To address this issue, it is strongly needed to enlarge the survey volume.  Motivated by this, we made a followup imaging survey of LAEs at $z=4.79$, i.e., LAEs located closer to us by $39h_{70}^{-1}$ Mpc than those at $z=4.86$, in exactly the same sky area. Using these data, we examine differences in the sky distribution of LAEs between the two redshifts. We adopt $\\Omega_0=0.3$, $\\lambda_0=0.7$, and $H_0 = 70 h_{70}$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "Hamana et al. (2004) found that the observed ACF of $z=4.86$ LAEs given in Ouchi et al. (2003a) cannot be reproduced by a simple halo model which assumes LAEs to be associated with dark haloes, because of too strong the observed correlation on scales $ \\gsim 2'$. We find that our $z=4.86$ LAEs have large-scale structure and that their ACF is high and not fit by a power law, while the clustering of $z=4.79$ LAEs is very weak. These results may suggest that the clustering of LAEs is typically weak, possibly tracing the dark-matter distribution, and that we happened to observe an unusual region in the $z=4.86$ universe where LAEs form a large, coherent structure of a size of $\\sim 50 h_{70}^{-1}$ Mpc. Conversely, if it turns out, from a larger survey, that the $z=4.86$ region we observed is relatively common in high-$z$ universes, this will suggest that LAEs and LBGs are separate populations in terms of clustering properties, since the clustering of LBGs has been found to be approximated well by halo models (Hamana et al. 2004). Detailed modeling of the clustering of LAEs based on an enlarged sample will give important hints on the nature of LAEs.  Although we did not find in our two LAE samples a clear difference in the number density, a large field-to-field variance in the clustering of LAEs, including velocity structures, can influence measurements of the number density of LAEs based on narrow-band surveys, especially if the survey volumes are smaller than ours and if $b_g$ is much larger than unity; the $b_g$ value derived from the clustering of the $z=4.86$ LAEs is as large as $\\sim 10$. Shallower surveys will suffer from larger variances, since LAEs with brighter Lyman $\\alpha$ luminosities tend to be clustered more strongly (Ouchi et al. 2003a). For instance, Ajiki et al. (2003) found that the number density of $z=5.7$ LAEs in their sample of $2 \\times 10^5 h_{70}^{-3}$ Mpc$^3$ is three times higher than that estimated by Rhoads \\& Malhotra (2001) based on a similar survey volume.  To summarize, our observations show that it is necessary to survey a much larger volume than ours in order to derive the average clustering properties of LAEs."
    },
    "0403/astro-ph0403100_arXiv.txt": {
        "abstract": "s{ Recent \\chandra\\ and \\xmm\\ surveys have confirmed that the cosmic X-ray background is mostly due to accretion onto super-massive black holes, integrated over cosmic time. Here we review the results obtained from the photometric and spectroscopic follow-up observations of the 122 X-ray sources detected by the HELLAS2XMM 1dF Survey down to a 2--10~keV flux of $\\approx10^{-14}$~\\cgs. In particular, we focus on the multiwavelength properties of a few intriguing classes of X-ray sources: high X-ray-to-optical flux ratio sources, Type~2 quasars, and XBONGs.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403499_arXiv.txt": {
        "abstract": "We present FUSE spectroscopy and supporting data for star-forming regions in nearby galaxies, to examine their massive-star content and explore the use of abundance and population indicators in this spectral range for high-redshift galaxies. New far-ultraviolet spectra are shown for four bright H II regions in M33 (NGC 588, 592, 595, and 604), the H II region NGC 5461 in M101, and the starburst nucleus of NGC 7714, supplemented by the very-low-metallicity galaxy I Zw 18. In each case, we see strong Milky Way absorption systems from H$_2$, but intrinsic absorption within each galaxy is weak or undetectable, perhaps because of the ``UV bias\" in which reddened stars which lie behind molecular-rich areas are also heavily reddened. We see striking changes in the stellar-wind lines from these populations with metallicity, suggesting that C II, C III, C IV, N II, N III, and P V lines are potential tracers of stellar metallicity in star-forming galaxies. Three of these relations - involving N IV, C III, and P V - are nearly linear over the range from O/H=0.05--0.8 solar. The major difference in continuum shapes among these systems is that the giant H II complex NGC 604 has a stronger continuum shortward of 950 \\AA\\  than any other object in this sample. Small-number statistics would likely go in the other direction; we favor this as the result of a discrete star-forming event $\\approx 3$ Myr ago, as suggested by previous studies of its stellar population. ",
        "introduction": "The history of galaxies is, in large part, the history of star formation. Massive stars play key roles both as highly visible tracers of star formation and as players in altering surrounding star formation and both energy and chemistry of the interstellar medium. It is these stars which dominate the observed properties of actively star-forming galaxies. The massive part of stellar populations is most clearly observed in the ultraviolet, where their energy distributions peak and competing light from cooler stars is minimal, as long as the foreground extinction allows escape of enough of this radiation. In such environments, ultraviolet studies of star-forming regions have proven fruitful in understanding these populations. The recent opening of the far-ultraviolet window, between Lyman $\\alpha$ and the Lyman limit, for deep observations, allows study of massive hot stars in a range where they fully dominate the spectrum. This relatively narrow band contains an embarrassment of spectral riches, with numerous lines from stellar winds as well as interstellar material both atomic and molecular. These include the strong and highly-ionized lines of O VI and the unique ability to measure cold H$_2$. In addition, this piece of the spectrum is accessible for high-redshift galaxies, at least for composite samples where the Lyman $\\alpha$ forest can be averaged adequately, allowing the possibility of direct comparisons of stellar populations over a large span of cosmic time.  The very sensitivity of the far-UV light to star formation and reddening makes it a purer probe of some properties of star-forming regions than observations at longer wavelengths. Since only short-lived stars contribute significantly in the far-UV range, the details of star-forming history should matter only for very brief bursts (such as might be found in individual H II regions, but are less likely on galaxy scales). This makes the far-UV spectrum more sensitive to the stellar population itself than to its history. Furthermore, although the extinction is high, its differential effect across the far-UV band is modest, and paradoxically the effective reddening to the stars we see is smaller than found at longer wavelengths. In observing stars intermingled with highly structured dust distributions, the ``picket-fence\" effect (Heisler \\& Ostriker 1988) means that most of the stars are so reddened as to make no significant contribution in the deep ultraviolet; all the stars we see are only lightly reddened. This also reduces the effects of the forest of H$_2$ absorption features because of the mixing of molecular gas and dust.  Observational data on nearby galaxies, before the {\\it Far-Ultraviolet Spectroscopic Explorer} (FUSE), were limited to a handful of star-forming systems. Four starburst galaxies were observed using HUT on {\\it Astro-2} (Leitherer et al. 1995, 2002), largely to measure escaping radiation in the Lyman continuum, which provided initial data for comparison with models based on stellar spectra dating back to {\\it Copernicus}. The strongest features fall into two blends near 970 (Ly $\\gamma$ + C III) and 1030 \\AA\\  (O VI + Ly $\\beta$ + C II). A {\\it Voyager 2} observation of M33, with some spatial resolution in one direction, was analyzed by Keel (1998), showing that its far-UV continuum is virtually identical to those of the powerful starbursts, and that NGC 604 is bluer in this range than the overall disk. These data also suggested significant Lyman $\\alpha$ emission, with spatial profile suggesting an origin in the diffuse ISM rather than giant H II regions. Ironically, until the availability of FUSE observations, the richest information on galaxies shortward of Lyman $\\alpha$ came from objects at high redshift, particularly composite spectra of Lyman-break galaxies (Steidel et al. 2001, Shapley et al. 2003). To enable comparisons between local, well-studied star-forming systems and these powerful, young objects, we have undertaken a series of FUSE observations of star-forming regions innearby galaxies. We present here the analysis of these spectra in the context of their stellar populations and systematic changes with metallicity. A companion paper (Keel, Shapley, \\& Steidel 2004) considers the comparison with composite spectra of Lyman-break galaxies. ",
        "conclusions": "We have used FUSE spectra of star-forming regions in nearby galaxies, whose gas-phase metallicities range from 0.05-0.8 solar, to explore the utility of far-ultraviolet spectra in measuring the abundances in star-forming galaxies, as well as to probe the massive-star populations in these galaxies. The absorption lines from radiatively-driven winds prove to be very sensitive to metal abundance; all six strong and unblended species (including C IV from archival IUE data) have a strong, monotonic metallicity dependence. For N IV, C III, and P V, the relation between straighforward equivalent-width values and oxygen abundance from emission-line spectra is closely linear, suggesting that these lines will be useful in tracing the chemical history of galaxies from $z=3-4$, beyond which the Lyman $\\alpha$ forest makes even composite spectra progressively less informative.  The continuum of NGC 604 departs from the uniform shape of the other objects below 950 \\AA . After considering the effects of small-number statistics among the massive stars in these objects, we conclude that this difference probably traces to a discrete burst of star formation $\\approx 3$ Myr ago in NGC 604. This region had been considered by several previous studies to have hoisted such a burst, on grounds of both morphology of the gas and fitting of the H-R diagram."
    },
    "0403/astro-ph0403385_arXiv.txt": {
        "abstract": "{  Recently, combining radial  velocities from Keck/HIRES \\'echelle spectra with published proper motion membership probabilities, \\citet{Coe_02} observed a sample of $21$ stars, probable members of Palomar~13, a  globular cluster in  the Galactic  halo.  Their  projected velocity dispersion $\\sigma_\\mathrm{p}  = 2.2 \\pm 0.4$~km\\thinspace s$^{-1}$\\  gives a mass-to-light ratio  $\\mathcal{M}/L_V =  40^{+24}_{-17}$,  about one  order  of magnitude larger than the usual estimate for globular clusters.  We present  here radial velocities  measured from three  different CCD frames  of commissioning  observations obtained  with the  new ESO/VLT instrument  FLAMES  (Fibre Large  Array  Multi Element  Spectrograph). From these  data, now publicly  available, we measure  the homogeneous radial  velocities   of  eight  probable  members   of  this  globular cluster. A new projected velocity dispersion $\\sigma_\\mathrm{p} = 0.6$--$0.9 \\pm 0.3$~km\\thinspace s$^{-1}$\\ implies  Palomar~13 mass-to-light ratio $\\mathcal{M}/L_V =  3$--$7$,  similar to the  usual value for globular  clusters.  We discuss briefly the two most  obvious reasons for the previous unusual mass-to-light ratio finding: binaries,  now clearly detected, and more homogeneous data from the multi-fibre FLAMES spectrograph. ",
        "introduction": "All the dynamical studies of nearby globular clusters have established that these  dynamical systems contain  no dark matter, apart  from the expected  stellar remnants  such  as white  dwarves  and neutron  stars \\citep[e.g.,][]{PrrMen93,MenHee97}.  Consequently, globular clusters may be the  most massive stellar systems in which no non-baryonic  dark  matter is  dynamically  detected, while  dynamical evidence  for non-baryonic  dark matter  seems to  be present  in most galaxies, from the  faintest dwarf spheroidals (dSphs) to the brightest cDs galaxies, and clusters of galaxies as well.  Some  of  the  local  dSphs, around the Galaxy and M31, have integrated absolute luminosities  similar  or  fainter than those of the brightest Galactic globular clusters.  Since  there  is  evidence  that  some  Galactic  dSph  galaxies  are dark-matter-dominated,  it  is  therefore  reasonable  to check if some globular clusters  do  present  dynamical  evidence  for non-baryonic dark matter. Such a possibility  may  be  emphasized by the current predictions of Cold Dark Matter (CDM) numerical simulations of galaxy formation, in which the number of low-mass dark-matter  substructures orbiting the halo of massive galaxies largely exceeds the  number  of  dwarf galaxies observed in the halos of both our Galaxy and M31 \\citep[e.g., ][]{Kln_99,Moe_99,Moe_01}. However, recently improved CDM models may correctly predict the observed number of satellite galaxies \\citep{Biy03}.  Could some of the globular clusters in the outer parts of the Galactic halo be such dark-matter  substructure~?  These remote stellar systems have so  far been  poorly studied because  of the difficulties  in the acquisition  of high-quality  radial velocities  and  proper motions, direct  consequences of the  faintness and  sparsity of  these distant stellar systems. They  are nevertheless important probes of  the formation and evolution  of the  Galaxy,  as their  ages  and metallicities  provide direct constraints  on the duration  of halo formation process  and on the time-scale  for Galactic chemical enrichment, while  the shape and extent  of the  Galactic dark  halo are  constrained by  their orbital properties.  In  1998, a  program  was started  at  Californian Institute of Technology to  study the  internal dynamics  of  seven distant  halo  globular  clusters  using the  High Resolution   Echelle    Spectrometer   (HIRES)   at    the   W.~M.~Keck Observatory. The aim was the first direct measurements of the velocity dispersions and mass-to-light ratios for these clusters.  Six clusters in this sample exhibited  velocity dispersions $\\sigma_\\mathrm{p} \\sim 1$~km\\thinspace s$^{-1}$, translating  into  mass-to-light  ratio  values  typical  of  globular clusters $\\mathcal{M}/L_V\\sim 3$ (all mass-to-light ratios quoted in this paper are in solar units). Only one cluster, the halo globular cluster  Palomar~13,  displayed   a  velocity  dispersion  larger  than expected.  \\citet[ also referenced  below as  the  Keck study]{Coe_02} presented  a  careful   analysis,  combining  radial  velocities  from Keck/HIRES \\'echelle spectra  with published  proper  motion membership probabilities from \\citet{Sil_01}.  They obtained a sample of 21 stars,  probable members  of Palomar~13.   Their  projected, intrinsic velocity  dispersion  of  $\\sigma_\\mathrm{p}   =  2.2\\pm 0.4$~km\\thinspace s$^{-1}$\\  implied  a mass-to-light ratio $\\mathcal{M}/L_V= 40^{+24}_{-17}$, about one order of magnitude larger than  the usual value for globular clusters. C\\^ot\\'e et al.\\ discussed at length all  possible reasons for such an  unusual result:  (i) some velocity  ``jitter'' among  the red giants, (ii) a  few binary stars, (iii) a  non-standard mass function, (iv) process of dissolving into the Galactic halo through catastrophic tidal  heating during  a  recent perigalacticon  passage,  or (v)  the presence of a massive non-baryonic dark matter halo.  It  is worth  emphasizing that,  in C\\^ot\\'e et al., careful determination  of the error  bars made  the usual  mass-to-light ratio values for globular clusters $\\mathcal{M}/L_V\\sim 3$ at about two sigmas from the $\\mathcal{M}/L_V$ value obtained for Palomar~13.  Because  of  this  marginally  significant  and  puzzling  result  from C\\^ot\\'e et al., some more  spectroscopic data of stars in the field  of Palomar~13  were acquired  during the  commissioning  of the ESO/VLT instrument FLAMES.  We  present hereafter new high-quality and homogeneous radial velocities for 46 stars, 9 of them being members of Palomar~13, which provide new velocity dispersion and mass-to-light ratio values for  this globular cluster.  The remaining  of this paper is  as  follows:  Sect.~2  presents  the  observation  and  the  data reduction, Sect.~3 discusses the membership of the stars, Sect.~4 gives  the  new velocity  dispersion  and corresponding  mass-to-light ratio, and Sect.~5 discusses the plausible reasons for the difference between the present results and those obtained by C\\^ot\\'e et al. ",
        "conclusions": ""
    },
    "0403/astro-ph0403666_arXiv.txt": {
        "abstract": "{ We present the COMBO-17 object catalogue of the Chandra Deep Field South for public use, covering a field which is $31\\farcm5 \\times 30\\arcmin$ in size. This catalogue lists astrometry, photometry in 17 passbands from 350 to 930~nm, and ground-based morphological data for 63,501 objects. The catalogue also contains multi-colour classification into the categories {\\it Star}, {\\it Galaxy} and {\\it Quasar} as well as photometric redshifts. We include restframe luminosities in Johnson, SDSS and Bessell passbands and estimated errors. The redshifts are most reliable at $R<24$, where the sample contains approximately 100 quasars, 1000 stars and 10000 galaxies. We use nearly 1000 spectroscopically identified objects in conjunction with detailed simulations to characterize the performance of COMBO-17. We show that the selection of {\\it quasars}, more generally type-1 AGN, is nearly complete and minimally contaminated at $z=[0.5,5]$ for luminosities above $M_B = -21.7$. Their photometric redshifts are accurate to roughly 5000~km/sec. Galaxy redshifts are accurate to 1\\% in $\\delta z/(1+z)$ at $R<21$. They degrade in quality for progressively fainter galaxies, reaching accuracies of 2\\% for galaxies with $R \\sim 22$ and of 10\\% for galaxies with $R>24$. The selection of stars is complete to $R \\sim 23$, and deeper for M stars. We also present an updated discussion of our classification technique with maps of survey completeness, and discuss possible failures of the statistical classification in the faint regime at $R\\ga 24$. ",
        "introduction": "The Chandra Deep Field South (CDFS) is one of the most well-studied patches of sky. It is the target of enormous observational efforts across a wide range of photon energies. The variety of imaging and spectroscopic data sets shall improve our understanding of fundamental processes in galaxy evolution.  A large amount of public data are contributed by the {\\it Great Observatories Origins Deep Survey} (GOODS). This survey obtains deep images of the field using all of NASA's great space-based facilities: the Chandra X-ray observatory \\citep[CXO,][]{GiaCDFS}, the {\\it Advanced Camera for Surveys} (ACS) onboard the {\\it Hubble Space Telescope} \\citep[HST,][]{Giav04}, and the new infrared space telescope Spitzer. Further space-based observations include the {\\it Ultra Deep Field} (UDF) project targetting a small part of the field with a single ACS pointing, deep observations with ESA's X-ray observatory XMM-Newton (PI Bergeron), and the wider-area ACS imaging by the GEMS team \\citep{Rix04}.  In this paper, we publish data and results from ground-based observations of the CDFS. Our project, COMBO-17, has targetted the CDFS among four other fields. They are all observed with the {\\it Wide Field Imager} \\citep[WFI,][]{WFI1,WFI2} at the MPG/ESO 2.2m-telescope on La Silla, Chile. This camera covers an area of more than $0.5\\degr \\times 0.5\\degr$, which is larger than the field initially observed from space by GOODS. The footprint of this larger WFI-based image is occasionally called {\\it Extended CDFS} or E-CDFS, but we just call it CDFS here. The purpose of the later GEMS images was to cover this larger area with HST resolution.  COMBO-17 was mainly carried out to study the evolution of galaxies and their associated dark matter haloes at $z\\la 1$ as well as the evolution of quasars at $1\\la z\\la 5$. In order to obtain large samples of objects, four fields with a total area of $\\sim 1~\\sq\\degr$ were observed with a 17-band filter set covering the range of $\\lambda_{\\mathrm obs} \\approx 350 \\ldots 930$~nm. This provides {\\it very-low-resolution spectra} which allow a reliable classification into stars, galaxies and quasars as well as accurate photometric redshifts.   This paper publishes the full COMBO-17 catalogue \\footnote{Catalogue and images are available at CDS and at the COMBO-17 website, http://www.mpia.de/COMBO/combo\\_index.html} on the CDFS with astrometry and 17-filter photometry \\footnote{In this paper, magnitudes are always used with reference to Vega as a zero point.} of 63,501 objects found on an area of $31\\farcm5 \\times 30\\arcmin$. We also include classification, photometric redshifts and restframe luminosities \\footnote{Throughout the paper we use $H_0 = h~\\times$ 100~km/(s~Mpc) in combination with $(\\Omega_m,\\Omega_\\Lambda)=(0.3,0.7)$ and $h=1$ for luminosity distances and restframe absolute magnitudes.} whereever the data permit their derivation. We believe, the classification is mostly reliable at $R\\la 24$, where the sample contains $\\sim 100$ QSOs, $\\sim 1000$ stars and $\\sim 10000$ galaxies. Wolf et al. (2001c) published an earlier version of the catalogue containing only astrometry and BVR photometry. The version published here contains the same set of objects with identical astrometry. However, after the photometry has been processed with our final procedures, we include all 17 passbands, classifications and redshifts.  Our catalogue could be used directly to analyse aspects of galaxy evolution, and some results involving more COMBO-17 fields have already been published: Wolf et al. (2003a) studied the evolution of the galaxy luminosity function by spectral type from redshift 1.2 to 0.2. Bell et al. (2004) have focussed in particular on understanding the red sequence evolution over this redshift. Accurate photometric redshifts of QSOs allowed us to observe the evolution of faint AGN from redshift 5 to 1 (Wolf et al. 2003b) and calculate luminosity functions from the largest faint and unabsorbed AGN sample to date.  Another obvious application is the selection of sub-samples for detailed observations, e.g. high-resolution spectroscopy, while relying on the knowledge of redshift and spectral type of targets. A first example drawn from this catalogue is the measurement of velocity dispersions for $z\\sim 1$ red sequence galaxies in the CDFS by van der Wel et al. (2004).  A number of weak lensing studies took particular advantage of the accurate photometric redshifts provided by COMBO-17: Kleinheinrich et al. (2004) have used galaxy-galaxy lensing to study dark matter haloes of galaxies and their dependence on observed galaxy properties. Gray et al. (2004) have discussed the correlation of galaxy properties with the underlying dark matter density field, based on a weak lensing mass map obtained by Gray et al. (2002). Brown et al. (2003) derived the shear power sepctrum and constrained cosmological parameters from weak lensing and redshift distributions in COMBO-17. Heymans et al. (2004) have later removed intrinsic alignment signals based on our photometric redshifts. Bacon et al. (2004) have constrained the growth of dark matter density fluctuations with decreasing redshift. Taylor et al. (2004) have demonstrated the benefit of accurate redshifts through discovering a background galaxy cluster in projection behind a known cluster and estimating its mass purely from 3-D lensing analysis. Of course, the newly discovered cluster could also be confirmed independently from the redshift catalogue itself.  A second purpose of this paper is to serve as a reference for the methodology of the classification and redshift estimation in COMBO-17. It is an update to the earlier and more detailed paper by Wolf, Meisenheimer \\& R\\\"oser (2001), hereafter WMR. In conjunction with WMR this paper provides a full description of the technique. We describe the performance of the classification and redshift estimation in the COMBO-17 data set as far as we can assess it at this time. We assume that our catalogue will only be useful if we provide estimates of completeness, contamination and accuracy of redshifts in the star, galaxy and quasar sample. We believe that our redshifts for galaxies are accurate within $\\sigma_z/(1+z) < 0.01$ at the bright end ($R<20$) where we also expect small outlier rates around 1\\%. Redshift errors increase towards fainter levels and exceed $\\sim$0.05 at $R>23.5$. Without NIR data faint galaxies at $z>1$ pose a tough challenge for our approach. We believe that the accuracy of our QSO redshifts is $\\sigma_z/(1+z)\\approx 0.015$ at all magnitudes where QSOs can be identified.  In the future, we might be able to test the quality of the photometric redshifts more thoroughly. The ESO team of GOODS plans a large VIMOS programme to obtain low-resolution spectra of more than 5000 galaxies in the CDFS. Such a valuable resource would allow the most systematic test of photometric redshifts from deep images to date. A large number of redshifts have already been obtained by the VIRMOS VLT Deep Survey \\citep[VVDS,][]{LeF03} team in VIMOS GTO time.  In this paper we briefly describe the COMBO-17 observations (Sect.~2) and data reduction (Sect.~3), followed by an update of our classification procedure and template choice over WMR (Sect.~4). In Sect.~5, we present the data structure of the catalogue for the CDFS, and Sect.~6 gives an overview of the object samples. In Sect.~7 we discuss completeness and redshift errors which we estimate from simulations of the survey. Finally, Sect.~8 discusses redshift errors in detail given the comparison with almost 1000 spectroscopic IDs of stars, galaxies and QSOs. Whenever we present numbers in this paper, we refer specifically to the CDFS dataset as it is published here, but whenever we talk about techniques in general, they are applied to all fields observed in COMBO-17. ",
        "conclusions": "Wolf, Meisenheimer \\& R\\\"oser (2001) showed that medium-band surveys deliver more accurate object classifications and photometric redshifts than broad-band surveys, while not consuming more telescope time. As a result, the COMBO-17 survey was started to obtain a large redshift catalogue of galaxies and AGN for evolutionary population studies, including weak lensing observations.  In this paper, we have discussed in detail the quality of the classification and photometric redshifts of galaxies and quasars in COMBO-17. We have shown that the identification of stars is complete to $R\\la 23$ (deeper for M~stars). We have demonstrated that the identification of type-1 AGN is complete above luminosities of $M_B=-21.7$ at all redshifts from 0.5 to 5. The photometric redshifts of galaxies in COMBO-17 are better than 0.01 at bright magnitudes ($R<21$) and increase with photometric noise to $\\delta_z/(1+z) \\approx 0.06$ at $R=24$. Fainter than $R=24$, COMBO-17 is not particularly useful because the medium-bands are too shallow then. We have demonstrated that we routinely obtain photometric redshifts of quasars and luminous Seyfert-1 galaxies to an accuracy of $\\delta_z/(1+z) \\approx 0.015$.  We have demonstrated that the medium-band approach indeed delivers the expected performance which motivated the survey COMBO-17. In this paper we now deliver the COMBO-17 database from one particular patch of sky to the community for public exploitation. The published database includes images and a catalogue with 63,501 objects. Classification and redshifts are typically reliable at $R<24$, where we find $\\sim 100$ quasars, $\\sim 1000$ stars and $\\sim 10000$ galaxies.  We included the Chandra Deep Field South from the very beginning in the COMBO-17 project. A multitude of deep observations was expected across a wide range of photon energies, creating a unique data set for studies of galaxy evolution. The COMBO-17 approach has allowed us to get hold of a large redshift catalogue on the field. This includes the implicit selection of $\\sim 100$ luminous type-1 AGN with photometric redshifts as accurate as $\\sim 5000$~km/sec. We are now in a position to make this catalogue available for general use and hope to feed many more dedicated follow-up studies."
    },
    "0403/astro-ph0403516_arXiv.txt": {
        "abstract": "Relativistic current sheets have been proposed as the sites of dissipation in pulsar winds, jets in active galaxies and other Poynting flux dominated flows. It is shown that the steady  versions of these structures differ from their nonrelativistic counterparts because they do not permit transformation to a de~Hoffmann/Teller reference frame, in which the electric field vanishes. Instead, their generic form is that of a true neutral sheet: one in which the linking magnetic field component normal to the sheet is absent. Taken together with Alfv\\'en's limit on the total cross-field potential, this suggests plasma is ejected from the sheet in the cross-field direction rather than along it. The maximum energy to which such structures can accelerate particles is derived, and used to compute the maximum frequency of the subsequent synchrotron radiation. This can be substantially in excess of standard estimates. In the magnetically driven gamma-ray burst scenario, acceleration of electrons is possible to energies sufficient to enable photon-photon pair production after an inverse Compton scattering event. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403502_arXiv.txt": {
        "abstract": "The Six Degree Field Galaxy Survey (6dFGS) is a spectroscopic survey of the southern sky, which aims to provide positions and velocities of galaxies in the nearby Universe. When completed the survey will produce approximately 170000 redshifts and 15000 peculiar velocities. The survey is being carried out on the Anglo Australian Observatory's (AAO) UK Schmidt telescope, using the 6dF robotic fibre positioner and spectrograph system. We present here the adaptive tiling algorithm developed to place 6dFGS fields on the sky, and allocate targets to those fields. Optimal solutions to survey field placement are generally extremely difficult to find, especially in this era of large-scale galaxy surveys, as the space of available solutions is vast (2N dimensional) and false optimal solutions abound. The 6dFGS algorithm utilises the Metropolis (simulated annealing) method to overcome this problem. By design the algorithm gives uniform completeness independent of local density, so as to result in a highly complete and uniform observed sample. The adaptive tiling achieves a sampling rate of approximately 95\\%, a variation in the sampling uniformity of less than 5\\%, and an efficiency in terms of used fibres per field of greater than 90\\%. We have tested whether the tiling algorithm systematically biases the large-scale structure in the survey by studying the two-point correlation function of mock 6dF volumes. Our analysis shows that the constraints on fibre proximity with 6dF lead to under-estimating galaxy clustering on small scales ($<$1\\Mpc) by up to $\\sim$20\\%, but that the tiling introduces no significant sampling bias at larger scales. The algorithm should be generally applicable to virtually all tiling problems, and should reach whatever optimal solution is defined by the user's own merit function. ",
        "introduction": "\\label{tilingintro}  \\begin{figure*} \\includegraphics[width=18cm]{Fig1.ps} \\caption{The 6dFGS targets show strong clustering on the sky, as can be seen in this equal--area (Aitoff projection) greyscale map of the surface density of targets. As the 6dF field covers an area of 25.5\\,deg$^2$ and has up to 150 fibres, an optimal surface density would be approximately 6 targets per deg$^2$. The large, and spatially complex, density variations about this optimum illustrate one of the major difficulties in tiling the 6dFGS.} \\label{greyscale} \\end{figure*}   The advent of large-scale spectroscopic surveys, made possible by high multiplex spectroscopic systems, has necessitated the development of automated schemes for placing survey fields (`tiles') on the sky, and allocating survey targets to those fields. Adaptive tiling schemes take into account survey and instrument characteristics and provide efficient and optimal tile placement and target allocation. The recently completed 2dF Galaxy Redshift Survey (2dFGRS) successfully utilized adaptive tiling to obtain 221414 redshifts, using a 400 fibre spectrograph with a 2\\degree\\ field of view \\citep{Colless4}. The 2dFGRS covered 2000 deg$^2$ at a median depth of $\\bar{z}=0.11$. The Sloan Digital Sky Survey (SDSS) \\citep{York} aims to observe $\\sim$10$^6$ targets with a 640 fibre system and a 3\\degree\\ field of view, and is also employing adaptive tiling \\citep{Blanton2}. The SDSS will cover $\\sim$10000\\,deg$^2$ at a depth similar to the 2dFGRS.  The 6dFGS is a redshift and peculiar velocity survey that will cover the 17000 deg$^2$ of the southern sky with $|b|>10$\\degree \\citep{Watson2,Saunders,Wakamatsu}. The survey is being carried out on the AAO's Schmidt telescope, using the 6dF automated fibre positioner and spectrograph system \\citep{Parker,Watson1}. 6dF can simultaneously observe up to 150 targets in a circular 5.7\\degree\\ field of view. Survey observations are made with two different gratings for each field. These two spectral ranges are spliced together as part of the redshifting process, resulting in single spectra that span the range from 3900\\AA\\ to 7500\\AA\\ , at a resolution of $R=1000$ at 5500\\AA\\ and a typical signal-to-noise ratio of $S/N\\sim10$.  The goals of the survey are to map the positions and velocities of galaxies in the nearby Universe, providing new constraints on cosmological models, and a better understanding of the local populations of normal galaxies, radio galaxies, AGN and QSOs \\citep{Saunders}. The primary targets for the redshift survey are 113988 $K_s$-selected galaxies from the 2MASS near-infrared sky survey (\\citep{Jarrett}; {\\tt{http://www.ipac.caltech.edu/2mass/releases/allsky}}) down to $K_{tot}<12.75$ and with a median redshift $\\bar{z}=0.05$. The total magnitudes are estimated from the 2MASS isophotal $K_{20}$ magnitudes and surface brightness profile information \\citep{Jones}. Merged with the primary sample are 16 other smaller extragalactic samples, including targets selected from the HIPASS HI radio survey \\citep{Koribalski}, the ROSAT All Sky Survey of X-ray sources (\\cite{Voges1,Voges2}; {\\tt{http://heasarc.gsfc.nasa.gov/docs/rosat/ass.html}}), the IRAS Faint Source Catalogue (\\citep{Moshir}; {\\tt{http://irsa.ipac.caltech.edu/IRASdocs/iras.html}}), the DENIS near-infrared survey \\citep{Epchtein}, the SuperCosmos $b_J$ and $r_F$ optical catalogues \\citep{Miller}, the Hamburg-ESO QSO survey \\citep{Wisotzki} and the NVSS radio survey \\citep{Condon}. In total the survey will produce approximately 170000 redshifts.  The 6dFGS peculiar velocity survey will consist of all early-type galaxies from the primary redshift survey sample that are sufficiently bright to yield precise velocity dispersions. These galaxies are observed at higher signal-to-noise ratio ($S/N>25$), in order to obtain velocity dispersions to an accuracy of 10\\%. Peculiar velocities will be obtained using the Fundamental Plane for early-type galaxies \\citep{Djorgovski,Dressler} by combining the velocity dispersions with the 2MASS photometry. Based on the high fraction of early-type galaxies in the $K_s$ sample and the $S/N$ obtained in our observations to date, we expect to measure distances and peculiar velocities for 10--15000 galaxies out to distances of at least \\mbox{$cz=15000$\\kms}.  Observations have so far been made for $40\\%$ of the survey fields and completion is expected mid--2005. The data is non-proprietary and an Early Data Release for some 14000 objects can be accessed at {\\tt{http://www-wfau.roe.ac.uk/6dFGS/}}.  This paper describes the adaptive tiling algorithm developed for the 6dFGS. It is organised in the following manner: \\S2 outlines the functional requirements for the tiling algorithm and the context in which it was developed; \\S3 gives a detailed explanation of the algorithm; \\S4 outlines the process of parameter selection and application of the algorithm to the 6dFGS catalogue; \\S5 presents an investigation of possible systematic effects introduced by the tiling, and their impact on subsequent analyses of survey data; \\S6 concludes with a summary of the tiling algorithm and its performance.   \\begin{table} \\caption{The distribution of 6dFGS targets in terms of the numbers of neighboring targets within the fibre-button proximity exclusion limit. Only 60\\% of the catalogue are without close neighbours (as compared with $\\sim$90\\% in the SDSS), meaning a significant proportion have multiple close neighbours, the most extreme being one target with 40 neighbours within 5.71\\,arcmin.} \\begin{tabular}{crc} \\# Neighbours  & \\# Targets & Sample fraction \\\\ 0      & 102252 & 59.2\\% \\\\ 1      &  43196 & 25.0\\% \\\\ 2      &  15695 &  9.1\\% \\\\ $\\ge3$ &  11604 &  6.7\\% \\\\  \\label{neighbours} \\end{tabular} \\end{table} ",
        "conclusions": "\\label{conc}  Utilizing an optimization method based on simulated annealing, we have successfully developed an adaptive tiling algorithm to optimally place 6dFGS fields on the sky, and allocate targets to those fields. The algorithm involves a four-stage process: (i)~establishing individual target weights based on target surface density and sample observational priorities; (ii)~creating a database of all possible conflicts in allocating neighbouring targets closer than the radius of a 6dF fibre button; (iii)~creating an initial tiling by centering tiles on randomly selected targets, and then allocating targets to those tiles in order of decreasing numbers of neighbours and increasing separation from tile centres; (iv)~and finally, using the Metropolis method in randomly shifting the position of tiles, and then reallocating targets, to maximise the objective function of the tiling and hence provide an optimal tiling solution.  In order to maximise the uniformity of sampling of the 6dFGS targets, we weight inversely with the surface density of 2MASS $K_s$ galaxies. Our results showed this gave us superior uniformity when compared with a simple uniform density weighting scheme, most noticeably in reducing the number of targets not allocated to tiles along the edges of the survey volume.  Despite the challenges of highly clustered targets and large fibre buttons, tiling solutions generated using the algorithm are highly complete and uniform, and employ an efficient use of tiles. The tilings consistently give sampling rates of around 95\\%, with variations in the uniformity of sampling of less than 5\\%. Tiles typically have more than 90\\% of their available fibres allocated to targets. The algorithm has also proved itself highly flexible, able to perform on highly irregularly shaped distributions of targets.  An analysis of the two-point correlation function, calculated from 6dF mock volumes tiled with the algorithm, revealed that the constraint on fibre proximity due to the large size of the fibre buttons produces a significant under-sampling of close pairs of galaxies on scales of 1 \\Mpc\\ and smaller; on larger scales, however, the tiling algorithm does not lead to any detectable sampling bias."
    },
    "0403/astro-ph0403028_arXiv.txt": {
        "abstract": "{ In comparison with other well studied star formation regions, Taurus is unusual in several respects. (i) Its stellar initial mass function (IMF) peaks at relatively high mass ($\\sim 0.8 M_\\odot$), but contains very few stars much more massive than $1 M_\\odot$, and is relatively deficient in brown dwarfs. (ii) It has a higher binary fraction, particularly at large separations. (iii) Its core mass function is strongly peaked at a few $M_\\odot$, and the cores have extended envelopes and relatively low levels of turbulence.  \\hspace{0.25cm} We present here the results of an ensemble of hydrodynamic simulations which suggest that the unusual stellar IMF in Taurus is a direct consequence of the unusual properties of the cores there. By following the collapse and fragmentation of cores having properties typical of Taurus, we find that roughly 50\\% of the objects formed in a core, predominantly the low-mass ones, are ejected from the core to form a population of low-mass stars and brown dwarfs with a flat mass function. The remaining objects form multiple systems within the core, accreting until their masses approach $1 M_\\odot$; this produces a population of intermediate-mass stars whose mass function peaks at $\\sim 0.8 M_\\odot$. Together these two populations reproduce the IMF in Taurus very well. This demonstrates, for the first time, a direct causal link between the core mass function and the stellar IMF in a star formation region. ",
        "introduction": "Stars form within dense molecular cores (e.g. Andr\\'{e} et al., 2000), and the densest and most centrally condensed cores, i.e. those closest to forming stars, are known as prestellar cores (Ward-Thompson et al., 1994). In Ophiuchus (Motte, Andr\\'e \\& Neri, 1998), Serpens (Testi \\& Sargent, 1998) and Orion (Motte et al., 2001), the mass function of pre-stellar cores is remarkably similar to the IMF for field stars and clustered star formation regions. This suggests that there is a simple mapping from the core mass function into the stellar IMF, with the mean masses of the stars forming within a core being proportional to the core mass. However, the details of how this works are still uncertain.  In this paper we present numerical simulations of core collapse and fragmentation which demonstrate a causal link between the core mass function and the stellar IMF in the Taurus star formation region. We choose to study the Taurus region, because it is nearby and extended on the sky, and has therefore been studied in detail. Molecular-line mapping has yielded estimates of the Taurus core mass function (Onishi et al., 2002), and deep surveys of its stellar content have revealed the Taurus IMF down to the deuterium burning limit (Luhman et al., 2003a). ",
        "conclusions": "Taurus has an unusual stellar initial mass function and an unusual core mass function. We have modelled the hydrodynamical evolution of an ensemble of cores with masses based on the Taurus core mass function and levels of turbulence based on those observed in Taurus. We find that the unusual stellar IMF in Taurus can be explained as a direct result of the unusual core mass function and intrinsic core properties in Taurus.  In each core an initially unstable, multiple system forms, with between 2 and 9 members. Typically, 2 or 3 objects are ejected before they can accrete a significant amount of material. These ejected low-mass stars and brown dwarfs constitute the flat, low-mass `tail' of the Taurus IMF. There are almost twice as many low-mass stars ($0.08 M_\\odot < M < 0.5 M_\\odot$) as brown dwarfs ($M < 0.08 M_\\odot$). The ejection velocities ($\\sim 1\\,-\\,2\\,{\\rm km}\\,{\\rm s}^{-1}$) are essentially independent of mass, so the low-mass stars are as dispersed as the brown dwarfs and there is no significant mass segregation.  The remaining objects stay near the centre of the core and continue to accrete until their masses are $\\sim 0.8 M_{\\odot}$, by which stage there is not much mass left to accrete. These more massive stars constitute the Gaussian peak centred at $\\sim 0.8 M_\\odot$ in the Taurus IMF. Almost all of them are in binary or triple systems. Thus the simulated cluster of cores has an IMF very similar to that of Taurus, viz. a narrow Gaussian peak at $\\sim 0.8 M_\\odot$, and a flat tail at lower masses, extending into the brown dwarf regime.  Further support for the hypothesis that the form of the IMF in Taurus is a direct result of the typical masses of the cores in Taurus comes from Goodwin et al. (in preparation) who find that a core with low levels of turbulence typically forms a number of objects approximately equal to the number of initial Jeans masses in the core (the initial Jeans mass being roughly $1 M_{\\odot}$).  Thus lower-mass cores form fewer objects and are unable to populate the tail of the IMF as they do not eject significant numbers of brown dwarfs and low-mass stars.  Higher-mass cores over-populate the tail through the ejection of more brown dwarfs and low-mass stars.  In addition the stars that remain within a core have a larger resovoir of gas to accrete and become larger than the observed $0.8 M_{\\odot}$ peak.  The agreement between the two histograms of Figure 3 (observation and theory) is remarkably close. Hence we believe we have shown, for the first time, a direct causal link between the core mass function of a star-forming region and the stellar IMF produced in that region."
    },
    "0403/astro-ph0403672_arXiv.txt": {
        "abstract": " ",
        "introduction": "The aim of this article is to demonstrate the useful role that can be played by spectropolarimetric observations of young and evolved emission line stars that analyse the linearly polarized component in their spectra.  At the time of writing, this demonstration has to be made on the basis of optical data since there is no common-user infrared facility, in operation, that offers the desired combination of spectral resolution and sensitivity.  If the new ESO instrument, CRIRES, can be characterised to sufficient precision, it may become the first capable of performing such science.  The case for high signal-to-noise (S/N) spectropolarimetry on the largest available telescopes has been made before at an earlier ESO conference (\\cite{Schmid}, \\cite{Donati}).  Because the polarised fraction of the light received from astrophysical sources is often only of the order of 1\\% or so, it is clear that observations seeking to separate out and characterise this component suffer at least a 5-magnitude disadvantage relative to total light spectroscopy: expressed in S/N terms, one can only begin to achieve anything at S/N approaching 1~000.  It was shown by both Schmid et al~\\cite{Schmid} and Donati et al~\\cite{Donati} that some of the most exciting results, particularly from the analysis of circular-polarised light, will only come as the achievable S/N rises to 100~000.  Here we focus on what can be learned from linear spectropolarimetry alone at reasonably high spectral resolution (R of order 20000 is already valuable) and at $10^3 < $S/N$ < 10^4$.  And we remind that the near infrared (1$\\umu$m--2$\\umu$m) has the potential to out-perform the optical as a domain to work in because of the greatly reduced interstellar obscuration at these wavelengths. ",
        "conclusions": ""
    },
    "0403/astro-ph0403391_arXiv.txt": {
        "abstract": "We present the results from an automated search for damped \\lya\\ (DLA) systems in the quasar spectra of Data Release 1 from the Sloan Digital Sky Survey (SDSS-DR1). At $z\\approx 2.5$, this homogeneous dataset has greater statistical significance than the previous two decades of research. We derive a statistical sample of \\ndla\\ damped \\lya\\ systems ($>50$ previously unpublished) at $z>2.1$ and measure H\\,I column densities directly from the SDSS spectra. The number of DLA systems per unit redshift is consistent with previous measurements and we expect our survey has $>95\\%$ completeness. We examine the cosmological baryonic mass density of neutral gas $\\Omega_g$ inferred from the damped \\lya\\ systems from the SDSS-DR1 survey and a combined sample drawn from the literature. Contrary to previous results, the $\\Omega_g$ values do not require a significant correction from Lyman limit systems at any redshift. We also find that the $\\Omega_g$ values for the SDSS-DR1 sample do not decline at high redshift and the combined sample shows a (statistically insignificant) decrease only at $z>4$. Future data releases from SDSS will provide the definitive survey of DLA systems at $z\\approx 2.5$ and will significantly reduce the uncertainty in $\\Omega_g$ at higher redshift. ",
        "introduction": "It has now been two decades since the inception of surveys for high redshift galaxies through the signature of damped \\lya\\ (DLA) absorption in the spectra of background quasars \\citep{wolfe86}. Owing to large neutral hydrogen column densities $\\N{HI}$, these absorption lines exhibit large rest equivalent widths ($W_\\lambda > 10$\\AA) and show the Lorentzian wings characteristic of quantum mechanic line-damping. Through dedicated surveys of high and low redshift quasars with optical and ultraviolet telescopes, over 300 damped \\lya\\ systems have been identified.  These galaxies span redshifts $z=0$ (the Milky Way, LMC, SMC) to $z=5.5$ where the opacity of the \\lya\\ forest precludes detection \\citep{songaila02}.  Statistics of the DLA systems impact a wide range of topics in modern cosmology, galaxy formation, and physics. These include studies on the chemical enrichment of the universe in neutral gas \\citep{pettini94,pro03b}, nucleosynthetic processes \\citep{lu96,phw03}, galactic velocity fields \\citep{pw97}, the molecular and dust content of young galaxies \\citep{vladilo98,ledoux03}, star formation rates \\citep{wpg03}, and even constraints on temporal evolution of the fine-structure constant \\citep{webb01}. Perhaps the most fundamental measurement from DLA surveys, however, is the evolution of the cosmological baryonic mass density in neutral gas \\ohi\\ \\citep[][ hereafter PMSI03]{storrie00,rao00,peroux03}. Because the DLA systems dominate the mass density of neutral gas from $z=0$ to at least $z=3.5$, a census of these absorption systems determines directly $\\Omega_g$. These measurements express global evolution in the gas which feeds star formation \\citep{pei95,mathlin01} and are an important constraint for models of hierarchical galaxy formation \\citep[e.g.][]{spf01,nagamine04a}.  The most recent compilation of damped \\lya\\ systems surveyed in a `blind', statistical manner combines the effects of observing programs using over 10 telescopes, 10 unique instruments, and the data reduction and analysis of $\\approx 10$ different observers (PMSI03). In short, the results are derived from a heterogeneous sample of quasar spectra derived from heterogeneous quasar surveys. While considerable care has been paid to collate these studies into an unbiased analysis, it is difficult to assess the completeness and potential selection biases  of the current sample. These issues are particularly important when one aims to address the impact of effects like dust obscuration \\citep{ostriker84,fall93,ellison01}.   In this paper we present the first results in a large survey for damped \\lya\\ systems drawn from a homogeneous dataset of high $z$ quasars with well-defined selection criteria. Specifically, we survey the quasar spectra from Data Release 1 of the Sloan Digital Sky Survey  (SDSS-DR1) restricting our search to SDSS-DR1 quasars with Petrosian magnitude $r' < 19.5$\\,mag. The DR1 sample alone (the first of five data releases from SDSS) offers a survey comparable to -- although not strictly independent from -- the efforts of 20 years of work.  We introduce algorithms to automatically identify DLA candidates in the fluxed (i.e.\\ non-normalized) quasar spectra and perform Voigt profile analyses to confirm and analyze the DLA sample.  This survey was motivated by a search for `metal-strong' DLA systems like the $z$=2.626 damped \\lya\\ system toward FJ$0812+32$ \\citep{phw03}.  A discussion of the `metal-strong' survey will be presented in a future paper (Herbert-Fort et al.\\ 2004, in preparation).  This paper is organized as follows.  In $\\S$~2, we present the quasar sample and discuss the automatic DLA candidate detection. In $\\S$~3, we present the Voigt profile fits to the full sample. We present a statistical analysis in $\\S$~4 and a summary and concluding remarks are given in $\\S$~5.  \\clearpage ",
        "conclusions": "In this paper, we have introduced an automated approach for identifying DLA systems in the SDSS quasar database.  We have applied our method to the Data Relase 1 quasar sample and have identified a statistical sample of \\ndla\\ DLA systems including $>50$ previously unpublished cases. Remarkably, the SDSS Data Release 1 exceeds the statistical significance of the previous two decades of DLA research at $z \\approx 2.5$.  More importantly, this sample was drawn from a well defined, homogeneous dataset of quasar spectroscopy. We present measurements of the number per unit redshift $n(z)$ of the DLA population and the contribution of these systems to the cosmological baryonic mass density in neutral gas $\\Omega_g$.  Although the SDSS-DR1 sample does not offer a definitive assessment of either of these quantities, future SDSS data releases will provide a major advancement over all previous work.  Our measurements of $n(z)$ are consistent with previous results suggesting a high completeness level for our DLA survey of the SDSS-DR1. We find $\\Omega_g$ increases with redshift to at least $z=3$ and is consistent with increasing to $z=4$ and beyond. This latter claim, however, is subject to significant uncertainty relating to sample size. Perhaps the most important result of our analysis is that the full DLA sample no longer shows significantly fewer DLA systems with large $\\N{HI}$ at $z>3.5$.  This contradicts the principal result of PMSI03 from their analysis of the pre-SDSS DLA compilation.  Apparently, their maximum likelihood approach failed to adequately assess uncertainty related to sample size.  With the inclusion of only 6 new DLA, we no longer find that Lyman limit systems with $\\N{HI} < 2 \\sci{20} \\cm{-2}$ are required in an analysis of $\\Omega_g$.  Before concluding, we offer several additional criticisms of the PMSI03 analysis and the role of sub-DLA systems. First, these authors assumed a three parameter $\\Gamma$-function for the column density frequency distribution of absorption systems with $\\N{HI} > 10^{17.2} \\cm{-2}$, \\\\ $f(N) = (f_*/N_*) (N/N_*)^{-\\beta} {\\rm e}^{-N/N_*}$. Although this function gives a reasonable fit to the column density frequency distribution of the DLA systems, it is not physically motivated\\footnote{In fact this curve does not smoothly connect to the power-law derived for quasar absorption lines with $\\N{HI} < 10^{17.2} \\cm{-2}$.} and, more importantly, places much greater emphasis on sub-DLA than other functions (e.g.\\ a broken power-law). Future assessments must include other functional forms to examine this systematic uncertainty. Second, the authors did not fit for the normalization of the distribution function $f_*$.  The uncertainty in this parameter could easily contribute an additional $>50\\%$ to the error budget. Third, their treatment did not account for sample variance; the uncertainties these authors reported were severe underestimates. Finally (and perhaps most importantly), a recent analysis of a sub-DLA sample by \\cite{dessauges03} has shown that these absorption systems have very high ionization fractions (see also Howk \\& Wolfe 2004 in preparation).  Although these absorption systems may ultimately make an important contribution to the total H\\,I mass density of the universe, they are intrinsically different from the DLA systems. Indeed, a more appropriate title for this sub-set of Lyman limit systems is the `super-LLS'. This gas -- in its present form -- cannot contribute to star formation and is unlikely to be directly associated with galactic disks or the inner regions of protogalactic `clumps'.  Any interpretation of results related to the super-LLS must carefully consider these points \\citep[e.g.][]{maller03,peroux03}."
    },
    "0403/astro-ph0403358_arXiv.txt": {
        "abstract": "Explaining the origin of the orbit of 2000~CR$_{105}$ ($a\\sim 230$~AU, $q\\sim 45$~AU) is a major test for our understanding of the primordial evolution of the outer Solar System.  Gladman et al$.$~(2001) showed that this objects could not have been a normal member of the scattered disk that had its perihelion distance increased by chaotic diffusion.  In this paper we explore four seemingly promising mechanisms for explaining the origin of the orbit of this peculiar object: (i) the passage of Neptune through a high-eccentricity phase, (ii) the past existence of massive planetary embryos in the Kuiper belt or the scattered disk, (iii) the presence of a massive trans-Neptunian disk at early epochs which exerted tides on scattered disk objects, and (iv) encounters with other stars.  Of all these mechanisms, the only one giving satisfactory results is the passage of a star. Indeed, our simulations show that the passage of a solar mass star at about 800~AU only perturbs objects with semi-major axes larger than roughly 200~AU to large perihelion distances.  This is in good agreement with the fact that 2000~CR$_{105}$ has a semi-major axis of 230~AU and no other bodies with similar perihelion distances but smaller semi-major axes have yet been discovered.  The discovery of 2003 VB$_{12}$, ($a=450$~AU, $q=75$~AU) announced a few days before the submission of this paper, strengthen our conclusions. ",
        "introduction": "\\label{intro}  The trans-Neptunian population of small bodies is usually divided in two categories, the Kuiper belt and the scattered disk, although the partition between the two is not precisely defined.  In Morbidelli et al. (2003) we have introduced a partitioning based on the dynamics of orbits in the current Solar System. We called {\\it scattered disk} the region of the orbital space that can be visited by bodies that have encountered Neptune within a Hill's radius at least once during the age of the Solar System, assuming no substantial modification of the planetary orbits.  We then called {\\it Kuiper belt} the complement of the scattered disk in the $a>30$~AU region.  The bodies that belong to the scattered disk in this classification scheme do not provide us with any significant clue about the primordial architecture of the Solar System.  This is because their current orbits can be achieved by purely dynamical evolution in the current planetary system from objects that started in nearly-circular nearly-coplanar orbits in Neptune's zone. The opposite is true for the orbits of the Kuiper belt objects.  All bodies in the Solar System must have been formed on orbits typical of an accretion disk (e.g. with very small eccentricities and inclinations).  Therefore, the fact that most Kuiper belt objects have a non-negligible eccentricity and/or inclination reveals that some excitation mechanism, which is no longer at work, was active in the past.  In this respect, particularly interesting are the Kuiper belt bodies with large semi-major axis ($a>50$~AU), such as 2001~QW$_{297}$ ($a=51.3$~AU, $q=39.5$AU, $i=17.1^\\circ$), 2000~YW$_{134}$ ($a=58.4$~AU, $q=41.2$AU, $i=19.8^\\circ$), 1995~TL$_8$ ($a=52.5$~AU, $q=40.2$AU, $i=0.2^\\circ$), 2000~CR$_{105}$ ($a=230$~AU, $q=44.4$~AU, $i=22.7^\\circ$) and, last discovered, 2003~VB$_{12}$ ($a=531$~AU, $q=74.4$~AU, $i=11.9^\\circ$; Brown et al$.$~2004).  We call these objects {\\it extended scattered disk} objects for three reasons: {\\it (i)} they do not seem to belong (some caution is needed because of the uncertainties in the orbital elements of these objects) to the scattered disk according to our definition but are very close to its boundary (Gladman et al., 2001; Emel'yanenko et al., 2003; Morbidelli et al., 2004); {\\it (ii)} some of these bodies have sizes of several hundred kilometers, suggesting that they formed much closer to the Sun, where the accretion timescale was sufficiently short (Stern, 1996) and were subsequently transported to these current locations; {\\it (iii)} the lack of objects with $q>41$~AU and $50<a<200$~AU cannot be due to observational biases (given that many classical belt objects have been discovered up to distances of 45--50~AU), suggesting that these extended scattered disk objects are {\\it not} the high-eccentricity members of an excited belt beyond 50~AU.  These considerations indicate that in the past the scattered disk extended beyond its present boundary in perihelion distance.  Perhaps that most promising idea for the formation of these extended scattered disk objects (or at least most of them) was recently studied by Gomes (2003b). Gomes (2003b) investigated whether the scenario proposed in Gomes (2003a) for the origin of the `hot Kuiper belt population' (the population of non-resonant bodies with large inclinations) could also be responsible for the extended scattered disk. We remind the reader that in Gomes's scenario, the hot population was originally a part of the primordial, massive scattered disk population.  During Neptune's migration, a small fraction of these objects had their perihelion distances increased and thus they became permanently trapped on stable orbits. Gomes (2003b) found several particles that increased perihelion distance well beyond $40$~AU in the $a>50$~AU region.  However, in all cases, the semi-major axis was smaller than 200~AU.  Therefore, Gomes's mechanism implies the existence of bodies with $q\\!sim\\!45\\,$AU spread from $a\\!\\sim\\!50\\,$AU to $\\sim\\!200\\,$AU, but no objects has been discovered at the small semi-major axis end of this range.  This, in spite of the fact that observational biases favor the discovery of small semi-major axis objects.  Indeed, 2000~CR$_{105}$ is special for a couple of reasons.  Until the recent discovery of 2003~VB$_{12}$, it had the largest semi-major axis of the extended scattered disk, by a large margin.  It also had a significantly larger perihelion distance than any other extended scattered disk object.  Although it is possible that 2000~CR$_{105}$ is just an outlaying member of the extended scattered disk, the fact that no objects with perihelion distance comparable to that of 2000~CR$_{105}$ but with a smaller $a$ had been discovered seemed significant to us.  This is particularly true considering that observational biases sharply favors the discovery of objects with smaller semi-major axes.  Thus, we were motivated to look for dynamical mechanisms that preferentially raised the perihelion distance of scattered disk objects at large semi-major axis. The discovery of 2003~VB$_{12}$ came a few days before the submission of this paper, and confirmed that our investigation was well motivated.  In fact, the orbit of this body definitely falls beyond the distribution produced in Gomes model.  Some of the mechanisms investigated in this paper have been already suggested by Gladman et al. (2001), but never have been quantitatively explored.  In Section~\\ref{nep_ecc} we consider the case where Neptune was more eccentric in the past, as proposed by Thommes et al$.$~(1999).  It is obvious that a more eccentric Neptune would produce a extended scattered disk, but it is not known, {\\it a priori}, what eccentricity would be required to produce objects on 2000~CR$_{105}$-like orbits, and over what timescale.  In Section~\\ref{planet} we investigate the effects of the presence of terrestrial mass planet(s) in the Kuiper belt or in the scattered disk, as proposed by Morbidelli and Valsecchi~(1997) and Brunini and Melita~(2002).  In Section~\\ref{dtides} we propose a new model for the origin of 2000~CR$_{105}$, in which the tides raised by a massive disk beyond $\\sim 70$~AU increased the perihelion distance of high inclined scattered disk objects.  Finally, in Section~\\ref{senc} we investigate the stellar passage scenario.  This scenario has been first proposed by Ida et al. (2000) to explain the structure of the inner Kuiper belt.  Although we disagree that the all of the sculpting of the Kuiper belt could be due this mechanism (see Levison et al$.$, 2004), it is still possible that a more gentle encounter could have formed objects like 2000~CR$_{105}$. ",
        "conclusions": "We have analyzed with numerical simulations four seemingly promising mechanisms for explaining the origin of the peculiar extended scattered disk object 2000~CR$_{105}$: (i) a high eccentricity phase of Neptune, (ii) the existence of a rogue planet in the Kuiper belt or in the scattered disk, (iii) the tide exerted by a massive and dynamically cold trans-Neptunian disk, and (iv) the passage of a star near the Solar System.  Of these, only the early passage of a Solar-mass star at about 800~AU from the Sun appears satisfactory. This is also the only scenario that we have studied that can easily explain the origin of the newly found object 2003~VB$_{12}$.  Another scenario for the origin of extended scattered disk objects has been proposed by Gomes (2003a, 2003b).  In it, a small fraction of the objects in an early massive scattered disk population permanently acquire a large perihelion distance during the outer migration of Neptune.  This mechanism predicts that there should be large-$q$ objects with semi-major axis all over the range from $\\sim\\!50$ to $\\sim\\!200\\,$AU.  The fact that the extended scattered disk bodies with the largest perihelion distances, 2000~CR$_{105}$ and 2003~VB$_{12}$, both have $a>200$~AU, argues against a scenario like that of Gomes.  Since observational biases (given an object's perihelion distance and absolute magnitude, and a survey's limiting magnitude of detection) sharply favors the discovery of objects with small semi-major axis, we believe that it would be unlikely that the first two discovered body with $q>44$~AU had $a\\!>\\!200$~AU, if the real semi-major axis distribution in the extended scattered disk were skewed toward small $a$.  An yet, Gomes's mechanism and all of the mechanisms we have studied, except for the passing star, lead to such a distribution. Indeed, it is intuitive that creating a distribution where the objects with the largest $a$ also have the largest $q$ requires a perturbation `from the outside', whose magnitude decreases with decreasing heliocentric distance.  Thus, there are not many alternatives to the stellar encounter scenario.  In the current galactic environment, the closest stellar encounter that should occur over the age of the Solar System is at $\\sim 900$~AU (Garcia-Sanchez et al., 2001). This should typically happen with a star that is about 1/10th the mass of the Sun.  Therefore, the encounter with a Solar-mass star at $\\sim 800$~AU at early times requires that the environment in which that Sun formed was significantly denser, such as that of a stellar cluster (Bate et al., 2003).  In the future the existence of 2000~CR$_{105}$, 2003~VB$_{12}$, and their cohorts my supply important clues to the exact environment in which the Sun and Solar system formed.  On a final note, since its discovery, passing stars have been used to capture objects into the Oort cloud (Oort 1950).  Indeed, the process invoked for the Oort cloud is identical to the one employed here. Thus, if a passing star is indeed responsible for the formation of objects like 2000~CR$_{105}$ and 2003~VB$_{12}$, then it is perhaps more appropriate to characterize these objects as the inner edge of the Oort cloud rather than the outer edge of the scattered disk.  Indeed, recent simulations of the Oort cloud in a star cluster (Eggers~1997; 1998; Fernandez and Brunini~2000) actually produce objects like 2003~VB$_{12}$."
    },
    "0403/astro-ph0403444_arXiv.txt": {
        "abstract": "{Combining precise $B$,$V$ photometry and radial velocities, we have been able to derive a firm orbital solution and accurate physical parameters for the newly discovered eclipsing binary HD~23642 in the Pleiades open cluster. The resulting distance to the binary and therefore to the cluster is 132$\\pm$2~pc. This closely confirms the distance modulus obtained by classical main sequence fitting methods ($m - M$=5.60 or 132 pc), moving cluster techniques and the astrometric orbit of Atlas. This is the first time the distance to a member of the Pleiades is derived by orbital solution of a double-lined eclipsing binary, and it is intended to contribute to the ongoing discussion about the discordant Hipparcos distance to the cluster. ",
        "introduction": "The distance to the Pleiades has been derived by main-sequence fitting methods several times over the years, and consistently found to cluster around $m - M$=5.60~$\\pm$0.04 (e.g. Turner 1979, Meynet et al. 1993), corresponding to a distance of 132~($\\pm$2) pc. Values of $m - M$=5.52 (127~pc) by Mitchell \\& Johnson (1957) and $m - M$=5.75 (141~pc) by Eggen (1950) bracket the range of published distance moduli. Therefore, it came as a surprise the shorter distance of 116~$\\pm$3.3~pc ($m - M$=5.33~$\\pm$0.06) derived from the parallaxes of 54 Pleiades members when the results of the Hipparcos astrometric mission became available (van Leeuwen \\& Hansen Ruiz 1997). The $\\sim$0.3~mag difference has far reaching consequences in many areas of astrophysics, and it soon prompted extensive observational and theoretical work to account for it.  \\begin{figure*}[!t] \\centerline{\\psfig{file=Gb231_f1.ps,width=16.7 cm,angle=270}} \\caption[]{Our photo-electric $V$ (top) and $B$ (bottom) observations of HD~23642 are phase plotted in the upper panel following the ephemeris of Table~2.  The lower panel displays our radial velocities from Table~1 (filled circles) and those from Pearce (1957, squares) and Abt (1958, crosses).  The curves over plot the orbital solution given in Table~2.} \\end{figure*}  An anomalous abundance of Helium was discussed as a possible explanation of the large difference between Hipparcos and ground-based photometric distances to the Pleiades (Mermilliod et al. 1997a). The required helium over-abundance ($Y$=0.37) is however too large to be a feasible explanation according to Pinsonneault et al. (1998). A sub-solar metallicity could reconcile the distances, but the metallicity of the Pleiades has been measured several times and with different methods, which generally converged to a mean value of [Fe/H]=0.00$\\pm$0.03 (e.g. Stauffer et al. 2003). Castellani et al. (2002) were able to obtain a good fit of the Pleiades main sequence with their theoretical isochrones (with updated physical inputs) and the Hipparcos distance, employing a sub-solar metallicity of [Fe/H]=--0.15. Working differentially with a sample of nearby G and K stars, Percival et al. (2003) presented a re-calibration of photometric metallicity indexes and argued that they support a sub-solar metallicity of the Pleiades in spite of the spectroscopic evidences (e.g. Boesgaard \\& Friel 1990) for a solar one, therefore removing the difference between main-sequence fitting and Hipparcos distances. However, Stauffer et al. (2003) remarked about the anomalous blue colors of K stars in the Pleiades, that they ascribed to fast rotation and spotted surfaces as a consequence of the young age of the cluster.  Narayanan \\& Gould (1999) derived a distance modulus to the Pleiades of $m - M$=5.58~($\\pm$0.18) using a variant of the moving cluster method, the gradient in the radial velocity of the cluster members in the direction of the proper motion of the cluster. Their techniques relies on the assumption that the velocity structure of the Pleiades is not significantly affected by rotation. Narayanan \\& Gould concluded that the errors in the Hipparcos parallaxes toward Pleiades are spatially correlated over angular scales of a few degrees, with an amplitude up to 2~mas. Almost simultaneously, however, van Leeuwen (1999) compared the results on 9 clusters and concluded instead that the Hipparcos parallaxes of the Pleiades were basically unaffected by systematics and that the distance modulus to the latter is $m - M$=5.37~($\\pm$0.07). Makarov (2002) suggested that, while Hipparcos parallaxes are overall of excellent quality, those for the Pleiades suffered from data analysis procedures that could have introduced systematics when a rich and bright cluster was crossing one of the two Hipparcos fields of view while the other was essentially deprived of stars, as in the case of Pleiades. From intermediate Hipparcos data and under some assumptions, Makarov has recomputed the distance to the Pleiades as 129~$\\pm$3~pc. F. van Leeuwen (private communication) is currently working on a further iteration on original Hipparcos reductions that should assess and quantify is any systematics affect the Pleiades data.  The question about the distance to the Pleiades has not yet been firmly solved, and to tackle it new, robust and independent approaches (as much geometric as possible) are required. A first one, advocated by Pacz\\'{y}nski (2003), combines astrometric and spectroscopic observations of Atlas (HD~23850), one of the brightest members of the Pleiades, which is an astrometric binary with a period of 291~day, a semi-major axis of 12.9~mas and 0.246 eccentricity. Pan et al. (2004) did not have the radial velocities, but they anyway derived the distance by combining the astrometric orbit with a mass-luminosity relation. They derived a distance of 135$\\pm$2~pc, which is in close agreement with the results of main-sequence fitting methods.  Another method involves double-lined eclipsing binaries (SB2 EB), which are the distance indicators now providing the most reliable distances to Magellanic Clouds and other galaxies in the Local Group. The recent discovery by Torres (2003) that HD~23642 in the Pleiades is an SB2 eclipsing binary (the only one so far known in the cluster) prompted us to use it to measure the distance to the Pleiades. In this {\\em Letter} we present accurate new $B$ and $V$ photoelectric photometry and radial velocities of HD~23642, and use them to derive an orbital solution and infer the distance to the star, and therefore to the cluster. ",
        "conclusions": ""
    },
    "0403/hep-th0403073_arXiv.txt": {
        "abstract": "Dirac-Born-Infeld type effective actions reproduce many aspects of string theory classical tachyon dynamics of unstable Dp-branes.  The inhomogeneous tachyon field rolling from the top of its potential forms topological defects of lower codimensions. In between them, as we show, the tachyon energy density fragments into a p-dimensional web-like high density network evolving with time. We present an analytic asymptotic series solution of the non-linear equations for the inhomogeneous tachyon and its stress energy. The generic solution for a tachyon field with a runaway potential in arbitrary dimensions is described by the free streaming of noninteracting massive particles whose initial velocities are defined by the gradients of the initial tachyon profile. Thus, relativistic particle mechanics is a dual picture of the tachyon field effective action. Implications of this picture for inflationary models with a decaying tachyon field are discussed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403608_arXiv.txt": {
        "abstract": "At $z=0.1055$, the gamma-ray burst GRB 031203 is one of the two nearest GRBs known. Using observations from the Very Large Array (VLA) and \\chandralong\\, we derive sub-arcsecond localizations of the radio and X-ray afterglow of this GRB. We present near-infrared observations of the supernova SN 2003lw, which exploded in the host galaxy of the GRB 031203. Our deep, high resolution Magellan/PANIC data establish that this SN is spatially coincident with the radio and X-ray localizations of the afterglow of GRB 031203 to sub-arcsecond precision, and is thus firmly associated with the GRB. We use image differencing to subtract the bright emission from the host galaxy, and measure the time evolution of the SN between $\\sim5$ and $\\sim50$ days after the GRB. The resulting light curve has a shape which is quite different from that of the two SNe previously associated with GRBs, SN 1998bw and SN 2003dh. With SN 2003lw securely associated with this burst, we confirm that all three GRBs with redshifts $z<0.3$ were accompanied by SN explosions. ",
        "introduction": "The emerging association between long-duration gamma-ray bursts (GRBs) and type Ic supernovae (SNe Ic) is perhaps the most significant breakthrough in our understanding of GRBs, and may also provide new insights into the physics of core-collapse SNe and the deaths of massive stars. The initial strong evidence for this connection came from the spatial and temporal coincidence between GRB 980425 and SN 1998bw (Galama et al. 1998). More recently, this picture was convincingly affirmed by the detection of SN features in the optical afterglow spectrum of GRB 030329 (Stanek et al. 2003; Hjorth et al. 2003; Matheson et al. 2003). Both of these milestone discoveries resulted from the study of the nearest GRBs yet identified, GRB 980425 at $z=0.0085$ and GRB 030329 at $z=0.1685$. Indeed, each relatively rare occurrence of a GRB at low redshift ($z<0.3$) provides a unique opportunity for further study of the GRB-SN connection.  GRB 031203\\footnotemark~was detected by IBIS on board the INTEGRAL spacecraft on 2003 December 3, at 22:01:28 UT (Gotz et al. 2003). A fading X-ray afterglow was discovered by  XMM-Newton on 2003 December 4 (Santos-Lleo et al. 2003; Campana et al. 2003; Rodriguez-Pascual et al. 2003) and found to be consistent with the location of a radio transient (Frail 2003; Soderberg, Kulkarni, \\& Frail 2003). Prochaska et al. (2003a; 2003b; 2004) identified the host galaxy of the X-ray and radio transients, determined its redshift ($z=0.1055$) and studied its properties in detail. Recently, several groups (Bersier et al. 2004; Tagliaferri et al. 2004a; Thomsen et al. 2004; Cobb et al. 2004) reported optical photometric and spectroscopic observations of this GRB, which apparently reveal the signatures of an associated SN (designated SN 2003lw; Tagliaferri et al. 2004b). In spite of the low redshift of this event, its low Galactic latitude (less than $5^{\\circ}$ from the plane) and the resulting Galactic extinction ($E[B-V]=1.04$; Schlegel, Finkbeiner, \\& Davis 1998), as well as its bright host galaxy, make the study of this event challenging.  \\footnotetext{The possible classification of this event as an X-ray flash (XRF) was debated in recent literature (e.g., Prochaska et al. 2004; Watson et al. 2004; Thomsen et al. 2004). However, the high-energy observations presented by Sazonov, Lutovinov, \\& Sunyaev (2004) conclusively show that this event does not fit any of the commonly used definitions of XRFs.}  To overcome these difficulties, we have undertaken near infrared (NIR) observations. The results of this effort are reported in this {\\it Letter}. Coordinated radio and X-ray observations, which enabled us to probe the total energy output of this sub-energetic GRB, are reported elsewhere (Soderberg et al. 2004). ",
        "conclusions": "In this {\\it Letter}, we have presented $J$-band observations of SN 2003lw. The sensitivity and resolution of our data enabled us to pinpoint the location of SN 2003lw within its host galaxy, and to show it is consistent with sub-arcsecond localizations of the radio and X-ray afterglow of GRB 031203, thus confirming the association of these two events. The precise NIR localization of this event also puts it within $0.2$~kpc from the host galaxy center. The $J$-band light curve of SN 2003lw shows a rapid initial rise ($5-7$ days after the GRB) and evidence for bright emission more than $50$ days after the GRB. The fast early rise of SN 2003dh, associated with GRB 030329, has been interpreted by Woosley \\& Heger (2003) and Mazzali et al. (2003) as evidence for asymmetry in the explosion. A thorough investigation of this possibility will probably require an analysis of our data in combination with other extensive data sets of optical and NIR photometry and spectroscopy collected by other groups (e.g., Thomsen et al. 2004; Tagliaferri et al. 2004; Bersier et al. 2004; Cobb et al. 2004).  Cobb et al (2004) have recently reported $I$ and $J$-band observations of this event, obtained with the SMARTS 1.3m telescope. A direct comparison between our observations and the SMARTS data, obtained on numerous epochs, is complicated by the fact that these authors do not present the light curve of SN 2003lw. Instead, they plot the temporal evolution of the combined light of the SN and its bright host galaxy, derived from aperture photometry, which shows considerable scatter. It is thus hard to say whether their data show the same early fast rise we detect. Our observation that SN 2003lw had similar flux levels 7 and 50 days after the GRB is consistent with the reported SMARTS data. Furthermore, comparing their $I$-band data with model light curves of SN 1998bw, these authors arrive at the conclusion that the light curve shape of SN 2003lw does not resemble that of SN 1998bw. This is in accord with our analysis of the $J$-band data (Cobb et al. 2004 do not attempt to compare their $J$-band data with a model of SN 1998bw). It appears that both our observations and the Cobb et al. (2004) data set suggest that SN 2003lw had a light curve quite unlike that of SN 1998bw; with a fast rise to maximum, which appears broader, and perhaps showing a secondary peak in the IR.  The above discussion provides further evidence for the diversity of SNe associated with GRBs [see, e.g., Thomsen et al. (2004) and Lipkin et al. (2004) for recent reviews]. With SN 2003lw, all three nearby GRBs are firmly associated with SNe, apparently supporting suggestions made, e.g., by Podsiadlowski et al (2004), that all long GRBs are accompanied by SNe. It may well be the case that the focus of future studies should now move from proving the association between SNe and GRBs to an attempt to characterize the properties of this population of SNe. Such studies may provide new clues about the progenitors and engines of GRBs, by requiring viable GRB models to be able to produce the large quantities of nickel derived from the SN observations, as well as valuable insights into possible explosion mechanisms of core-collapse SNe, which may involve GRB-like aspherical effects [e.g., Khokhlov et al. (1999), see Gal-Yam et al. (2004) for further discussion]. Low-redshift GRBs, expected to be localized in larger numbers by the upcoming {\\it SWIFT} mission, as well as systematic studies of local core-collapse SNe (e.g., Berger et al. 2003; Soderberg et al. 2004, in preparation; Stockdale et al. 2004) will probably shed more light on these intriguing questions."
    },
    "0403/astro-ph0403114_arXiv.txt": {
        "abstract": "We report on observations of the blazar W Comae (ON+231) with the Solar Tower Atmospheric Cherenkov Effect Experiment (STACEE), a wavefront-sampling atmospheric Cherenkov telescope, in the spring of 2003. In a data set comprising 10.5 hours of ON-source observing time, we detect no significant emission from W Comae. We discuss the implications of our results in the context of the composition of the relativistic jet in W Comae, examining both leptonic and hadronic models for the jet. We derive 95\\% confidence level upper limits on the flux at the level of 1.5--3.5$\\times 10^{-10}$ cm$^{-2}$ s$^{-1}$ above 100 GeV for the leptonic models, or 0.5--1.1$\\times 10^{-10}$ cm$^{-2}$ s$^{-1}$ above 150 GeV for the hadronic models. ",
        "introduction": "The current catalog of extragalactic TeV gamma-ray sources consists of blazars, which are among the brightest and most rapidly variable objects in the sky.  The unusual observational properties of blazars are usually explained in terms of relativistic bulk motion of the emitting region in a jet parallel to the line of sight \\citep{up95}.  The continuum emission of blazars is typically nonthermal and contains two broad peaks, one at lower energies (radio to X-ray) and one at higher energies (keV to TeV). Although the low-energy peak is generally assumed to result from synchrotron radiation from high-energy electrons in the jet, the mechanism for the production of the high-energy radiation is still a subject of debate, and several competing models exist.  In ``leptonic'' jet models, the high-energy radiation is produced by inverse Compton scattering of low-energy photons from a population of electrons and/or positrons; these models are favored for the well-studied blazars Mrk 421 and 501 \\citep{ca99,konopelko03}.  In the synchrotron self-Compton model (SSC), a single population of electrons and positrons produce the target photon field via synchrotron radiation, and subsequently upscatter the synchrotron photons to gamma-ray energies \\citep{bm96}.  There may also be an external Compton (EC) component of target photons from the accretion disk or ambient medium \\citep{dsm92,sbr94}.  Alternatively, in ``hadronic'' jet models, protons play a central role.  In hadronic models the high-energy radiation is attributed to photomeson processes \\citep{mannheim93} or synchrotron radiation from protons or muons, as in the synchrotron proton blazar (SPB) model \\citep{mp00,aharonian00,mucke03}. The study of hadronic models was originally motivated by the hypothesis that the highest-energy cosmic rays are produced in blazar jets \\citep{mannheim95}.  The BL Lac object W Comae (ON+231) may provide an excellent test case for hadronic jet models \\citep{bmr02}.  The transition between the low-energy and high-energy peaks in the continuum of W Comae appears clearly in X-ray data taken by the BeppoSAX satellite \\citep{tagliaferri00}. These high-quality observations of the transition region place tight constraints on leptonic models, requiring the predicted gamma-ray emission to cut off sharply above 100 GeV. In contrast, hadronic models may allow for significant emission above 100 GeV. Observations by the EGRET detector aboard the Compton Gamma Ray Observatory show a hard power law spectrum (photon spectral index $\\alpha = 1.73 \\pm 0.18$) extending up to about 10 GeV with no sign of any cutoff \\citep{hartman99}. Yet the object has not been detected at energies above 300 GeV, despite repeated observation by the Whipple 10-m instrument \\citep{horan03}. At a redshift of $0.1$, absorption of gamma rays in this energy range by pair production \\( \\gamma \\gamma \\longrightarrow e^+e^- \\) against the extragalactic background light (EBL) may be significant, but only at energies above about 500 GeV \\citep{primack99,primack01,ms01,aharonian01}, so the intrinsic emission spectrum of W Comae should be directly observable at energies lower than this.  The Solar Tower Atmospheric Cherenkov Effect Experiment (STACEE) currently operates above an energy threshold of about 100 GeV for gamma rays. W Comae has been observed in the past with previous versions of the STACEE detector \\citep{theoret-thesis}. New observations of W Comae were carried out in the spring of 2003 by the most recent version of the detector, STACEE-64. ",
        "conclusions": "The STACEE-64 observations in the spring of 2003 were made at a lower energy threshold than any other atmospheric Cherenkov telescope has yet attained for W Comae.  STACEE detects no significant emission from W Comae, resulting in 95\\% CL upper limits on the integral flux above this threshold in various hadronic emission models at the level of $10^{-10}$ cm$^{-2}$ s$^{-1}$. While leptonic models predict a flux which falls below this level, extrapolations of the best-fit EGRET power law, and some synchrotron-proton hadronic models, predict an integral gamma-ray flux above the energy threshold close to, or exceeding, the upper limit from STACEE observations. Additional STACEE observations planned for the spring of 2004 ought either to exclude these models at a significantly higher confidence level, or to detect gamma-ray emission from W Comae if these models provide an adequate description of the source."
    },
    "0403/astro-ph0403264_arXiv.txt": {
        "abstract": "Using the SPectrometer for Infrared Faint Field Imaging (SPIFFI) on the ESO VLT, we have obtained $J$, $H$, and $K$ band integral field spectroscopy of the $z = 2.565$ luminous submillimeter galaxy SMM\\,J14011+0252.  A global spectrum reveals the brighter of this spatially resolved system's two components as an intense starburst that is remarkably old, massive, and metal-rich for the early epoch at which it is observed.  We see a strong Balmer break implying a $\\geq 100\\,{\\rm Myr}$ timescale for continuous star formation, as well as nebular emission line ratios implying a supersolar oxygen abundance on large spatial scales.  Overall, the system is rapidly converting a large baryonic mass into stars over the course of only a few hundred Myr.  Our study thus adds new arguments to the growing evidence that submillimeter galaxies are more massive than Lyman break galaxies, and more numerous at high redshift than predicted by current semi-analytic models of galaxy evolution. ",
        "introduction": "Deep imaging in the optical and near-IR has made it possible to constrain both the evolution of the cosmic star formation rate density (e.g., Madau et al. 1996) and its time integral, the growth of the cosmic stellar mass density (e.g., Dickinson et al. 2003).  Short-wavelength studies give an incomplete picture of these trends, however, since an important population of high-redshift galaxies is too obscured to be easily detected.  These submillimeter galaxies (SMGs; see Blain et al. 2002 and references therein) have large IR luminosities that are predominantly generated by star formation rather than accretion \\citep{barg01,alma03}; as a result, they contribute substantially to cosmic star formation \\citep{barg00}.  Moreover, as the strikingly different appearances of the Hubble Deep Field at $0.83\\,{\\rm \\mu m}$ \\citep{will96} and $850\\,{\\rm \\mu m}$ \\citep{hugh98} exemplify, SMGs are rarer and forming stars much more intensely than typical optically selected systems.  To understand the evolutionary state of SMGs, we have conducted a detailed rest-frame optical case study using SPIFFI, the SPectrometer for Infrared Faint Field Imaging \\citep{tecz00,eise00,eise03}.  Our target is SMM\\,J14011+0252 (hereafter J14011), a $z = 2.565$ galaxy lying behind the $z = 0.25$ cluster A1835 and the second SMG to have an optical counterpart identification \\citep{barg99} validated by CO interferometry \\citep{fray99}.  Recent maps of J14011's molecular gas have led to divergent conclusions about its intrinsic size: \\citet{ivis01} suggest that a large background source undergoes a factor ${\\cal M} \\sim 2.5$ magnification by A1835 as a whole, while \\citet{down03} argue that a fortuitously located cluster member further boosts magnification of a small background source to ${\\cal M} \\sim 25$.  We defer a full discussion of lensing to a future paper; here, we focus primarily on the lensing-independent conclusions of high age and metallicity that can be drawn from J14011's global spectrum alone. Throughout the paper we assume a flat $\\Omega_\\Lambda = 0.7$ cosmology with $H_0 = 70\\,{\\rm km\\,s^{-1}\\,Mpc^{-1}}$. ",
        "conclusions": ""
    },
    "0403/astro-ph0403578_arXiv.txt": {
        "abstract": "\\noindent Obtaining simultaneous radio and X--ray data during the outburst decay of soft X--ray transients is a potentially important tool to study the disc -- jet connection. Here we report results of the analysis of (nearly) simultaneous radio (VLA or WSRT) and {\\it Chandra} X--ray observations of XTE~J1908+094 during the last part of the decay of the source after an outburst. The limit on the index of a radio -- X--ray correlation we find is consistent with the value of $\\sim0.7$ which was found for other black hole candidates in the low/hard state. Interestingly, the limit we find seems more consistent with a value of 1.4 which was recently shown to be typical for radiatively efficient accretion flow models. We further show that when the correlation--index is the same for two sources one can use the differences in normalisation in the radio -- X--ray flux correlation to estimate the distance towards the sources if the distance of one of them is accurately known (assuming black hole spin and mass and jet Lorentz factor differences are unimportant or minimal).  Finally, we observed a strong increase in the rate of decay of the X--ray flux. Between March 23, 2003 and April 19, 2003 the X--ray flux decayed with a factor $\\sim$5 whereas between April 19, 2003 and May 13, 2003, the X--ray flux decreased by a factor $\\sim$750. The source (0.5--10 keV) luminosity at the last {\\it Chandra} observation was ${\\rm L\\approx 3\\times10^{32} (\\frac{d}{8.5 kpc})^2 erg\\,s^{-1}}$. ",
        "introduction": "\\label{intro}  Low--mass X--ray binaries (LMXBs) are binary systems in which a $\\approxlt 1\\,M_{\\odot}$ star transfers matter to a neutron star or a black hole. These systems form one of our main windows on the physical processes taking place around black holes and hence they can provide us with information about the fundamental properties of spacetime. One reason for this is that the great majority of Galactic black hole candidates (BHCs) are found in transient LMXB systems.  Over the last few years it has become apparent that jets are an integral and energetically important part of these BHC systems (especially) when these systems are in the so called low/hard state (\\pcite{2001MNRAS.322...31F}; \\pcite{2001MNRAS.327.1273S}).  Recently, it was found that there exists a correlation between the radio and X--ray flux in the low/hard state of several BHCs over 3--4 decades in X--ray flux showing that there must be some form of disc--jet coupling (\\pcite{2003A&A...400.1007C}; \\pcite{2003MNRAS.344...60G}). \\scite{2003MNRAS.343L..59H}, \\scite{falcke2003}, and \\scite{merheinzdim2003} review the disc--jet connection in terms of different accretion disc and jet models.  \\scite{merheinzdim2003}, building on previous work of \\scite{2003MNRAS.343L..59H}, showed that inefficient accretion flow models can reproduce the observed radio -- X--ray correlation index for the initial parameter space they covered. \\scite{2003A&A...397..645M} showed that the jet--model explaining the observed X--rays in terms of synchrotron emission from the jet of \\scite{2001A&A...372L..25M} can reproduce both the observed correlation index as well as the normalisation. \\scite{2003MNRAS.343L..99F} used the observed radio -- X--ray correlation for BHCs to argue that there is no need to advect energy across a black hole event horizon in order to explain the observed difference in quiescent luminosity between the neutron star and BHC transient systems as was proposed by e.g.~\\scite{2001ApJ...553L..47G}. An important assumption in the work of \\scite{2003MNRAS.343L..99F} is that the observed radio -- X--ray correlation holds down to X--ray luminosities as low as L$_X \\sim10^{30-32}$ erg s$^{-1}$.  XTE~J1908+094 was discovered serendipitously during {\\it RXTE} observations of the soft gamma--ray repeater SGR~1900+14 by \\scite{2002IAUC.7856....1W}. The source flux is absorbed (${\\rm N_H\\sim2.3\\times10^{22} cm^{-2}}$), the spectrum is well--fit with a hard power--law with a photon index of 1.55.  Subsequent {\\it BeppoSAX} observations (\\pcite{2002IAUC.7873....1I}; \\pcite{2002A&A...394..553I}) confirmed both the hard spectrum (the source was detected up to 250 keV) and the high Galactic absorption. A broadened iron emission line was present in spectra extracted from both the RXTE and the BeppoSAX observations.  In 't Zand et al. (2002) presented strong evidence for a low/hard -- high/soft state change. The fact that the source displayed both a low/hard and a high/soft state during the outburst is confirmed by the timing and spectral analysis of the RXTE/PCA observations by \\scite{2002xrb..confE..11G} (see also \\pcite{gogus2004}). A radio counterpart was discovered by \\scite{2002IAUC.7874....1R} whereas a near--infrared counterpart was found by \\scite{2002MNRAS.337L..23C}. These authors also found that the optical upper limits (\\pcite{2002ATel...86....1W} and \\pcite{2002IAUC.7877....4G}) are fully consistent with the near--infrared colours of and the high extinction towards the source.  In this paper we report the findings of our (nearly) simultaneous Very Large Array (VLA)\\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.}  and Westerbork Synthesis Radio Telescopes (WSRT)\\footnote{The Westerbork Synthesis Radio Telescope is operated by the ASTRON (Netherlands Foundation for Research in Astronomy) with support from the Netherlands Foundation for Scientific Research NWO} radio and {\\it Chandra} X--ray observations of XTE~J1908+094 during the last part of the decay of the source after an outburst. ",
        "conclusions": "We have obtained (nearly) simultaneous VLA and WSRT radio and {\\it Chandra} X--ray observations of the BHC XTE~J1908+094 during the decay after an outburst. We find that: \\newline {\\it (i)} Limits on a power--law correlation between radio and X--ray flux suggest that the power--law index may be larger than that found before in other BHCs.\\newline {\\it (ii)} The rate of decay increases from a factor of $\\sim$5 in $\\sim$25 days to a factor of $\\sim$750 in $\\sim$25 days. \\newline {\\it (iii)} The source spectrum hardens during the decay. \\newline Below we will discuss these findings in more detail.  Previously, \\scite{2003A&A...400.1007C} and \\scite{2003MNRAS.344...60G} found that there is a correlation between the radio and X--ray flux over more than 4 orders of magnitude. The index of the power--law relation they fitted was consistent with being the same for several sources at a value of 0.7$\\pm$0.1 (Gallo et al.~2003). If we assume that the radio flux of the source on April 19, 2003 was at the 1~$\\sigma$ level of the WSRT upper limit the index was 1.5$^{+0.45}_{-0.3}$ for a two--point power--law decay for XTE~J1908+094. This is consistent at the $\\sim3\\sigma$ level with the value of 0.7 found before. However, there are some caveats. This index was determined over a limited range in X--ray flux and using very few measurements only. Furthermore, we assumed that the radio spectral index is flat with $\\alpha=0$ (S$_\\nu\\propto\\nu^{\\alpha}$). Finally, the radio flux on Qpril 19, 2003 may have been much lower than what we have assumed; this would make the correlation index steeper.  Previously, \\scite{2003MNRAS.342L..67M} found an index of $\\sim$1.4 for the radio -- X--ray correlation over a small range in X--ray flux in the neutron star system 4U~1728--34. Possibly, XTE~J1908+094 is a neutron star as well. However, \\scite{2002A&A...394..553I} and \\scite{2002xrb..confE..11G} argue strongly in favour of a BHC nature for XTE~J1908+094 on the basis of the observed outburst X--ray spectral and timing properties (see also Gogos et al.~2004, submitted), but obviously a dynamical mass estimate showing the mass of the compact object to be more than 3 M$_\\odot$ would settle the issue.  If we follow the line of reasoning laid--out in \\scite{2003MNRAS.343L..99F} but taking ${\\rm L_{Radio}\\propto L_X^{1.4}}$ instead of ${\\rm L_{Radio}\\propto L_X^{0.7}}$ we find that ${\\rm L_{Jet}\\propto L_X}$. Hence, the ratio between jet and accretion power (which is assumed to be tracked by the X--ray luminosity) remains the same as the source flux decays. We note that the fact that the index was found to be 1.4 for a neutron star system does not affect the conclusion of \\scite{2003MNRAS.343L..99F} that the difference in quiescent X--ray luminosity between BHCs and neutron star soft X--ray transients can be explained without the need of advection of energy across the event horizon as long as the index for the BHCs is 0.7.  Recently, \\scite{2003MNRAS.343L..59H} showed that the index of 0.7 follows naturally for several jet--models if one assumes that those jet models are scale invariant (i.e.~the Schwarzschild radius, r$_s$, is the only relevant length scale for jet formation).  Building on the work of \\scite{2003MNRAS.343L..59H}, \\scite{merheinzdim2003} also showed that other indexes for the radio -- X--ray correlations could be found, i.e.~for both a gas and a radiation pressure dominated disc the index would be $\\sim$1.4, whereas it would be close to 0.7 for radiatively inefficient accretion flows. \\scite{2003A&A...397..645M} show that the model explaining part of the X--ray emission as jet synchrotron emission (\\pcite{2001A&A...372L..25M}) can explain the 0.7 index as well as the normalisation of the radio -- X--ray correlation.  Hence, it seems that in XTE~J1908+094, during the part of the decay that we covered with our radio and X--ray observations, a standard geometrically thin optically thick disc plus a corona could have been present. Why accretion would proceed via a geometrically thin disc in XTE~J1908+094 whereas in other sources it is thought that the standard disc is not present in the low/hard state is unclear. Perhaps it has something to do with the luminosity levels at which the various sources are observed so far, but since the distance to XTE~J1908+094 is ill--constrained (\\pcite{2002A&A...394..553I} argue that the distance must be larger than 3 kpc) the source luminosity is not well known.  From Fig.~\\ref{xdecay} it is clear that the rate of decay increased enormously after April 12, 2003. Such a steep decrease has been observed before for several BHCs and neutron star soft X--ray transients (e.g.~\\pcite{1997ApJ...491..312C}). In neutron star systems this has been interpreted as evidence for the onset of the propeller effect (\\pcite{1998ApJ...499L..65C}), however, since such a drop in luminosity seems to be common for BHCs as well, this interpretation may need to be revised (\\pcite{2003MNRAS.341..823J}).  Finally, we find that the X--ray spectrum hardens between the first and second X--ray observation. Spectral hardening is often observed during/just after a transition to the low/hard state (cf.~\\pcite{2001ApJ...563..229T}). However, the findings of Gogos et al.~(2004) show that XTE~J1908+094 was already in the low/hard state several months before the first {\\it Chandra} observation was made. Perhaps the source changed from the low/hard state back to a soft state in between the {\\it RXTE} and the {\\it Chandra} observations. Such a change would be consistent with the fact that the limit on the radio -- X--ray correlation index is close to 1.4. We conclude that more observations at these low flux levels (and likely low luminosities) are necessary to determine the behaviour of these sources when they return back to quiescence."
    },
    "0403/astro-ph0403052_arXiv.txt": {
        "abstract": "The High Energy Stereoscopic System (\\hess) - is a system of four, 107~m$^{2}$ mirror area, imaging Cherenkov telescopes under construction in the Khomas Highland of Namibia (1800~m asl). The \\hess~ system is characterised by a low threshold ($\\sim$ 100~GeV) and a $\\sim$1\\% Crab flux sensitivity resulting from the good angular resolution and background rejection provided by the stereoscopic technique.  The first two telescopes are operational and first results are reported here. The remaining two telescopes (of \\hess~Phase-I) will be commissioned early in 2004. ",
        "introduction": "\\begin{figure} \\begin{center} \\includegraphics[height=19pc]{sensitivity.eps} \\end{center} \\caption{Expected point-source sensitivity of \\hess~Phase I compared to that of the VERITAS array\\cite{VERITAS} (reproduced from \\cite{HESS_ICRC01}).} \\label{fig_mc} \\end{figure}  \\hess~is an array of 4 identical telescopes arranged in a square with 120~m sides. Each telescope has a focal length of 15 m and a 13~m diameter. The reflectors comprise 380 quartz-coated round facets (60~cm diameter), arranged with Davies-Cotton optics.  The \\hess~cameras consist of 960 pixels of $0.16^\\circ$ angular size providing a total field of view of $5^\\circ$. A pixel consists of a photomultiplier tube (PMT) with a Winston cone light collector. The PMTs (Photonis XP2960) are organised into {\\em drawers} of 16 PMTs with associated read-out electronics. All triggering and read-out electronics are contained inside the camera body.  Monte-Carlo simulations predict a sensitivity for the 4 telescope system of around 1\\% of the Crab flux (5$\\sigma$, 50 hours) and an energy threshold of $\\approx$100 GeV (see Fig.~\\ref{fig_mc}). ",
        "conclusions": "Two of the four telescopes of \\hess~Phase-I are complete and operational. The optical and mechanical performance of these telescopes meets all specifications. $\\gamma$-ray observations of the Crab confirm the simulated performance of a single \\hess~telescope. We confirm the BL-Lac PKS 2155-304 as a strong VHE $\\gamma$-ray source.  The full \\hess~system should begin operations early in 2004."
    },
    "0403/astro-ph0403347_arXiv.txt": {
        "abstract": "s{We discuss some of the most unusual active galactic nuclei (AGN) discovered to date by the Sloan Digital Sky Survey (SDSS): the first broad absorption line quasar seen to exhibit He\\,{\\sc ii} absorption, several quasars with extremely strong, narrow UV Fe\\,{\\sc ii} emission, and an AGN with an unexplained and very strange continuum shape. } ",
        "introduction": "The Sloan Digital Sky Survey (York et al. 2000; Fukugita et al. 1996; Gunn et al. 1998; Hogg et al. 2001; Stoughton et al. 2002; Smith et al. 2002; Pier et al. 2003) is obtaining optical spectra for $\\sim$10$^5$ quasars over $\\sim$$\\frac{1}{4}$ of the entire sky. Through careful target selection (Richards et al. 2002) and sheer size, the SDSS includes numerous AGN with unconventional properties. The high-quality, moderate-resolution SDSS spectra can be used to set the stage for the detailed multiwavelength studies often needed to understand interesting quasar subclasses. We illustrate this fact using several unusual AGN included in the SDSS Second Data Release (Abazajian et al. 2004). ",
        "conclusions": ""
    },
    "0403/astro-ph0403171_arXiv.txt": {
        "abstract": "s{ We show that within a recently developed nonlocal, chiral quark model the critical densities for a phase transition to color superconducting quark matter under neutron star conditions can be low enough that these phases occur in compact star configurations with masses below $1.4~M_\\odot$. We study the cooling of these objects in isolation for different values of the gravitational mass and thus different composition and structure of the interior. Our equation of state allows for a  2SC phase with a large quark gap $\\Delta \\sim 100~$MeV for $u$ and $d$ quarks of two colors, a normal quark matter phase and their coexistence in a mixed phase within the hybrid star interior.  We argue that, if  the phases with unpaired quarks were allowed, the corresponding hybrid stars would cool too fast to describe the neutron star cooling data existing by today. We incorporate other attractive channels permitting a weak pairing of the residual quarks which remained unpaired in the 2SC phase and demonstrate that the model does not contradict the cooling data if the weak pairing gaps are of the order of $0.1~$ MeV. } ",
        "introduction": "In the recent paper \\cite{bgv2004}, hereafter BGV, we have reinvestigated the cooling of neutron stars (NS) within a purely hadron model, i.e., ignoring the possibility of quark cores in NS interiors. We have demonstrated that the NS cooling data available by today can be well explained within the {\\em \"Nuclear medium cooling scenario\"}, i.e., if one includes medium effects in the emissivity and takes into account a suppression of  the $3P_2$ neutron gap. In a subsequent work \\cite{bgv2004q} we have shown that this result does not exclude the possibility that neutron stars might possess large quark matter cores that extend up to more than half of the star radius. Such a hybrid structure gives room for a whole variety of additional scenarios of compact star cooling which fall into two classes: either nuclear and quark matter phases have similar cooling behavior (homgeneous cooling) or the faster cooling of the one phase is compensated by the slower cooling of the other (inhomogeneous cooling). In the present contribution we will report on our results within the former, homgeneous cooling scenario of hybrid stars and what implications the comparison with present-day cooling data may provide for the EoS and transport properties of quark matter. ",
        "conclusions": ""
    },
    "0403/astro-ph0403492_arXiv.txt": {
        "abstract": "\\noindent If the causality condition [the speed of sound always  remains less  than  that  of  light  in vacuum, i. e., $v \\leq c = 1$] is imposed on the spheres of homogeneous energy density,  the `ratio of  the  specific  heats',  $\\gamma \\leq  2.59457$,  constraints  the compaction parameter, $u [\\equiv (M/a)$, mass to size ratio in geometrized units]  of  the  dynamically  stable  configurations $ \\leq  0.34056 $ [corresponding to a surface redshift $(z_a ) \\leq 0.771$]. Apparently, The maximum value of $u$ obtained in this  manner belongs to an absolute upper bound, and gives: (i)  The  maximum value for static neutron star masses as $5.4 M_\\odot$, if  we  substitute the density at the surface  of  the  configuration  equal  to  the average nuclear density, $E = 2 \\times 10^{14}$ g\\, cm$^{-3}$ [e.g. {\\it Nature}, {\\bf 259}, 377 (1976)]. (ii) However,  if  the  density  of  the  static configuration is constrained to the value $1.072 \\times 10^{14}$ g\\, cm$^{-3}$, by imposing the empirical result that the minimum rotation period  of the fastest rotating pulsar known to date, PSR 1937 + 21, is 1.558 ms, the maximum mass value for static neutron  stars  exceed  upto $7.4 M_\\odot$. These masses have important implications for  the  massive compact objects like Cyg X-1, Cyg XR-1, and LMC-X3 etc., which may not,  necessarily,  represent  black  holes.  (iii)  The   minimum rotation periods for a static $1.442 M_\\odot$  neutron star to  be  0.3041 ms. (iv) A suitable stable model of ultra-compact objects [$u \\ > (1/3)$] which has important astrophysical significance. ",
        "introduction": "%  Incompressible fluid spheres of uniform energy density $E$ in General Relativity were first discussed by  Schwarzschild.$^1$ The importance of this solution in  General  Relativistic  stellar structures is apparent, because it gives an absolute  upper  limit on compaction parameter, $u (\\equiv M/a$, mass to size  ratio  of  entire configuration in geometrized units) $ \\leq (4/9) $ for any regular static solution in hydrostatic equilibrium.$^2$  Chandrasekhar$^{3,4}$   discussed   the   condition   of hydrostatic equilibrium under the small  adiabatic  perturbations, and showed that  for  each  value  of  the  compaction  parameter, corresponding  to  the  {\\em compressible}  homogeneous  spheres,  there exists a critical  (minimum)  value  of  the  ``ratio  of  specific heats'', $\\gamma (= \\gamma_{crit}$) such that for $\\gamma \\ <  \\gamma_{crit}$, the  configuration becomes  dynamically  unstable.  For  the  limiting  case  of  the compaction parameter approaching  the  Schwarzschild  limit  ($u  = 4/9), \\gamma$ becomes infinity.  For dynamically stable superdense objects like neutron  stars one may expect a finite value of $\\gamma$. However, for such objects  the equations of state are not well  known  [empirically]  beyond  the density $ \\cong 10^{14} $ g\\, cm$^{-3},^5$ and one  can  only  extrapolate the equations of state (available  in  the  literature)$^6$ far beyond this density. As  a  way  out,   one  can impose some restrictions upon the known physical quantities,  such that, the speed of sound inside the configuration, \\begin{displaymath} v \\equiv \\sqrt {(\\partial {P} / \\partial {E})_s} \\end{displaymath} (Where $P$ is the pressure, $E$ is the energy density and $s$ stands for specific entropy) does not exceed the speed of light in vacuum, i.e., $v \\leq c = 1$  (in geometrized units), and obtain an upper bound  on  stable  neutron star masses.$^{7-9}$  In the present paper, we have  obtained  an  upper  bound  on compaction parameter ($u  \\leq  0.34056$  corresponding  to  a  surface redshift  of  0.771)  for  the  compressible  homogeneous  spheres,$^{3,4}$ by imposing constraint on the ``ratio  of specific heats'', $\\gamma [ \\leq 2.59457)]$, compatible with causality ($v \\leq 1$) and dynamical stability. This value of the compaction parameter is an absolute maximum because, for an  assigned  value  of  $\\gamma$,  the maximum compactness would correspond  to  a  compressible  uniform density sphere, and can be used  to  obtain  an  upper  bound  on neutron star masses, as well as the minimum rotation period  of  a $1.442 M_\\odot$  neutron star  (the  maximum  mass  of  the  neutron  star accurately known at present).$^{10}$ ",
        "conclusions": "An absolute upper bound on compaction parameter, $u  \\leq  0.34056 $ [or the surface redshift $ \\leq 0.771$], compatible with  causality  and the ratio of the specific heats, $\\gamma \\leq 2.59457$, is obtained by  using the dynamically stable compressible homogeneous sphere. This upper limit of compaction parameter gives  (i) The maximum static mass of conventional model of neutron stars [taking $E = 2 \\times 10^{14}$ g\\, cm$^{-3}]^{7,8}$ as $5.4 M_\\odot$. This is greater than $4.8 M_\\odot$   considered  as  an upper limit earlier.  (ii) The maximum mass of static neutron star exceeds to the  value of $7.387 M_\\odot$  which is greater than the upper limit of $5.3 M_\\odot$ for neutron stars (so called $Q$-star models) obtained by Hochron,  Lynn and Selipesky.$^{26}$ This may have important implications for  the heavy compact objects like Cyg X-1, Cyg XR-1, and LMC-X3 which may not, necessarily, be black holes.  (iii) The minimum rotation periods for a static $1.442 M_\\odot$   neutron star to be 0.3041 ms with a uniform energy density $E$ as $2.813  \\times 10^{15}$ g\\, cm$^{-3}$.  (iv) A causally consistent and dynamically stable model  of  ultra-compact objects [$u \\ >  (1/3)$]  which  has  important  astrophysical significance [see, e.g. Ref. 16 and  references therein]."
    },
    "0403/astro-ph0403201_arXiv.txt": {
        "abstract": "Protostellar sources in star forming regions are responsible for driving jets with flow velocities ranging between 300 and 400 km s$^{-1}$. This class of jets consists of highly collimated outflows which include  thermal knots with number densities estimated to be greater than that of their ambient medium. On the other hand, extragalactic FR I jets consist of light fluid with low  Mach number burrowing  through a denser medium as the magnetized  jets   radiate  nonthermal emission. Both protostellar and extragalactic jets are believed to be launched by accretion disks. Here we consider a jet model in which the characteristics common to both protostellar and extragalactic jets are used to explain the origin of  nonthermal filaments in the Galactic center region. We argue that these filaments are analogous to FR I extragalactic sources but are launched by embedded young stars or clusters of stars in star-forming regions. ",
        "introduction": "It has been 20 years since the discovery of the nonthermal radio filaments (NRFs) associated with the Galactic center Arc was first reported (Yusef-Zadeh, Morris, \\& Chance 1984).  These observations showed evidence of linear, magnetized features running perpendicular to the Galactic plane. A number of NRFs with similar characteristics to the prototype NRFs have been discovered in the intervening years (Yusef-Zadeh 2003 and references therein). Several models have suggested that the filaments trace the illuminated component of a large-scale poloidal magnetic field pervasive throughout the Galactic center region. However, the presence of a number of filaments oriented at large angles to the normal to the Galactic plane does not support the above interpretation and indicates a different origin. Unlike most models that predict a global, static geometry of the magnetic field around the Galactic center, the model described here argues for a local origin. In the proposed picture, the NRFs originate in star-forming regions. The filaments behave like jets extracting mass and energy from embedded young stars or clusters of stars as the jets propagate in a dense ISM of the Galactic center region. A more detailed account of this model will be given elsewhere. ",
        "conclusions": ""
    },
    "0403/astro-ph0403037_arXiv.txt": {
        "abstract": "The principal results of daily observations with the RATAN-600 radio telescope of X-ray binary with relativistic jets microquasar SS433  in 1986--2003 are presented. We have measured the flux densities  at 0.96, 2.3, 3.9, 7.7, 11.2 and 21.7 GHz in different sets, duration from a week to some months. In general there are 940 observations of SS433 and more than 4500 flux density measurements in the period. Observations show that radio spectra are well fitting by a power law. The mean spectral index remained the same, $-0.60\\pm0.14$ during almost 20 years at least, and mean accuracy of the index determination was better than 0.1 in our multi-frequency observations,  i.e. it was higher than in the intensive two-frequency monitoring of SS433  with the three-element GBI interferometer. Flux density data and spectra `on-line' plotting   are accessible on the CATS data base site: http://cats.sao.ru/. \\keywords {X-rays: binaries -- stars: flare -- stars: individual: SS433 -- jets -- radio continuum: stars -- monitoring } ",
        "introduction": "The X-ray binaries (XBs) have long been studied in radio band. After an identification of Cyg X-3 in the beginning of the 70-s it became evident that XBs could have powerful variable radio emission. A sample of X-ray binaries with relativistic jets that Mirabel {\\&} Rodriguez (1999) named microquasars consists of 15--20 objects. The brightest of them were observed actively under a monitoring program of XBs with the RATAN-600 radio telescope (Trushkin 2000, Trushkin \\& Bursov 2001).  SS433 --- a bright variable emission star --- was identified by Clark and Murdin (1978) with a rather bright compact radio source 1909+048 located in the center of a supernova remnant W50. When in 1979 mobile optical emission lines were discovered in the spectrum of this bright star SS433 --- a radio source 1909+04 (Margon et al. 1979), it became apparent that a new class of objects in the Galaxy was found. At the same time Spencer (1979) was the first to discover an extended structure: a compact core and 1 arcsec long aligned jets in the radio image of SS433. At present such a structure in microquasars is commonly named a radio jet. Different data do indicate a presence of a very narrow (about 1$\\degr$) collimated beam at least in X-ray and optical ranges. At present there is no doubt that SS433 is related to W50. SS433 is probably a stellar remnant of a SN that exploded in a binary system about 30000~years ago. An explosion of one of the components did not destroy the binary. A distance to SS433 of 4.8 kpc was later determined by different ways including the direct measurement of proper motions of the jet radio components.  As is now known, this unique variable X-ray, IR and radio source is an eclipsing binary cosisting of a compact object and an early type massive star. The system has a pair of opposite relativistic (in which matter is ejected at 0.26c) jets, which show themselves in different ways in different ranges from X-ray to radio. The VLBI (Very Long Baseline Radio Interferometry) observations of the jets proper motion showed that they consist of separate blobs that move ballistically at the same velocity  0.26c. Resultant from the 164-day precession of a thick accreting disk the jets rotate along conic surface and look like structures similar to a twin corkscrew on VLA radio maps (Hjellming {\\&} Johnston 1981).  There are strong evidences for believing that this binary system consists of a relativistic star with an accreting disk and a massive optical component filling its Roche lobe. The supercritical accretion onto the relativistic star is likely to lead to initiation of the SS433 main feature --- two opposite jets of matter moving from the accretion disk poles at a velocity of about a quarter of the speed of light. The recent intensive spectral investigation made it possible to determine rather precisely the SS433 mass function and to estimate the relativistic component mass. It is equal to 11~$\\pm $~5~M$_{\\sun}$ (Gies et al. 2002a,b) what suggests that this is a black hole. Absorption lines in the SS433 spectrum belonging to the optical component are almost identical with those in the spectrum of an evolved star of a mid-A spectral class with a mass of 19~$\\pm $~7~M$_{\\sun}$.  These jets that had been first discovered in spectral optical observations are seen also in X-ray and radio wavelengths. Activity phenomena in SS433 give rise to strong variability of all its electromagnetic spectrum.  As is seen from the jet X-ray images and from an oblong form of the radio remnant W50, the jet structure extends up to size scales of $1\\degr$. It gives an estimate of jets age as $\\sim$ 1000 years. A  simple kinematic model of SS433 (Abell {\\&} Margon 1979) with collimated jets was verified as a whole by many observational data (Margon 1984).  SS433 is still the unique object in the Galaxy, in which the relativistic opposing jets appear in Doppler-shifted optical and X-ray emissions, i.e. beside the relativistic radio-emitting electron-positron plasma a baryonic matter with high atomic numbers moves in the SS433 jets (Kotani et al. 1996, Marshall et al. 2002).  The major contribution to the study of temporal and spectral properties of microquasars was made by long-term  program of variable sources monitoring at two frequencies of 2.25 and 8.3 GHz with GBI (Green-Bank Interferometer, NRAO). The publicly available data totally include about 16000 measurements of flux density of SS433. Fig.\\ref{gbi} shows the  light curve of SS433 in 1979--2000 which cover all its GBI-observations at a frequency of 2.25~GHz. The results of these studies were partly published by Johnston et al. (1984) and Fiedler et al. (1987).  From recent interesting results of the SS433 study noteworthy is the circular polarization of its emission in cm range detected by Fender et al. (2000) and a sign change of the Stokes parameter V detected by McCormick et al. (2003) in the 1--9~GHz range. It might be related either to the presence of a gyro-synchrotron radiation emitted by relatively low-energy electrons or to the change in a configuration of collimated magnetic field in the jets.  Space observatory INTEGRAL first registered a variable hard X-ray emission (25--100~keV) (Cherepashchuk et al. 2003). It was shown that the emission level is influenced by the precession and orbital motions. These observations in high energy range confirm indirectly once more a presence of a black hole in this XB.  In radio band a quasisteady synchrotron emission of SS433 with spectrum S$_{\\nu }$~$\\sim $~S$_{0}\\nu ^{\\alpha }$ is superimposed by a non-thermal flares, when the flux density can exceed the quiescent level 10 times. The SS433 radio flares are not rare, although the periods of relative quiescence can be as long as 100--200~days.  At the flares onset there are evidences of synchrotron source opacity (a flat spectrum) and with further gradual decay of the flare the spectrum becomes steeper. The maximum flux of a flare comes later at lower frequencies as was shown for SS433 by Fiedler et al. (1987), Vermeulen et al. (1993a,b), Bursov {\\&} Trushkin (1995); for Cyg X-3 see Waltman et al. (1994, 1995, 1996) and for GRS~1915+105 --- Trushkin et al. (2001) and Fender et al. (2002). It corresponds to evolution of an adiabatically expanding cloud of relativistic emitting electrons. Shklovskij (1960) and van der Laan (1966) were the first to formulate basic equations describing this evolution.  \\begin{figure*} \\centerline{\\vbox{ \\psfig{figure=1909gbi.eps,width=16cm,angle=-90} }} \\caption{% The SS433 light curves at a frequency of 2.25 GHz by data of GBI (NRAO/NASA).} \\label{gbi} \\end{figure*}  Different modifications of this model necessary for observations fitting include a hollow conic geometry of jets, an account for synchrotron losses or reverse Compton scattering, a dynamic motion outwards a thermally absorbing envelope. Computations of flares  based on  a model of twin conic jet show a remarkable coincidence with observations for Cir\\,X-1 (Garcia 1995),   SS433 (Hjellming {\\&} Johnston 1988), LSI+61$\\degr$\\,303 (Paredes et al. 1991), {\\mbox Cyg\\,X-3} (Marscher at al. 1975, Marti et al. 1992) and SAXJ1819-254 (Hjellming et al. 2000).  Based on data of ten SS433 flares in 1986--87, Trushkin (1989) showed that on average they have close frequency power relations for the maximum flare flux and the time of this maximum occurrence: $$\\Delta S_{m}{\\rm(Jy)}=1.3 \\nu^{-0.4\\pm 0.15},$$ $$\\Delta t_{m}{\\rm (days)}=5\\nu^{-0.4\\pm 0.1},$$ where $\\nu$ is in GHz.  A gradual flare decay follows one of two time laws. An original model of Shklovskij -- van der Laan provides the power law. In many microquasars there was actually detected the radio emission decay according to the law $$S_\\nu = S_o\\,t^{-2p}\\nu^{\\alpha},$$ where $p\\approx 2$ is an index of a power electron distribution that is related to the spectral index $$\\alpha = (1-p)/2\\approx-0.5$$ for an optically thin source. It is in excellent agreement with a possible diffusive mechanism of relativistic electrons acceleration in shock waves. At a compression index in a strong shock wave $\\rho\\le4$ the spectral index of these electrons is $$p = (\\rho+2)/(\\rho-1) = 2.$$  \\begin{table*} \\caption { A flux density sensitivity of the ``Northern Sector'' of RATAN-600} \\begin{center} \\medskip \\begin{tabular}{l|l|l|l|l|l|l|l} \\noalign{\\smallskip} \\hline Wavelength (cm) &31.2  &13.0 & ~7.6  &~6.2  & ~3.9 &  ~2.7& 1.38 \\\\ \\hline Frequency (GHz)    & 0.96 &~2.3 & ~3.9  &~4.9  & ~7.7 &  11.2& 21.7  \\\\ \\hline $\\Delta$S(mJy)   & 40   & 15  &   ~3  &  ~3  &   8  &   10 &  20   \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}  Other flares and fluxes of plasmons in jets were decaying according to the exponential law $$S_{\\nu }=S_{0}e^{-t/ \\tau }$$ (for SS433 see Jowett \\& Spencer 1995). A diminution of the characteristic time $\\tau $ with frequency was often but not always observed. A relation $$\\tau {\\rm(days)}=11.5~\\nu ^{-0.29\\pm 0.03}$$ where $\\nu$ is in GHz was observed  in Cyg X-3 (Trushkin 1998). In Cyg X-3 a transit from the exponential law to the power one was detected several times (Hjellming et al. 1974, Marscher et al. 1975, Bursov {\\&} Trushkin 1995, Trushkin 1998) what is interpreted as a transit from a mode with a dominated radiative losses to a mode of adiabatic expansion.  Spencer (1996) calculated the energy of the jet components based on the condition for minimum energy of relativistic particles and of magnetic field what is fulfilled at an equipartition of energy of these components. He showed that a considerable not to say prevailing part of energy released during flares belonged to radio emitting particles and magnetic field of jets.  Below we discuss a few large monitoring programs carried out at the Northern Sector of RATAN-600 having resulted in the SS433 light curves at 4--6 frequencies. ",
        "conclusions": "This paper compile all observational data over many sets from December 1986 to January 2004 containing 940 observations of SS433 at two--six frequencies simultaneously. In all more than 4500 measurements of flux density at frequencies from 960 to 21700~MHz are presented. The average spectral index by all data is equal to --0.60~$\\pm $~0.14 with an average measurement error of 0.09 at an average flux density of 1.5~Jy at 960~MHz. In the indicated frequency range there were no cases of the spectrum inversion, when the spectral index would be positive. Thus, if there exists such a stage of the SS433 spectrum evolution, it is essentially shorter than one day not to get in such a moment during almost 1000 observations.  It was demonstrated that there are some earlier unknown regularities of the SS433 activity in the measured light curves. By the data on many flares the delay of flare maximum towards lower observational frequencies was detected. On the other side, the value of this maximum falls with frequency according to a power law. The flares fading with time also follows a power law. The indices of these relations change from flare to flare, but median values of indices coincide astonishingly well. They are equal to --0.4~$\\pm $~0.1. It is evidently an indication that the integral emission of the SS433 jets is due to similar properties of synchrotron radiation of the blobs--plasmons moving separately. In the end it is rigidly related with processes inside this XR and jets on the whole. From this point of view the models with inner shock waves propagating  and increasing radio emission along the jet could be more preferable than the models in which the evolutions of separate radio components, blobs, are independent from each other.  The accumulated data of the flux density measurements are presented as publicly-accessible Web-programs for constructing the SS433 spectra in the Home Page of the data base CATS: http://cats.sao.ru/cgi-bin/ss433.cgi.  The complete data of all measurements carried out with RATAN-600 are presented in Table~3."
    },
    "0403/astro-ph0403084_arXiv.txt": {
        "abstract": "We present the results of a wide-area mapping of the far-infrared continuum emission toward the Orion complex by using a Japanese balloon-borne telescope. The 155-$\\mu$m continuum emission was detected over a region of 1.5 deg$^2$ around the KL nebula with $\\timeform{3'}$ resolution similar to that of the IRAS 100-$\\mu$m map. Assuming a single-temperature model of the thermal equilibrium dust, maps of the temperature and the optical depth were derived from the 155 $\\mu$m intensity and the IRAS 100 $\\mu$m intensity. The derived dust temperature is 5 -- 15 K lower and the derived dust optical thickness were derived from the 155-$\\mu$m intensity and the IRAS 100-$\\mu$m intensity. The derived dust temperature is 5 -- 15 K lower and the derived dust optical depth is 5 -- 300 times larger than those derived from the IRAS 60 and 100-$\\mu$m intensities due to the significant contribution of the statistically heated very small grains to the IRAS 60-$\\mu$m intensity. The optical-thickness distribution shows a filamentary dust ridge that has a $\\timeform{1D.5}$ extent in the north -- south direction and well resembles the Integral-Shaped Filament (ISF) molecular gas distribution. The gas-to-dust ratio derived from the CO molecular gas distribution along the ISF is in the range 30 -- 200, which may be interpreted as being an effect of CO depletion due to the photodissociation and/or the freezing on dust grains. ",
        "introduction": "The Orion region at a distance of 470 pc \\citep{Brown1994}, is the nearest active star-forming region containing massive stars. It contains a massive molecular cloud complex, bright H\\emissiontype{II} regions, and many young stellar objects. Many studies of this region as a representative case of a massive-star-forming region have been made and have provided insight into the massive-star-formation process (see reviews by \\cite{Goudis1982, Genzel1989, O'Dell2001}). Even now, it is the most important region for investigating the massive-star-formation phenomena. M 42 is the brightest H\\emissiontype{II} region in this region, and is excited by the Trapezium star cluster. A dense molecular cloud complex lies behind the H\\emissiontype{II} region. In the densest part of this molecular could (OMC-1) are the Orion KL and BN objects \\citep{KL, BN}. They are embedded deeply in OMC-1, and energize the surrounding molecular gas of OMC-1. Two other dense molecular clouds, OMC-2 and -3, lie along the dense molecular ridge north of Orion KL. This dense ridge has a filamentary structure, $\\timeform{1.5D}$ long. From the apparent shape seen in the ${}^{13}$CO high-resolution map by \\citet{Bally1987}, it was named the Integral-Shaped Filament (hereafter ISF). The cloud cores of OMC-1, -2, and -3 lie on the dense filamentary feature. At the northernmost position of ISF, there is a H\\emissiontype{II} region, NGC 1977, and the south of it, there are small H\\emissiontype{II} regions, L 1641-N and NGC 1999. A H\\emissiontype{II} region, M42 containing the Orion KL, lies on the middle of the ISF. The ISF has so far been observed and identified in various molecular lines of CO (J = 2--1) by \\citet{Sakamoto1994}, ${}^{13}$CO (J = 1--0) by \\citet{Nagahama1998}, C${}^{18}$O (J = 1--0) by \\citet{Dutrey1991}, NH${}_3$ by \\citet{Cesanori1994}, and CS by \\citet{Tatematsu1993}, and in the neutral atomic carbon [C\\emissiontype{I}] line by \\citet{Ikeda2002}.  The continuum emission in the far-infrared and sub-millimeter wavelength regions is radiated by interstellar dust grains that are mixed with the molecular gas. Thanks to recent progress in instrumental techniques in this wavelength region, relatively wide areas of the sky have been mapped with ground-based telescopes \\citep{Chini1997, Lis1998, Johnstone1999} and with balloon-borne telescopes \\citep{Ristorcelli1998, Mookerjea2000-a, Dupac2001}. Previous investigations successfully revealed the dust distribution of this region, but covered only a part of the ISF-OMC-1 and its vicinity, OMC-2, and OMC-3. \\citet{Mookerjea2000-a} mapped this region at 138 and 205 $\\mu$m with a spatial resolution of $\\timeform{1.5'}$ using a one-meter balloon-borne telescope, and showed that a dust ridge extends from Orion KL to the north, tracing the molecular ridge. \\citet{Ristorcelli1998, Dupac2001} obtained wider-area maps at 200, 260, 360, and 580 $\\mu$m with the large balloon-borne telescope, PRONAOUS. The observed area was from Orion KL to the north. Their observations revealed several cold condensations to the west of OMC-1, which may be pre-collapse phase cloud cores. \\citet{Johnstone1999} mapped the southern ridge with SCUBA of JCMT, and showed that the dust ridge really extends along the ISF at least up to $\\timeform{30'}$ from Orion KL. However, the observed area ($\\timeform{50'} \\ \\times \\ \\timeform{10'}$) did not cover the whole of the ISF. It is of great interest to map the interstellar dust distribution over the whole region of the ISF. As is well known, IRAS completely mapped this region at four infrared bands of 12, 25, 60, and 100 $\\mu$mm. By using the IRAS 60 and 100 $\\mu$m maps, \\citet{Bally1991} derived the maps of dust temperature and the dust mass. However, they could not identify the ISF-like structure in the derived maps. The main reason for this is the unsuitability of the IRAS 60 $\\mu$m band for the study of a cold dust region like the ISF, because in this band, the emission from stochastically-heated, very small grains significantly contaminates the thermal emission from the large grains at a steady state temperature (e.g. \\cite{Nagata2002}). Moreover, as is mentioned later, the IRAS 60 and 100 $\\mu$m band signals are saturated at and near the Orion KL position and have a spurious pattern due to the insufficient baffle design. COBE/DIRBE had achieved highly reliable, well-calibrated intensity maps at various infrared bands for the whole sky. The 100, 140, and 240 $\\mu$m intensity maps of DIRBE can be used for deriving the dust temperature and mass distribution. However, the spatial resolution of $\\timeform{0.7D}$ of the DIRBE maps is insufficient for resolving the dust distribution in the Orion region \\citep{Wall1996}. The aim of the present work is to observe a very large area covering the whole ISF at wavelengths between 100 and 200 $\\mu$m, where the interstellar dust (large grains) has its peak emission, with a spatial resolution similar to that of IRAS 100 $\\mu$m band in order to reveal the dust temperature and mass distribution and, consequently, the energy source of individual parts of molecular clouds in this region.  The instrumentation and observation are described in section 2; the data reduction, analysis, and calibration are explained in section 3. The intensity map and the temperature distribution are presented in section 4. In section 5, the column-density distribution of the dust grain is derived, and the heating sources for this region are discussed. ",
        "conclusions": "This paper presents the result of a wide-area mapping of the 155 $\\mu$m continuum emission in the Orion molecular cloud complex. The Japanese balloon-borne telescope (FIRBE) surveyed a sky area over 50 deg$^2$ and detected far-infrared emission over the region of 1.5 deg$^2$ around the KL nebula with $\\timeform{3'}$ resolution. The peak emission at Orion KL has a flux value of 42.5 $\\pm$ 4.8 kJy in the $\\phi \\ \\timeform{3'}$ beam. The far-infrared distribution of intense region is similar that of the IRAS 100 $\\mu$m image. The dust temperature and optical depth of the thermal-equilibrium large grain are estimated by applying a single-temperature dust model modified with a $\\lambda^{-2}$ emissivity law to the IRAS 100 and FIRBE 155 $\\mu$m intensities. They range from 15 to 30 K and from 0.15 to 0.0015, respectively. Due to the significant contribution of the statistically heated very small grain in the IRAS 60 $\\mu$m intensity, the dust temperature T(60/100) derived from the IRAS 60 and IRAS 100 $\\mu$m intensities is 5 -- 15 K higher than the thermal-equilibrium dust temperature. As a result, the dust optical depth at 100 $\\mu$m, derived from T(60/100) has a 5 -- 300 times smaller value. Toward the H\\emissiontype{II} region, M 42, the IRAS 60 $\\mu$m and also IRAS 25 $\\mu$m intensities are relatively stronger than those outside of the H\\emissiontype{II} region. This tendency in the H\\emissiontype{II} region may be attributed to the result of the statistically heating of very small grains by double photon incidence \\citep{Okumura1999}. The region where the dust temperature is greater than 20 K around M 42 has an extent similar to that of the optical region. This fact suggests that the far-infrared emission around M 42 is mainly provided by the Trapezium Cluster. The optical-depth distribution shows a filamentary dust ridge that has a $\\timeform{1.5D}$ extent in the north -- south direction. The shape of this dense dust ridge resembles the Integral-Shaped Filament (ISF) molecular gas distribution. The gas-to-dust ratio derived from the CO molecular gas distribution along the ISF shows 100 -- 200 around the M 42 ionized region on the ISF. However, the gas-to-dust ratio decreases in other regions of the ISF, which may be interpreted as being an effect of CO depletion due to the photodissociation and/or freezing on dust grains.   \\bigskip We would like to acknowledge Dr. T. Nagahama and Dr. Y. Fukui for allowing us to use the ${}^{13}$CO data and Dr. X. Dupac for providing us with the PRONAOUS data. We also express our special thanks to the technical staff of Tata Institute of Fundamental Research and National Balloon Facility of India. This work was supported by the Grant-in Aid of the Japan Society for the Promotion of Science (10147102) and by the Research Fellowship of Japan Society for the Promotion of Science.    \\clearpage \\begin{figure} \\begin{center} \\FigureFile(80mm,80mm){Figure1.eps} \\end{center} \\caption{ The thick curve shows the relative systematic spectral response of the FIRBE 155 $\\mu$m band. The thin-curve and the broken-curve are the spectral responses of the IRAS 100 $\\mu$m band and the COBE/DIRBE 140 $\\mu$m band, respectively (IRAS and COBE/DIRBE Explanatory Supplement). All curves are normalized at their peaks. } \\end{figure}   \\clearpage \\begin{figure} \\begin{center} \\FigureFile(80mm,80mm){Figure2.eps} \\end{center} \\caption{ Example of raw detector signals of eight pixels in one column of the 4 by 8 pixel array at crossing Orion KL. The horizontal axis shows the azimuth offset relative to the peak position that corresponds to Orion KL. } \\end{figure}   \\clearpage \\begin{figure} \\begin{center} \\FigureFile(160mm,160mm){Figure3.eps} \\end{center} \\caption{ (a) 155 $\\mu$m continuum intensity map observed by the FIRBE telescope. The spatial resolution is $\\timeform{3'}$ and the intensity unit is MJy sr$^{-1}$. The lowest contour levels are 352 (=2.6 $\\sigma$), 488 (= 3.6$\\sigma$), 600, 700, 800, 900, 1000, 1300, 1000 interval up to 10000, 10000 interval up to peak intensity of 71080 MJy sr$^{-1}$ corresponding to the flux of 42.5 kJy in $\\phi \\ \\timeform{3'}$ beam. (b) The IRAS 100 $\\mu$m intensity map of the same area after correction for the DIRBE -- IRAS calibration scale difference (multiplied by 0.72). The contour levels are 70, 100, 150, 100 interval from 200 up to 1000, and 1000 interval up to 10000. The IRAS 100 $\\mu$m intensity around Orion KL is saturated and the sprious patterns appear at an angular distance around $\\timeform{1D}$ away from Orion KL. } \\end{figure}   \\clearpage \\begin{figure} \\begin{center} \\FigureFile(160mm,160mm){Figure4.eps} \\end{center} \\caption{ (a) Dust temperature ($T_{\\mathrm{LG}}$) map derived from the 155 $\\mu$m intensity of the present work and the IRAS 100 $\\mu$m intensity. The contours indicate 16, 20, 24, and 28 K. (b) Same as (a) from the IRAS 60 and 100 $\\mu$m intensities. The contours indicate at intervals of 4 K from 24 to 40 K. (c) Red-band image of the Digital Sky Survey overlaid on the contour map of (b). The white-colored region represents the H\\emissiontype{II} region determined from the 330 MHz radio continuum map by \\citet{Subrahmanyan2001}. } \\end{figure}   \\clearpage \\begin{figure} \\begin{center} \\FigureFile(100mm,100mm){Figure5.eps} \\end{center} \\caption{ Schematic diagram for defining the three regions around M 42 superposed on the red-band image. The gray central rectangular region indicates the region where the IRAS 100 $\\mu$m band is saturated. } \\end{figure}   \\clearpage \\begin{figure} \\begin{center} \\FigureFile(160mm,160mm){Figure6.eps} \\end{center} \\caption{ (a)Gas column density map derived from the dust optical depth and a gas-to-dust mass ratio of 100. Logarithmic contours are shown along the ISF at the interval of 0.2 from 22.5 up to 23.5. (b) ${}^{13}$CO (J = 1--0) total integrated intensity map with $\\timeform{3'}$ spatial resolution \\citep{Nagahama1998}, where the velocity range is from 0 to 15 K km s$^{-1}$, is overlaid on the color map described in figure 6a. Contour levels are at 3.0, 5.0, 7.0, and at the intervals of 5 K km s$^{-1}$ from 10.0 K km s$^{-1}$. } \\end{figure}   \\clearpage \\begin{figure} \\begin{center} \\FigureFile(160mm,160mm){Figure7.eps} \\end{center} \\caption{ (a) The gray area indicates the area for which the dust mass was derived as in (b). The contour is the ${}^{13}$CO integrated intensity map. The vertical axis is the declination offset relative to the OMC-1 position. (b) Mean dust mass distribution along the ISF ridge, averaged in each $\\timeform{3'}$ bin in declination. (c) Same as (b) for the gas mass. (d) Gas-to-dust distribution. } \\end{figure}    \\clearpage \\begin{table} \\caption{Intensity ratio of the `High' Temperature Region to the surrounding reference region.} \\begin{center} \\begin{tabular}{lccccc} \\hline \\hline & \\multicolumn{4}{c}{\\textit{IRAS}} & \\textit{FIRBE} \\\\ \\cline{2-6} Band & 12 $\\mu$m & 25 $\\mu$m & 60 $\\mu$m & 100 $\\mu$m & 155 $\\mu$m \\\\ \\hline Intensity ratio & 1.30 $\\pm$ 0.20 & 1.72 $\\pm$ 0.42 & 1.71 $\\pm$ 0.36 & 1.26 $\\pm$ 0.21 & 1.25 $\\pm$ 0.25 \\\\ \\hline \\end{tabular} \\end{center} \\end{table}   \\clearpage"
    },
    "0403/astro-ph0403617_arXiv.txt": {
        "abstract": "{We report a flare on the M9 dwarf DENIS 104814.7-395606.1, whose mass places it directly at the hydrogen burning limit. The event was observed in a spectral sequence during 1.3 hours. Line shifts to bluer wavelengths were detected in $\\mathrm{H_{\\alpha}}$, $\\mathrm{H_{\\beta}}$, and in the Na\\,{\\sc i} D lines, indicating mass motions. In addition we detect a flux enhancement on the blue side of the two Balmer lines in the last spectrum of our series. We interpret this as rising gas cloud with a projected velocity of about 100 $\\mathrm{km s^{-1}}$ which may lead to mass ejection. The higher Balmer lines $\\mathrm{H_{\\gamma}}$ to $\\mathrm{H_{8}}$ are not seen due to our instrumental setup, but in the last spectrum there is strong evidence for $\\mathrm{H_{9}}$ being in emission. ",
        "introduction": "DENIS 104814.7-395606.1 (hereafter DENIS 1048-39) was discovered in the DEep Near Infrared Survey (DENIS), which covered the southern sky \\citep{Epchtein} in two near infrared bands (J and $\\mathrm{K_{s}}$) and one optical band (I).  With this choice of bands DENIS is very sensitive to very low mass stars and brown dwarfs and thus excellently suited for searches for hitherto unknown low mass stellar or substellar objects in the solar vicinity. With a distance of only $4.6\\pm0.3$\\,pc \\citep{DENISinfr} DENIS 1048-39 is extremely close to the Sun. Classified as a M9 star, the lithium resonance line at \\mbox{6708 \\AA} could not be detected, therefore  substantial lithium depletion must have taken place \\citep{DENISdelfosse}. Compared to LP~944-20, which is also classified as M9, DENIS~1048-39 should be older and more massive. Theoretical models place DENIS 1048-39 directly at the hydrogen burning boundary with an estimated mass of 0.075 up to 0.09 $\\mathrm{M_{Sun}}$ and an age of 1 - 2 Gyrs \\citep{DENISinfr} assuming a solar chemical composition. Thus DENIS 1048-39 may either be among the most massive brown dwarfs or among the least massive stars.  Interestingly, the $\\mathrm{H_{\\alpha}}$-line was found to be variable \\citep{DENISdelfosse}, implying that DENIS 1048-39 exhibits activity despite its old age and low mass. There are only a few very late type objects showing strong and persistent H$_{\\alpha}$ emission.  H$_{\\alpha}$ emission, a well established activity indicator, is not detected in the majority of the stars. Instead one observes a steep decline of the strength of H$_{\\alpha}$ emission for stars later than spectral type M7. This effect shows up in  smaller equivalent widths (EW) of the H$_{\\alpha}$ line for these objects. However, since the H$_{\\alpha}$ line is seen against an increasingly faint photosphere for later objects, a better activity indicator is the ratio of the H$_{\\alpha}$ luminosity to the bolometric luminosity. This ratio was found to drop in only three subclasses (M8-L1) by one order of magnitude \\citep{Gizis}. The same authors also found that the activity of these late type objects is primarily related to temperature and shows only a secondary dependence on mass and age. Thus strong H$_{\\alpha}$ emission does not necessarily imply youth; on the contrary, strong H$_{\\alpha}$ emitters in the field are more likely to be old.  Despite this decline in activity there are many reports of flare activity even among L dwarfs. For very late-type M dwarfs \\citet{Gizis} estimated a flaring time fraction of about 7 \\%, while \\citet{Liebert1} found a flare duty cycle of about 1 \\% for L dwarfs, suggesting that there must be some ongoing magnetic activity in these dwarfs despite the absence of a persistent chromosphere or corona.   In this paper we report on a huge flare on DENIS 1048-39 detected in the H$_{\\alpha}$, H$_{\\beta}$, and sodium emission lines. In section 2 we describe the VLT data and their analysis, in section 3 we discuss the timing behaviour of the flare. ",
        "conclusions": "\\subsection{Emission geometry}  The measured half widths of the Na\\,{\\sc i} D and the H$_{\\beta}$ lines always stay below 20 $\\mathrm{km\\,s^{-1}}$, i.e., they are smaller than the measured rotational velocity. The same applies to the two main components of the H$_{\\alpha}$ line. This suggests that the emission for the Balmer lines as well as for the Na\\,{\\sc i} D lines is confined to a restricted region on the star. Since the H$_{\\alpha}$ line consists of two components and there is evidence that the H$_{\\beta}$ line consists of two components as well, the star could host even two active regions producing the observed Balmer emission. Now the question arises what the nature of these two active regions may be. There are two main possibilities: a dynamic scenario or a static scenario.  \\subsection{Static or dynamic scenario ?}  The dynamic scenario invokes two active regions located on the surface of the star and mass motions within these active regions. The Balmer emission comes from the chromosphere, nevertheless very close to the surface of the star.  At least in one of the two regions a brightening in the Balmer lines takes place on a timescale of about 1.5 hours. H$_{\\alpha}$ and H$_{\\beta}$ emitting material moves towards the observer and causes the blue shift.  The static scenario interprets the two regions as emitting gas clouds corotating in some distance to the stars surface as was first proposed for the K0 star AB~Dor by \\citet{Cameron}. Since for one component the line shift becomes bluer this one must be rising while the other one must be about to set behind the star since its line shift is moving towards the blue. The blue shifted component must then be captured during its rise above the horizon of the star to account for the flux increase. This latter static scenario can be excluded because of the rapid rotation of the stars as follows: For an estimated  radius of $\\mathrm{R_{\\star}} = 0.1 \\mathrm{R_{\\odot}}$ the maximal rotation period of DENIS 1048-39 is 4.9 hours. Therefore in the four consecutive spectra lasting together 80 minutes the star completes about a fourth of its rotation. If the emitting gas is confined in a corotating cloud, the measured velocity shift should exceed the star's rotational velocity in at least one spectrum for orbits near the equatorial plane. Moreover the flux increase cannot be explained by a cloud rotating into view because the star rotates too fast for this interpretation. A more natural explanation is therefore the scenario of active regions with mass motions at the surface of the star that brighten during the spectral series.  \\subsection{Blue-shifted Balmer emission}  Let us now consider the broad emission features bluewards of the two Balmer lines seen only in the fourth spectrum. Again a dynamic or static scenario may be considered. Let us first assume a static interpretation with a corotating cloud.  Since its radial velocity of about 100 $\\mathrm{km\\,s^{-1}}$ is too high to be interpreted by a region on or near the surface of the star, the emission would have to come from a cloud at some distance to the star. Such a cloud must be confined then by magnetic loops above the surface of the star and has rotated just into the field of view. If one assumes a stellar inclination close to $90^{\\circ}$ and a cloud in the equatorial plane, one can compute its distance $R$ from the rotation axis with $R=\\frac{\\mathrm{v} P}{2\\pi}$ with $\\mathrm{v}$ denoting the radial velocity and $P$ the rotation period of DENIS~1048-39. We find a distance of $R=4 \\mathrm{R_{\\star}}$ with the measured radial velocity of 100 $\\mathrm{km\\,s^{-1}}$, which is below the Keplerian corotation radius of $6 \\mathrm{R_{\\star}}$ like the clouds found on AB~Dor in $\\mathrm{H_{\\alpha}}$ \\citep{Donati} and other chromospheric lines \\citep{Brandt}. But since the star rotates fast, the feature should be visible in more than one spectrum since for its distance of four stellar radii it can be shadowed by the star only about 10 \\% of the time, i.\\,e. about 30 minutes, thus again a dynamic interpretation appears more likely.  In the dynamical interpretation the line shift is explained in terms of material ejected by the star. Since the emission in the main body of the $\\mathrm{H_{\\alpha}}$ brightens substantially in the same spectrum it is suggestive that the broad emission feature is connected to the active region producing the $\\mathrm{H_{\\alpha}}$ line arguing further in favor of a dynamic scenario. In the context of interpreting the broad emission feature as the signature of a rising cloud the question remains unsettled whether this rise leads to mass ejection since we do not know the longitude of the cloud. Assuming again an radius of $0.1 \\mathrm{R_{\\odot}}$ we find an escape velocity of about 550 $\\mathrm{km\\,s^{-1}}$ for a mass between 0.075 and 0.09 $\\mathrm{M_{\\odot}}$ while the projected velocity of the cloud is about 100 $\\mathrm{km\\,s^{-1}}$.  Let us now consider the width of the emission feature. This width can be due to temperature and turbulent broadening if one thinks of a confined gas cloud rising towards the observer. A first estimate of the temperature of a $\\mathrm{H_{\\alpha}}$ emitting gas is about 10\\,000 K leading to a thermal broadening of 9 $\\mathrm{km\\ s^{-1}}$, far less than the observed total line broadening  of 50 $\\mathrm{km\\ s^{-1}}$. In this scenario with a rising cloud one can estimate the height of the cloud after the twenty minutes exposure. Assuming a constant cloud velocity of the cloud and an ejection start right at the beginning of the exposure leads to a cloud height of 1.7 $\\mathrm{R_{\\star}}$. It is probable that the cloud has been decelerated during such a rise. Therefore a second interpretation of the line width is that different velocities of the cloud during deceleration are integrated over the exposure time. Moreover, the cloud may consist of more than one component with different velocities.  Besides these uncertainties in the event geometry there is no doubt that a dynamical interpretation of the spectra is needed. The last spectrum is then quite suggestively interpreted as  the onset of a flare on DENIS~1048-39 as reported for AD~Leo by \\citet{Houdebine} who found a similar flux enhancement in the far blue wing of the $\\mathrm{H_{\\gamma}}$ line during a particular violent flare on AD~Leo. Since these authors found projected velocities of up to 5800 $\\mathrm{km\\,s^{-1}}$ this event was clearly associated with a mass ejection, while this question remains unsettled for DENIS~1048-39.  \\subsection{Summary}  In conclusion, we find at least two active regions on DENIS~1048-39 contributing to the bulk of the Balmer line flux. Mass motions directed towards the observer are found for the emission lines of the Balmer series as well as for the Na D lines. In the last spectrum of our observations the onset of a flare seems to take place, since substantial brightening and blueshifts can be seen in the lines. In addition in the $\\mathrm{H_{\\alpha}}$ and the $\\mathrm{H_{\\beta}}$ line there is a broad emission feature on the blue side of the line. This can be interpreted as a rising  cloud.   Since DENIS~1048-39 seems to be located directly at the hydrogen burning limit this flare gives evidence that such events may be more ubiquitous than previously assumed. It is consistent with X-ray detections of brown dwarfs \\citep{Mokler} and the X-ray flare event found on the similar late-type star LHS~2065 \\citep{LHS-2065}."
    },
    "0403/hep-th0403234_arXiv.txt": {
        "abstract": "We discuss an inflation model, in which the inflation is driven by a single scalar field with exponential potential on a warped DGP brane. In contrast to the power law inflation in standard model, we find that the inflationary phase can exit spontaneously without any mechanism. The running of the index of scalar perturbation spectrum can take an enough large value to match the observation data, while other parameters are in a reasonable region. ",
        "introduction": "Astronomical observations in resent years, especially by Wilkinson Microwave Anisotropy Probe (WMAP) \\cite{wmap}, lead to a high precision era of cosmology. All the results imply there exists an accelerating moment (inflationary phase) in the very early universe. The angle power spectrum of cosmological microwave background (CMB) provides some evidences that the universe is almost spatially flat and that large scale structure is formed from a primordial spectrum of adiabatic, normal and nearly Harrison-Zeldovich (scale invariant) density perturbations. This can be explained by the simplest model of inflation~\\cite{liddlelyth}.    Despite the great success of the big bang standard model of cosmology together with inflation, there are still several serious problems in our present scenario of cosmology.  The cosmological constant (dark energy), unexpected low power spectrum at large scales, egregious running (to the common inflationary models) of power spectrum index are distinctive ones. The latter two might relate to the very high energy physics. Although many approaches have been made in the literature, fairly speaking, these problems still stick on.    In view of achievements and shortcomings of inflationary scenario in Einstein gravity, it is necessary to further deepen our understanding of the inflationary scenario from a theoretical perspective. In nonperturbative string/M theories there are topological solitons, called branes,  in $10$ or $11$ dimensional spacetime.  In the Horava-Witten model \\cite{horava}, gauge fields of the standard model are confined on two 9-branes located at the end points of an $S^1/Z_2$ orbifold, i.e., a circle folded on itself across a diameter. Inspired by the Horava-Witten model, the idea that our universe is a 3-brane embedded in a higher dimensional spacetime has received a great deal of attention in recent years~\\cite{braneworld}. In this brane world scenario, the standard model particles are confined on the 3-brane, while the gravitation can propagate in the whole space. In this picture, the Friedmann equation of the brane universe gets modified~\\cite{branecos}, compared to the one for the standard model. The chaotic inflation model on the RSII brane has been studied in~\\cite{chaotic}.  The inflation model in the brane world scenario with a Gauss-Bonnet term in the bulk has also been discussed recently by Lidsey and Nunes~\\cite{LN}.     In the RSII model~\\cite{RSII},  a brane with positive tension is embedded in five dimensional anti de Sitter space. Due to the warped effect of bulk geometry, general relativity on the brane can be recovered in low energies, while gravity on the brane is five dimensional in high energy limit. Compared to the RSII model, the brane world model proposed by Dvali, Gabadadze and Porrati (DGP)~\\cite{dgpmodel} is also very interesting. In the DGP model, the bulk is a flat Minkowski spacetime, but a reduced gravity term appears on the brane without tension. In this model, gravity appears 4-dimensional at short distances but is altered at distance large compared to some freely adjustable crossover scale $r_0$ through the slow evaporation of the graviton off our 4-dimensional brane world universe into an unseen, yet large, fifth dimension. The late-time acceleration is driven by the manifestation of the excruciatingly slow leakage of gravity off our four-dimensional world into an extra dimension. This offers an alternative explanation for the current acceleration of the universe.    In the DGP model, the gravitational behaviors on the brane are commanded by the competition between the 5-dimensional curvature scalar ${}^{(5)}R$ in the bulk and the 4-dimensional curvature scalar $R$ on the brane. At short distances the 4-dimensional curvature scalar $R$ dominates and ensures that gravity looks 4-dimensional. At large distances the 5-dimensional curvature scalar ${}^{(5)}R$ takes over and gravity spreads into extra dimension. As a result, Newton-like force law becomes 5-dimensional one. Thus, gravity begins weaker at cosmic distances. So it is natural that such a dramatic modification affects the expansion velocity of the Universe. In fact the DGP model has been applied to cosmology immediately after DGP putting forward their model \\cite{dgpcosmology}.    In this paper, we will study an inflation model in a brane world scenario, which combines the RSII model and DGP model.  That is, an induced curvature term will also appear on the brane in the RSII model. We call the brane as warped DGP brane. In this  model, inflation of the universe is driven by a single scalar field with exponential potential on the brane. In contrast to the power law inflation in standard model, we find that the inflationary phase exits spontaneously. The running of the index of the scalar perturbation spectrum can take negative values, which content the high precision observations by WMAP. At the same time other parameters are in a reasonable region.  The organization of this paper is as follows. In the next section, we present the Friedmann equation on the warped DGP brane. In Secs.~III and IV the inflation model is introduced. We analyze the parameter space allowed by the observation data in Sec. V. The paper ends in Sec.~VI with conclusions and discussions. ",
        "conclusions": "Various inflationary models have been proposed since Guth's seminal work~\\cite{Guth}. Going with the observations of high precision, we realize that a successful inflation model must at least possess the following properties: 1) a sufficient large number of e-folds, 2) a near Harrison-Zeldovich (scale invariant) spectrum, 3) a negative running of spectrum index. The model we discussed in this paper contents these. In particular, it is worth noting that in this inflation model based on the warped DGP brane world scenario, even for an exponential potential, the inflationary phase can exit naturally, and within a reasonable parameter region, the model can give us a negative running of the scalar spectrum index. By choosing more appropriate parameters in the model, we may obtain much better consistence with the current observation data. These features are quite attractive. Other properties of this model deserves further study. For example, it would be interesting to investigate carefully the relations of the spectrum index and its running to the comoving wave number $k$. In addition, in the warped DGP brane model, it is certainly of interest to further construct inflation models with scalar spectrum index larger than one at larger scales.   {\\bf Note added:} While we were finishing this paper, a paper~\\cite{new} appears on the net, which also discusses an chaotic inflation model on a brane with induced gravity. More recently, two related and interesting papers occur on the net. In \\cite{BMW} the authors calculate the amplitude of gravitational waves from brane-world inflation with induced gravity, while the authors of \\cite{BW} study the induced gravity with a non-minimally coupled scalar field on the brane.   {\\bf Acknowledgments:} This work was supported in part by a grant from Chinese academy of sciences, a grant No. 10325525 from NSFC, and by the ministry of science and technology of China under grant No. TG1999075401."
    },
    "0403/astro-ph0403329_arXiv.txt": {
        "abstract": "{ Three-dimensional hydrodynamic calculations are performed in order to investigate mass transfer in a close binary system, in which one component undergoes mass loss through a wind. The mass ratio is assumed to be unity. The radius of the mass-losing star is taken to be about a quarter of the separation between the two stars. Calculations are performed for gases with a ratio of specific heats $\\gamma=1.01$ and $5/3$.  Mass loss is assumed to be thermally driven so that the other parameter is the sound speed of the gas on the mass-losing star.  Here, we focus our attention on two features: flow patterns and mass accretion ratio, which we define as the ratio of the mass accretion rate onto the companion, $\\dot{M}_{acc}$, to the mass loss rate from the mass-losing primary star, $\\dot{M}_{loss}$.  We characterize the flow by the mean normal velocity of wind on the critical Roche surface of the mass-losing star, $V_R$.  When $V_R<0.4 A\\Omega$, where $A$ and $\\Omega$ are the separation between the two stars and the angular orbital frequency of the binary, respectively, we obtain Roche-lobe over-flow (RLOF), while for $V_R>0.7 A\\Omega$ we observe wind accretion. We find very complex flow patterns in between these two extreme cases. We derive an empirical formula of the mass accretion ratio as $0.18 \\times 10^{-0.75V_R/A\\Omega}$ ~in the low velocity regime and $0.05\\,(V_R/A\\Omega)^{-4}$ in the high velocity regime. ",
        "introduction": "Wind accretion plays an essential role especially in the evolution of detached binary systems such as symbiotic stars, precursor of peculiar red giants, $\\zeta$ Auriga stars and massive X-ray binaries. In the majority of symbiotic stars, the system contains a hot component - believed to be a mass-gaining white dwarf - and a cool component, the mass-losing red giant (e.g., Miko{\\l}ajewska \\cite{Miko} for a review). Recent studies indicate that many of the red giant components in symbiotic stars do not fill their Roche lobe, i.e., $R_{\\rm RG} \\la \\ell_1 /2$, where $R_{\\rm RG}$ is the radius of the red giant component and $\\ell_1$ is the distance of the inner Lagrangian point from the centre of the red giant (e.g., M\\\"urset \\& Schmid \\cite{1999A&AS..137..473M}). This possibly discloses that symbiotic stars are - with probably only the exception of T CrB (e.g., Anupama \\& Miko{\\l}ajewska \\cite{1999A&A...344..177A}) - well detached binary systems. Thus, mass transfer in symbiotic stars seems to be driven by wind, and wind mass loss is therefore a key ingredient for triggering symbiotic activity on a white dwarf companion. On the other hand, the mass accretion rates expected in the Bondi-Hoyle picture (Bondi \\& Hoyle \\cite{1944MNRAS.104..273B}) for such detached binary systems are usually much lower than those required to maintain the symbiotic activity of the systems. For example, Nussbaumer (\\cite{Nuss}) pointed out that in most symbiotic stars the typical rates of wind mass loss red giant are in the range $1 \\times 10^{-6} M_\\odot$ yr$^{-1}$ to $1 \\times 10^{-8} M_\\odot$ yr$^{-1}$ and that the expected accretion efficiency of wind capture is probably no better than one percent.  Hence, the accretion rates associated with most symbiotic white dwarfs are probably in the range of $1 \\times 10^{-8} M_\\odot$ yr$^{-1}$ to $1 \\times 10^{-10} M_\\odot$ yr$^{-1}$, which are too low to power the typical luminosities of symbiotics ($\\ga 1000 L_\\odot$) unless they are all in a late phase of hydrogen shell-flash (novae) on a white dwarf (e.g., Sion \\& Ready \\cite{1992PASP..104...87S}; Sion \\& Starrfield \\cite{1994ApJ...421..261S}). It will nevertheless be shown in the present paper that the mass accretion rate is smaller than that expected by a simple Bondi-Hoyle picture.   It has recently been suggested that some of the mass-accreting white dwarfs in symbiotic stars are the progenitors of Type Ia supernovae (e.g., Hachisu \\& Kato \\cite{2001ApJ...558..323H}).  On the other hand, it has long been discussed that the efficiency, the ratio of the mass captured by the white dwarf to the mass lost by the red giant, is very small (say, a few per cent) and not enough to increase the white dwarf mass to the Chandrasekhar mass limit (e.g., Kenyon et al. \\cite{1993ApJ...407L..81K}).  Wind velocities in cool red giants are typically a few times 10 km s$^{-1}$ and comparable to the orbital velocity in most symbiotic binary systems. Therefore, flow patterns are probably something between Roche lobe overflow (RLOF) and  Bondi-Hoyle wind accretion kinds of flows. The efficiency of capture in symbiotics should therefore not be estimated by the Bondi-Hoyle picture but by direct simulations of intermediate, complicated flow patterns.  Symbiotic stars are not the only systems in which wind accretion play a significant role. The most adopted model for the formation of peculiar red giants, Barium, CH and S stars, require a binary system containing an Asymptotic Giant Branch star (AGB) transferring mass via its stellar wind to a main sequence companion which becomes polluted in Carbon and s-process elements (see e.g. Boffin \\& Jorissen \\cite{BJ88}). The AGB then evolve into a white dwarf while the companion will appear as Carbon or s-process rich and, when on the giant branch, as a peculiar red giant. Boffin \\& Zacs (\\cite{BZ94}) have shown that an accretion efficiency of only a few percent is enough to explain present Barium stars. Objects related to these are the extrinsic S stars which also show symbiotic activity (e.g. Carquillat et al. \\cite{Carq98}) and bipolar planetary nebulae (e.g. Mastrodemos \\& Morris \\cite{MM98}).  $\\zeta$ Aurigae systems are another kind of system where wind accretion plays a role. In these eclipsing double-lined binaries, a main sequence star has an accretion wake produced by the wind of its K supergiant companion. The eponymous system, $\\zeta$ Aur, is a binary with a 972 days orbital period and containing a 5.8 M$_\\odot$ K4 Ib and a 4.8  M$_\\odot$ K B5 V star (Bennett et al. \\cite{Benn96}).  Finally, accretion wakes are also reported for high mass X-ray binaries (HMXB). These systems consist of a compact object, neutron star or black hole, orbiting a massive OB primary star which has a strong stellar wind. The X-ray emission is believed to be due to accretion of matter on the compact companion.  There have been already many numerical studies of stellar wind in a close binary system. Biermann (\\cite{1971A&A....10..205B}) computed two-dimensional stellar wind using a characteristic method. Sorensen, Matsuda \\& Sakurai (\\cite{1975Ap&SS..33..465S}) performed two-dimensional finite difference calculation of stellar wind emitted from a Roche lobe filling secondary. They used the Fluid in Cell (FLIC) method with first order of accuracy and a Cartesian grid. They computed Roche lobe overflow (RLOF) as well. Sawada, Hachisu \\& Matsuda (\\cite{1984MNRAS.206..673S}) calculated two-dimensional stellar wind from a contact binary using a Beam-Warming time implicit finite difference scheme and a generalized curvilinear coordinate. Their main goal was to compute the angular momentum loss rate from the system, which is an important factor to define the evolution of the binary system. Sawada, Matsuda \\& Hachisu (\\cite{1986MNRAS.221..679S}) computed two-dimensional stellar wind from a semi-detached binary system using the Osher upwind scheme and a generalized curvilinear coordinate. They observed a transition from RLOF to stellar wind by increasing the wind speed on the mass-losing star. The present study is a three-dimensional version of their study.  Matsuda, Inoue \\& Sawada, (\\cite{1987MNRAS.226..785M}) also conducted a similar study and discovered a flip-flop instability of the bow shock formed around a compact mass accreting object (see also Boffin \\& Anzer, \\cite{BA94}, and on the stability issue, Foglizzo \\& Ruffert, 1997 and 1999, and Pogolerov, Ohsugi \\& Matsuda, 2000). However, it was found that the phenomenon was characteristic to two-dimensional case (Matsuda et al., \\cite{1992MNRAS.255..183M}, Ruffert, \\cite{Ruffert96}). In the present work we perform three-dimensional calculation in which we do not observe the flip-flop instability.  These simulations concerned the case of a compact object moving in a wind and were mostly aimed at estimating the validity of the Bondi-Hoyle accretion rate. Binary effects were partially simulated by including velocity or density gradient in the wind. There are a few simulations however which simulate in full mass transfer by wind in a binary system (Blondin et al. 1991, Theuns \\& Jorissen 1993, Blondin \\& Woo 1995, Theuns et al. 1996, Walder 1997, Mastrodemos \\& Morris 1998, Dumm et al. 2000, Boroson et al. 2001, Gawryszczak, Miko{\\l}ajewska, \\& R{\\' o}{\\. z}yczka 2002). Theuns \\& Jorissen (\\cite{TJ93}) and Theuns, Boffin \\& Jorissen (\\cite{TBJ96}, TBJ96 in the following) performed three-dimensional hydrodynamic simulations of a 3 M$_\\odot$ AGB transferring mass through its stellar wind to a 1.5 M$_\\odot$ main sequence companion, in order to test the wind accretion model for the origin of peculiar red giants. They performed simulations using a polytropic equation of state with $\\gamma =$1., 1.1 and 1.5 and found mass accretion ratios  (that is, mass accretion rate/mass loss rate) of 1-2\\% for the $\\gamma =1.5$ case and 8 \\% for the other models, i.e. about ten times smaller than the theoretical Bondi-Hoyle estimate. They also observed that this value is dependent on resolution.  Walder (1997) presented simulations of wind accretion in well separated binaries for three different cases: a HMXB, $\\zeta$ Aur and a barium star progenitor. He obtains mass accretion ratio (mass accretion rate/mass loss rate) of 0.6, 3 and 6 \\%, respectively. Walder \\& Folini (2000) also presented a nice review of wind dynamics in symbiotics binaries, with an emphasis on the influence of the radiation field of the accretor.  More recently, Dumm et al. (\\cite{Dumm00}) presented three-dimensional Eulerian isothermal simulations in order to represent the symbiotic system RW Hya which seems to present an accretion wake, as indicated by the reduced UV flux observed at phase $\\Phi = 0.78$. They found that 6\\% of the M-giant wind is captured by the companion.  In this paper we perform three-dimensional numerical simulations of mass transfer by wind in a binary system with a mass ratio of unity. The full gravitational forces of the two components are taken into account. This is therefore different from the e.g. TBJ96 or Dumm et al. (2000) studies, where in order to account for the not very well known acceleration mechanism of the wind, the gravitational force of the primary is partially or totally reduced.  Here, we will concentrate on the derivation of the mass accretion ratio. We will leave for a following paper the physical discussion of the flow, including a discussion of angular momentum loss from the system. ",
        "conclusions": "\\begin{enumerate} \\item The flow pattern is classified by means of the normal speed of gas on the critical Roche surface, $V_R$. For $V_R<0.4A\\Omega$, we have a RLOF type flow, while for $V_R>0.7A\\Omega$, wind accretion flows are realized. In the intermediate parameter range, $0.4A\\Omega<V_R<0.7A\\Omega$, we observe very complex flows.  \\item We construct an empirical formula for the mass accretion ratio given by $f = 0.18\\times 10^{-0.75 V_R/A\\Omega}$, valid for low $V_R$.  \\item The calculated mass accretion ratio shows a kink at $V_R \\sim 2.6A\\Omega$. By using the nested grid, we found that this is because the assumed size of the mass-accreting box is comparable to or larger than the accretion radius. By reducing the size of the mass-accreting box for the nested grid, the position of the kink is shifted to larger $V_R$.  \\item For larger $V_R$ we obtain another empirical formula for the mass accretion ratio  $f = 0.05(V_R/A\\Omega)^{-4}$.  \\item For typical wind accretion flow, the results based on the coarse grid and those of the nested grid are essentially the same except the mass accretion ratio described above. For intermediate flow, the computed flow pattern depends on the resolution slightly. The amplitude of the oscillation of the flow increases with increased resolution because of reduced numerical viscosity. \\end{enumerate}"
    },
    "0403/astro-ph0403603_arXiv.txt": {
        "abstract": "We describe a technique to derive constraints on the differential emission measure (DEM) distribution, a measure of the temperature distribution, of collisionally ionized hot plasmas from their X-ray emission line spectra.  This technique involves fitting spectra using a number of components, each of which is the entire X-ray line emission spectrum for a single ion.  It is applicable to high-resolution X-ray spectra of any collisionally ionized plasma and particularly useful for spectra in which the emission lines are broadened and blended such as those of the winds of hot stars.  This method does not require that any explicit assumptions about the form of the DEM distribution be made and is easily automated. ",
        "introduction": "\\label{sec:intro} The X-ray emission for a large class of astrophysical objects is dominated by emission from a hot, thin plasma in which the electrons have a thermal velocity distribution and ion-electron recombination is balanced by collisional ionization --- i.e., collisional ionization equilibrium.  This class of astrophysical objects includes both hot and cool stars, clusters of galaxies, and elliptical galaxies as well as some cataclysmic variable systems and active galactic nuclei \\citep[see, e.g.,][]{pae03}.  A large class of laboratory plasmas are also in collisional ionization equilibrium.  At temperatures above $\\sim10^6$\\,K, plasmas in collisional ionization equilibrium emit X-rays in the 1.5--30\\,\\AA{} (0.4--8\\,keV) band.  Up to temperatures of $\\sim3\\ee8$, that X-ray emission is characterized by strong line emission.  This line emission varies strongly with temperature changes of a few tenths of a decade or more.  The unprecedented ability of the grating spectrometers on the \\chandra{} and \\xmm{} observatories to resolve this X-ray line emission present new possibilities for measuring temperatures of plasmas in CIE and, therefore, new tests of theories of astrophysical objects.  Compilations of X-ray line emissivities have long been available (e.g., \\citealt{ray77}, \\citealt*{mew85}, \\citealt{smi01}) and given a detailed model of an astrophysical object, it is, in general, straightforward to compute the emission line spectrum of any optically thin CIE plasma.  For the purpose of determining luminosities of most emission lines, the three dimensional temperature and density distributions may be reduced to the differential emission measure (DEM) distribution, a single-valued function of temperature which we define in \\S\\ref{sec:dem}.  However, it is often the case that no explicit theoretical prediction exists for the DEM distribution of objects of interest.  Furthermore, even when theory does provide an explicit prediction of the differential emission measure, it is often desirable to make measurements of the differential emission measure without respect to any model. Therefore, in this paper, we develop a method for obtaining constraints on the DEM distribution of a collisionally ionized plasma from its X-ray emission line spectrum without respect to any physical model.  Our method consists of fitting the observed spectral data using a model consisting of a continuum plus a number of line emission components, with each line emission component containing all of the lines of a single ion in the observed wavelength band.  Since each ion emits only in a specific temperature range, the best fit magnitude for the line emission of a given ion gives a measure of a weighted average of the DEM distribution in that temperature range times the abundance of that element.  We plot the constraints for all of the ions as a function of the temperature at which emission from the ion peaks. These plots may be understood as one-dimensional ``images'' of the DEM distributions in that they consist of discrete ``pixels'' (one for each ion) and differ from the true DEM distribution by a convolution which may be understood as a ``temperature-spread'' function.  In addition, by fitting the entire X-ray spectrum, rather than attempting to measure fluxes of individual lines, we take advantage of the information available from blended lines.  This is particularly important in the analysis of spectra in which the lines are significantly broadened and blended due to insufficient instrumental spectral resolution, Doppler motion of the emitting plasma, or other effects.  This method can be automated, facilitating the analysis of large numbers of spectra.  This attribute will become increasingly important with the future \\conx{} mission which will be able to obtain high-resolution spectra for a large number of objects.  While we are unaware of any use of this method exactly as it is described here, it is quite similar to and, in fact, inspired by methods described and applied by \\citet{pot63} for the analysis of solar ultraviolet spectra, by \\citet{sak99} for the analysis of an X-ray spectrum of the photoionized wind of the high mass X-ray binary Vela~X-1, and by \\citet{bri01} and \\citet*{beh01} for analyses of the X-ray spectra of the corona of the cool stars HR~1099 and Capella, respectively.  In \\S\\ref{sec:dem} we motivate and describe our method in detail and in \\S\\ref{sec:simtest}, we apply it to a simulated spectrum of a plasma with a continuous temperature distribution and simulated spectra of several single temperature plasmas.  In \\S\\ref{sec:discuss}, we discuss the possible applications of our method.  The X-ray emission of most hot stars originates in supersonic stellar winds, resulting in broadened and blended emission lines. Therefore, our technique is particularly useful in the analysis of these spectra and in a companion paper (hereafter, ``Paper II'') we apply our technique to spectra of nine hot stars. ",
        "conclusions": "\\label{sec:discuss}  We have described a method for deriving constraints on the differential emission measure distribution of a collisionally ionized plasma from its X-ray emission spectrum and displaying these constraints in a way that indicates the DEM distribution.  We have designed our technique to account for the line pumping that occurs at high densities and with high radiation intensities.  We have applied our method to simulated spectra and demonstrated that, with this method, we can recover simulated DEM distributions in, approximately, the temperature range $10^6$--$3\\ee{8}$\\,K from their emission line spectra subject to a temperature resolution comparable to the width of the line power functions: a few tenths of a decade.  We expect that the technique described here will be useful in the analysis of high-resolution spectra of plasmas in collisional ionization equilibrium in several astrophysical contexts.  This technique is particularly useful in cases where lines are broadened to the point of being blended.  Therefore, we expect this technique to be especially useful in the analysis of the X-ray spectra of hot stars and we apply it to spectra of nine hot stars in Paper II. \\citet{muk03} have shown that the X-ray spectra of some cataclysmic variable stars are characteristic of plasmas in collisional ionization equilibrium.  Some of those spectra show lines as broad as 500\\,\\kms{}.  Therefore, this technique may be useful in the analysis of those objects.  Even for spectra in which lines are not broad and blended, this technique has advantages over others.  It does not require fitting of individual lines to determine their fluxes and is therefore easily automated.  Also with this technique, it is not necessary to make assumptions or introduce biases about the form of the DEM distributions.  Therefore, our technique may also be useful for analysis of high-resolution spectra of the coronae of cool stars obtained with \\xmm{} and \\chandra.  The {\\it Astro-E II} observatory will be able to obtain high-resolution spectra of extended objects. Therefore, with the data from that observatory, our technique may be useful in the analysis of the spectra of clusters of galaxies and halos of elliptical galaxies.  The fact that this algorithm can be easily automated may be of particular importance when the DEM distributions need to be determined for a large number of spectra. The fact that the anticipated {\\it Constellation-X} observatory will have a very large effective area for high-resolution spectroscopy will enable it to obtain spectra for a large number of objects that are too faint to be efficiently observed with the three previously mentioned observatories.  Therefore, our technique may be especially useful in the analysis of the large number of high-resolution X-ray spectra that will be obtained with that observatory."
    },
    "0403/astro-ph0403059_arXiv.txt": {
        "abstract": "{The magnetic moment of the $^{13}$C nucleus is shown to provide a potentially useful tool for analysing quiescent cold molecular clouds. We report discovery of hyperfine structure in the lowest rotational transition of H$^{13}$CO$^{+}$. The doublet splitting in H$^{13}$CO$^{+}$, observed to be of width 38.5$\\pm$5.2 kHz or 0.133~km~s$^{-1}$, is confirmed by quantum chemical calculations which give a separation of 39.8 kHz and line strength ratio 3:1 when H and $^{13}$C nuclear spin-rotation and spin-spin coupling between both nuclei are taken into account. We improve the spectroscopic constants of H$^{13}$CO$^{+}$ and determine the hitherto uncertain frequencies of its low-$J$ spectrum to better precision by analysing the dark cloud L\\,1512. Attention is drawn to potentially high optical depths (3 to 5 in L\\,1512) in quiescent clouds, and examples are given for the need to consider the (1~--~0) line's doublet nature when comparing to other molecular species, redirecting or reversing conclusions arrived at previously by single-component interpretations.\\\\ We further confirm the hyperfine splitting in the (1~--~0) rotational transition of $^{13}$CO that had already been theoretically predicted, and measured in the laboratory, to be of width about 46 kHz or, again, 0.13~km~s$^{-1}$. By applying hyperfine analysis to the extensive data set of the first IRAM key-project we show that $^{13}$CO optical depths can as for H$^{13}$CO$^+$ be estimated in narrow linewidth regions without recourse to other transitions nor to assumptions on beam filling factors, and linewidth and velocity determinations can be improved. Thus, for the core of L\\,1512 we find an inverse proportionality between linewidth and column density, resp. linewidth and square root of optical depth, and a systematic inside-out increase of excitation temperature and of the $^{13}$CO:C$^{18}$O abundance ratio. Overall motion toward the innermost region is suggested. ",
        "introduction": "It has long been recognised that cold interstellar clouds are excellent laboratories for determining basic physical quantities of molecular structure. These clouds also provide the only access to molecular species that are too unstable to permit sufficient terrestrial production for in-depth investigation. A good case in point here is the X-ogen of the early 1970s, alias oxomethylium or the formyl cation, HCO$^{+}$ (e.g.~Snyder et al. 1976) which has since gained prime importance for studying the interactions between interstellar gas and magnetic fields.  In some of these clouds, microwave emission profiles are so narrow that line frequencies can be determined to between seven and eight significant digits, relative to other lines, corresponding to an accuracy of some ten meters per second in velocity, or Mach 0.02 at typical temperatures. A recent example is the re-analysis of the N$_2$H$^+$(1~--~0) transitions (Caselli et al. 1995) which led to a frequency determination with precision 7 kHz, difficult to attain with laboratory spectroscopy. It thus becomes possible to study minute motions in such objects that may be on the brink of collapsing toward star formation.  With this aim we have observed the dark cloud Lynds 1512 in numerous molecular transitions. This object is a small accumulation of opaque blobs on the POSS plates, situated near the edge of the Taurus-Auriga-Perseus complex. Ever since Myers \\& Benson (1983) discovered ammonia radiation from one of these blobs, this ``core'' has been known as a source of particularly narrow emission lines and has consequently attracted much attention, among many others that of Fiebig (1990) who from NH$_3$ observations derived a kinetic temperature of 9.7~K at the center, Muders (1995) who drew attention to the cloud's well-structured velocity field, and more recently from the extensive IRAM Key Project on pre-star-forming regions (Falgarone et al. 1998).  Somewhat to the East of the L\\,1512 ammonia peak there exists a region particularly well-suited for high-precision line studies because there, several tracers of high gas density like CCS, HC$_{3}$N and C$_4$H show very narrow emission profiles whose widths indicate the kinetic temperatures of the gas to be not more than 9 to 10~K. There is a certain velocity gradient all over L\\,1512 which must be subtracted off when determining intrinsic linewidths, but many coexisting species are subject to this same gradient and can therefore be compared to each other very well over the whole region. Their intrinsic profiles then turn out to be basically identical.  There is one high-density tracer however which does not at all fit this general picture of uniformity: H$^{13}$CO$^{+}$. We found its linewidth to be considerably broader than that of the other species, and its profiles to be non-Gaussian all over this region. Since HCO$^{+}$, being tied by its electric charge to collapse-retarding magnetic field lines, could be the prime informer on the interaction between neutral and charged molecules in dynamic clouds, it seemed important to clarify the reason for this difference. We thus realised that the magnetic moment of the $^{13}$C nucleus -- in conjunction with the much less efficient moment of the single proton -- might cause an observable hyperfine splitting in the H$^{13}$CO$^{+}$(1~--~0) line, of sufficient separation to be developed into a useful tool for spectral analysis. Subsequently we searched for a similar splitting in the (1~--~0) line of the low-density tracer molecule $^{13}$CO as well. Below we will report on some possibilities for line interpretation that emerge from the fact that again the splitting is wide enough to be observed in quiescent cold clouds.   This paper is organised as follows. In section 2 we summarise the observations performed at various telescopes, then present some general results in section 3. Section 4 discusses the observed H$^{13}$CO$^{+}$(1~--~0) line shapes, the determination of the width of the hyperfine splitting, and the ensuing high optical depths of this isotopomer. In section 5 our quantum chemical calculations concerning the hyperfine structure of this molecule are reported and improved spectroscopic constants derived. Line frequencies for the low-$J$ rotational spectra of H$^{13}$CO$^{+}$ (and HC$^{18}$O$^{+}$) can then be determined with high precision in section 6. Some examples of the analytical possibilities provided by the seeming complication due to hyperfine structure are given in section 7. Section 8 investigates the corresponding splitting for the $^{13}$CO molecule, which allows several physical trends to be established for the core of L\\,1512. Finally, section 9 presents our conclusions. ",
        "conclusions": "We wanted to show in this paper that the $^{13}$C nucleus' magnetic moment can provide a useful additional tool for analysing radiation from both the H$^{13}$CO$^{+}$ and the $^{13}$CO molecules. These species have not previously been investigated with consideration of this important detail. This is because the separation of their hyperfine components, 0.13~km~s$^{-1}$ for both species, is as narrow as the thermal width of cold interstellar gas and therefore a rather quiescent cold astronomical target, high spectral resolution, and good signal-to-noise ratios are all required to observe the splitting. When these conditions are fulfilled, however, narrowness of separation is hardly a handicap: while partial overlap of the components does of course degrade information near the profile's center, the noise in the totality of the line is also less than if the components were distinctly apart. One should also note that line strength ratios 3:1 or 2:1, as in the present cases, are next to ideal for retrieving astronomical quantities like optical depth. Ratios much closer to one would not allow clear determination of depth by comparison of amplitudes of the hyperfine components, much larger ones on the other hand would carry the risk of correlating gas regimes that might in reality be quite distinct, one of high optical depth, the other transparent.  For narrow-line regions, the determination of velocities profits most directly from knowing the hyperfine structure underlying a profile. Therefore we were able to present H$^{13}$CO$^{+}$ rotational transition frequencies of considerably improved precision by comparing to other, well-defined lines in L\\,1512, as well as to verify astronomically the $^{13}$CO(1~--~0) doublet frequencies first predicted by theory. This precision in turn allows studies of the internal motions of very quiescent clouds that may be extremely subsonic but, as was shown above for the case of L\\,1512, may nevertheless indicate the onset of collapse toward the center of a potentially star-forming core.  The inadequacy of interpreting a close-doublet profile by singlet analysis becomes very evident in the derivation of linewidths. Here, neglect of hyperfine structure in the profile can easily entail considerable overestimates which in turn lead to erroneous values of column density and of virial mass (the latter being proportional to the square of intrinsic linewidth). This effect is compounded by the frequently necessary neglect of line broadening caused by nonzero optical depth. In the case of the two $^{13}$C-containing species discussed here, the direct handle on optical depth that is provided by the nuclear magnetic moment may permit a more reliable determination of the intrinsic linewidth and its derivatives. For this reason, we were able to establish a width-depth or width-column density relation for the core of L\\,1512 that may eventually make the transition from a completely quiescent central core to a turbulent envelope accessible to observations. This more direct method of obtaining optical depths, rather than by the standard method for $^{13}$CO, allowed us to recognise the L\\,1512 core as considerably more bound than suspected since core mass has to be revised upwards and virial mass downwards upon considering line multiplicity correctly. Likewise we tried to argue that measurements in the dark core of L\\,1544 which seem to set H$^{13}$CO$^{+}$ quite apart from numerous other molecular species may in reality fall nicely in line when hyperfine splitting is properly considered.  A reestimate of H$^{13}$CO$^+$ linewidths also suggests that chemical evolution of collapsing cores over time should indeed be observable through the comparison of these linewidths with those of other molecules. Further, recent attempts to derive magnetic field strengths in molecular clouds by comparing ion linewidths with those of neutrals (H$^{13}$CO$^+$ vs. H$^{13}$CN) should profit from taking the doublet nature of the ion into consideration. For all these reasons it is hoped that the hyperfine structure in the profiles of the species discussed in this paper will motivate further high-resolution studies of cold dark clouds."
    },
    "0403/astro-ph0403090_arXiv.txt": {
        "abstract": "{We initiated a comprehensive state of the art binary population synthesis study of white dwarf main-sequence star (WDMS) binaries to serve as a foundation for subsequent studies on pre-cataclysmic variables, double white dwarfs, and white dwarf + B-star binaries. We considered seven distinct formation channels subdivided into three main groups according to the evolutionary process that gives rise to the formation of the white dwarf or its helium-star progenitor: dynamically stable Roche-lobe overflow (Algol-type evolution), dynamically unstable Roche-lobe overflow (common-envelope evolution), or stellar winds (single star evolution). For each formation channel, we examine the sensitivity of the population to changes in the amount of mass lost from the system during dynamically stable Roche-lobe overflow, the common-envelope ejection efficiency, and the initial mass ratio or initial secondary mass distribution. In the case of a flat initial mass ratio distribution, the local space density of WDMS binaries is of the order of $\\sim 10^{-3}\\, {\\rm pc}^{-3}$. This number decreases to $\\sim 10^{-4}\\, {\\rm pc}^{-3}$ when the initial mass ratio distribution is approximately proportional to the inverse of the initial mass ratio. More than 75\\% of the WDMS binary population originates from wide systems in which both components essentially evolve as if they were single stars. The remaining part of the population is dominated by systems in which the white dwarf is formed in a common-envelope phase when the primary ascends the first giant branch or the asymptotic giant branch. When dynamically stable mass transfer proceeds highly conservative and the common-envelope ejection process is very efficient, the birthrate of WDMS binaries forming through a common-envelope phase is about 10 times larger than the birthrate of WDMS binaries forming through a stable Roche-lobe overflow phase. The ratio of the number of helium white dwarf systems to the number of carbon/oxygen or oxygen/neon/magnesium white dwarf systems derived from large samples of observed WDMS binaries by, e.g., future planet-search missions such as SuperWASP, COROT, and Kepler may furthermore constrain the common-envelope ejection efficiency. \\begin{keywords} binaries: general -- stars: evolution -- stars: white dwarfs -- methods: statistical \\end{keywords} } ",
        "introduction": "In recent years, substantial numbers of detached white dwarf main-sequence star (WDMS) binaries have been detected in large-scale surveys searching for pre-cataclysmic variables (Hillwig et al. 2000, Raymond et al. 2003), double degenerates (Saffer et al. 1998), low-mass white dwarfs (Marsh et al. 1995, Maxted et al. 2000), planetary nebulae (Bond \\& Livio 1990, Livio 1992), or even the dark matter content of the Galaxy's heavy halo (Silvestri et al. 2002). Upcoming surveys looking for transiting extrasolar planets are furthermore expected to contribute further to this rapidly growing database (Farmer \\& Agol 2003). So far, the majority of WDMS binaries found consist of a white dwarf with a low-mass companion. White dwarfs in binaries with main-sequence secondaries more massive than $0.3\\,M_\\odot$ are generally harder to detect due to the large luminosity difference between the component stars (Iben et al. 1997, Marsh 2000). Despite this obstacle, several claims have recently been made on the possible detection of white dwarf binaries with bright companions of spectral types as early as B(e) (e.g. Holberg et al. 1998, Vennes et al. 1998, Burleigh \\& Barstow 2000, Burleigh et al. 2001, Shobbrook et al. 2003; and references therein).  WDMS binaries with early B(e)-type companions more massive than $\\sim 10\\,M_\\odot$ are thought to be the end product of binaries in which a giant-type star transfers mass to its initially less massive companion. The mass-transfer phase is responsible for exposing the giant's core as a white dwarf and for increasing the mass of the secondary so that it evolves towards earlier spectral types. In addition, transport of angular momentum may spin the secondary up to form a rapidly rotating Be star (Waters et al. 1988, Pols et al. 1991). The predicted number of WDMS binaries forming in this way is, however, rather small (e.g. de Kool \\& Ritter 1993, Figs. 2 and 4) which may pose a problem if the recent claims on the detection of white dwarf + B-star binaries are confirmed. In addition, as mass transfer is expected to increase the orbital period it is hard to explain the formation of short-period white dwarf + early B-star candidates such as HD\\,161103 ($P_{\\rm orb} = 4.7$ days, Shobbrook et al. 2003) through this evolutionary channel.  The wealth of existing and expected new data combined with new up-to-date stellar evolution tools (see Hurley et al. 2000) render the WDMS binaries a potentially fruitful subject of a comprehensive binary population synthesis study. A systematic exploration of the formation channels in particular may shed new insights on the origin of short-period white dwarf + B-star binaries. Other interesting subclasses of detached WDMS binaries include pre-cataclysmic variables and progenitors of double degenerates.  To the best of our knowledge, de Kool \\& Ritter (1993) were the first to devote an entire study to the population of WDMS binaries as a class on its own. The authors used a Monte Carlo type simulation in which the initial binary parameters were randomly generated from observed orbital period, primary mass, and secondary mass distributions; and the evolution of the binaries was approximated by analytical fits to detailed stellar evolution tracks as described by Politano (1988) and de Kool (1992). The method resulted in distribution functions for the formation rates and space densities of detached WDMS binaries (or their direct progenitors) as a function of their expected orbital and stellar parameters.  In this paper, our aim is to update the population synthesis study by de Kool \\& Ritter (1993) by using the more recent and much more detailed analytical fits to stellar evolution derived by Hurley et al. (2000). We study in more detail the differences between the evolutionary channels leading to the formation of WDMS binaries, and derive the formation rates and absolute numbers of systems for each of the formation channels separately. In addition, we consider a much more elaborate set of population synthesis models characterised by different assumptions about the fate of the mass transferred during dynamically stable Roche-lobe overflow, different common-envelope ejection efficiencies, and different initial mass ratio or secondary mass distributions.  Our method furthermore differs from that by de Kool \\& Ritter (1993) in the implementation of the observed distributions of the initial masses and orbital periods: instead of randomly generating them from the observed distribution functions at the onset of the binary evolution calculations, we start our simulations from a logarithmically spaced 3-dimensional grid of initial masses and orbital periods in which each set of binary parameters has an equal probability. The likelihood of the different initial binary parameters is then taken into account by weighting the contribution of each WDMS binary found in the simulations according to the adopted primary mass distribution, initial mass ratio or secondary mass distribution, and initial orbital period distribution.  In doing so, we obtain a better and more uniform sampling of the entire parameter space, which particularly benefits the resolution of low-density tails in the distribution functions describing the masses and orbital periods of WDMS binary components. This procedure also allows us to vary the distribution functions for the initial masses and orbital periods without having to repeat the binary evolution calculations.  The results of this study are intended to serve as a foundation for more focused population synthesis studies of white dwarf + B-star binaries (Willems et al., in preparation), double white dwarfs (Willems et al., in preparation), and pre-cataclysmic variables. Hence, we here concentrate on the effect and the uncertainties associated with the different model parameters in general. In subsequent papers we hope to constrain these uncertainties by comparing our results with different types of WDMS binaries and WDMS binary descendants. The results presented here therefore do not necessarily represent the best possible fit to observationally derived binary parameters.   The plan of the paper is as follows. In Section~2, we present the basic ingredients and the assumptions adopted in the binary population synthesis calculations. In Section~3, we describe the different evolutionary channels leading to the formation of WDMS binaries as well as the ranges of stellar masses and orbital periods available to each formation channel. Section~4 deals with the effects of varying degrees of non-conservative mass transfer and different common-envelope ejection efficiencies on the population of WDMS binaries. In Section~5, we estimate the order of magnitude of the formation rates and the expected number of WDMS binaries currently populating the Galaxy. Section~6 briefly addresses the expected luminosity differences between the white dwarf and its companion, as well as their expected orbital radial-velocity variations. A brief summary of our results and some concluding remarks are presented in Section~7. ",
        "conclusions": "We used the BiSEPS binary population synthesis code described by Willems \\& Kolb (2002) to study the population of detached white dwarf main-sequence star binaries forming through seven distinct evolutionary channels. In six of these, the white dwarf is formed through binary interactions resulting from Roche-lobe overflow. The channels mainly differ in the stability and the outcome of the mass-transfer phase that eventually gives rise to the formation of the white dwarf. The six formation channels are depicted schematically in Figs.~\\ref{ch1} and \\ref{ch2}--\\ref{ch6}. The seventh evolutionary channel is characterised by the absence of any type of binary interactions other than mass and angular momentum exchange via a stellar wind. Binaries evolving through this channel are typically very wide, so that the white dwarf forms in much the same way as it would if the primary were a single star.  Our results show that the wide non-interacting binaries generally comprise more than 75\\% of the total WDMS binary population. The remaining part of the population is dominated by systems undergoing a common-envelope phase when the white dwarf progenitor ascends the first giant branch or the asymptotic giant branch. The total number of systems currently populating the Galaxy, the relative number of systems evolving through each formation channel, and the birthrates of systems forming through each channel are given in Tables~\\ref{total1}--\\ref{br1} for different different assumptions about the fate of the mass transferred during dynamically stable Roche-lobe overflow, different common-envelope ejection efficiencies, and different initial mass ratio or secondary mass distributions. The total number of systems is of the order of $\\sim 10^9$ when a flat initial mass ratio distribution $n(q) = 1$, for $0 < q \\le 1$, is considered. This number decreases to $\\sim 7 \\times 10^7$ for an initial mass ratio distribution of the form $n(q) \\propto q^{-0.99}$, for $0 < q \\le 1$. An even stronger decrease is found for the birthrates of systems forming through the different formation channels. Besides the different assumptions regarding the treatment of dynamically stable and unstable mass transfer, the total number of WDMS binaries currently populating the Galaxy also depends on the finite life time of the systems which, in the case of short-period systems, may be affected by the adopted magnetic braking prescription. A detailed analysis of the effects of different braking laws is, however, beyond the scope of this investigation.  The total number of Galactic WDMS binaries evolving through a common-envelope phase is of the order of $\\sim 2 \\times 10^8$ when a flat initial mass ratio distribution $n(q) = 1$ is considered, and of the order of $\\sim 7 \\times 10^6$ when an initial mass ratio distribution of the form $n(q) \\propto q^{-0.99}$ is considered. Under the assumption that the Galaxy has an effective volume of $5 \\times 10^{11}\\, {\\rm pc^3}$, these numbers correspond to approximate local space densities of $4 \\times 10^{-4}\\, {\\rm pc^{-3}}$ and $10^{-5}\\, {\\rm pc^{-3}}$, respectively. The latter space density is comparable to the space density of $6 - 30 \\times 10^{-6}\\, {\\rm pc^{-3}}$ derived by Schreiber \\& G\\\"ansicke (2003) from an observed sample of 30 post-common-envelope binaries with low-mass main-sequence secondaries. However, since the space densities derived by Schreiber \\& G\\\"ansicke (2003) are likely to be lower limits, the space density of $4 \\times 10^{-4}\\, {\\rm pc^{-3}}$ is not necessarily in disagreement with observations. We also note that 12 of the 30 post-common-envelope binaries considered by Schreiber \\& G\\\"ansicke (2003) contain an sdOB primary instead of a white dwarf. These sdOB + main-sequence star systems are not included in our simulated WDMS binary samples.  We furthermore find that the ratio of the number of He white dwarf systems to the number of C/O or O/Ne/Mg white dwarf systems is weakly dependent on the amount of mass lost from the system during dynamically stable Roche-lobe overflow and strongly dependent on the common-envelope ejection efficiency. We therefore propose that comparison of the observationally derived ratio from, e.g., large-scale planet-search mission as SuperWASP, COROT, and Kepler with the ratios presented in Fig.~\\ref{fratio} may yield constraints on the common-envelope ejection efficiency, provided that the efficiency is low. For higher efficiencies, the role of the ejection efficiency parameter cannot be distinguished from the role of the initial mass ratio or initial secondary mass distribution. We note, however, that the theoretically predicted ratios are derived under the assumption of a constant envelope binding-energy parameter and a constant envelope-ejection parameter. In addition, the derivation of constraints on the common-envelope ejection efficiency requires the detection of WDMS binaries with luminosity ratios down to $10^{-5}$, which poses a severe and not easily overcome observational challenge.  This paper is to serve as a first comprehensive step in a study of more specific subclasses of WDMS binaries and binaries descending from them. In future investigations we will address WDMS binaries in the context of pre-cataclysmic variables, double white dwarfs (Willems et al., in preparation), and white dwarf + B-star binaries (Willems et al., in preparation). Other potentially interesting applications are the study of type Ia supernovae progenitors through both the 'standard' evolutionary channel and the recently proposed thermal time scale mass-transfer channel (King et al. 2003)."
    },
    "0403/astro-ph0403573_arXiv.txt": {
        "abstract": "Cool giant stars are highly dynamical objects, and complex micro-physical processes play an important role in their extended atmospheres and winds. The interpretation of observations, and in particular of high-resolution IR spectra, requires realistic self-consistent model atmospheres. Current dynamical models include rather detailed micro-physics, and the resulting synthetic spectra compare reasonably well with observations. A transition from qualitative to quantitative modelling is taking place at present. We give an overview of existing dynamical model atmospheres for AGB stars, discussing recent advances and current trends in modelling. When comparing synthetic spectra and other observable properties resulting from dynamical models with observations we focus on the near- and mid-IR wavelength range. ",
        "introduction": "Cool giant stars are prime targets for high-resolution infrared spectroscopy: these stars are luminous, emit most of their radiation at near- and mid-infrared wavelengths, and they have highly dynamical atmospheres due to convection and pulsation, leading to strong spectral variations. The coolest and most luminous of these objects also have slow, massive stellar winds. They represent an important but not so well-understood stage in the evolution of low- and intermediate mass stars, and they contribute significantly to the chemical enrichment of the interstellar medium. These are good reasons for studying cool giants but the interpretation of the observed spectra requires detailed and consistent model atmospheres for this type of stars.  Several of the very reasons that make pulsating asymptotic giant branch (AGB) stars interesting targets for IR spectroscopy, however, cause major challenges for realistic, self-consistent modelling of their dynamical atmospheres and circumstellar envelopes. Many interacting processes have to be taken into account simultaneously: propagating shock waves caused by stellar pulsation modify the structures of the atmospheres on local and global scales, leading to strong deviations from hydrostatic stratification. The radiative fields are dominated by molecular opacities or even by dust grains forming in the cool outer layers of the atmospheres. Important micro-physical processes like gas-phase chemistry and dust formation may be severely out of equilibrium.  Pioneering models for dynamic atmospheres of AGB stars date back more than two decades, but these models were severely restricted by the lack of sufficiently powerful computers and micro-physical input data. Only recently it has become possible to calculate consistent dynamic models which produce reasonably realistic synthetic spectra for a wide range of stellar parameters. This progress is of particular interest in the context of current instrumental developments in IR spectroscopy. ",
        "conclusions": "Since the beginnings of time-dependent dynamical modelling more than two decades ago, dynamic model atmospheres for cool giant stars have overcome many physical and computational obstacles. Within the last few years a transition from qualitative to consistent quantitative modelling is taking place. Current models include detailed micro-physics and some non-equilibrium processes in the computations of atmospheric structures and stellar winds. The resulting synthetic spectra compare reasonably well with observations, but certain consistency issues remain unsolved so far. This has to be kept in mind when using dynamical models for the interpretation of observations.  In addition to the spherical models with detailed input physics, first prototypes of global 3D radiation-hydrodynamical models ('star-in-a-box') have been computed recently, investigating the dynamics of giant convection cells and their influence on atmospheric structure and dust formation. These 3D models are based on simpler micro-physics and have a lower spatial resolutions than the spherical models. This trade-off is necessary to keep computation times within acceptable limits.  In summary, the present status of dynamic model atmospheres and the developments that can be expected within the next few years look promising, regarding the interpretation of high-resolution IR spectra."
    },
    "0403/astro-ph0403023_arXiv.txt": {
        "abstract": "The evolution of a supermassive black hole in the center of galaxy is considered. We analyze  the  kinetic equation describing relaxation processes associated with stellar encounters. The initial distribution function of stars is assumed to be isothermal, and it changes in accordance with consumption of stars by the black hole. We calculate the consumption rate and obtain the law of growth of supermassive black hole, i.e. the dependence of the black hole mass on time and on the velocity dispersion of stars. The obtained values of mean black hole mass and mean rate of stars consumption by a black hole are in reasonable agreement with observational data. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403509_arXiv.txt": {
        "abstract": "{One of the major challenges for theoretical modeling of Type Ia supernova explosions is to explain the diversity of these events and the empirically established correlation between their peak luminosity and light curve shape. In the framework of the so-called Chandrasekhar mass models, the progenitors carbon-to-oxygen ratio has been suggested to be a principal source of peak luminosity variations due to a variation in the production of radioactive $^{56}$Ni during the explosion. The underlying idea is that an enhanced carbon mass fraction should result in a more vigorous explosion since here the energy release from nuclear reactions is increased. It was suspected that this would produce a higher amount of  $^{56}$Ni in the ejecta. In this letter we describe a mechanism resulting from an interplay between nucleosynthesis and turbulent flame evolution which counteracts such an effect. Based on three-dimensional simulations we argue that it is nearly balanced and only minor differences in the amount of synthesized $^{56}$Ni with varying carbon mass fraction in the progenitor can be expected. Therefore this progenitor parameter is unlikely to account for the observed variations in Type Ia supernova luminosity. We discuss possible effects on the calibration of cosmological measurements.} ",
        "introduction": "\\label{intro_sect}  The question of the origin of the observed diversity of Type Ia supernovae (SNe Ia) has been a long standing problem. Although these astrophysical events possess a surprising homogeneity of features, which led to the application of SNe Ia as distance indicators in cosmology, variations of  photometric and spectroscopic data have been reported. Nevertheless, corrections of the luminosity based on an empirical relation between peak luminosity and light curve shape \\citep{phillips1993a} made SNe Ia a major tool of observational cosmology (for a recent review see \\citealt{leibundgut2000a}). A theoretical understanding of this correlation is still lacking. Thus the challenge for astrophysical theory is to construct a SN Ia explosion model that is robust and---as a first step---able to explain the observed variations.  It is generally agreed on the fact that SNe Ia originate from thermonuclear explosions of white dwarf (WD) stars, although many scenarios have been proposed for the particular realization (see \\citet{hillebrandt2000a} for a recent review). The currently favored model is a binary system in which a WD composed of carbon and oxygen accretes matter from a non-degenerate companion until it reaches the Chandrasekhar mass. At this point, thermonuclear burning at the center of the star develops into a flame which propagates outward. From the viewpoint of hydrodynamics, two modes of flame propagation are possible: a subsonic deflagration, in which the reaction is mediated by thermal conduction of the degenerate electrons, and a supersonic shock-induced detonation. \\cite{arnett1969a} noted that to produce the observed intermediate mass elements a prompt detonation can be excluded. Therefore the flame starts out in the deflagration mode and may or may not develop into a detonation. The key feature of this model is that buoyancy-induced instabilities lead to a turbulization of the flame front in the deflagration mode. Only due to this effect the flame is accelerated sufficiently to power SN Ia explosions.  \\begin{figure*}[t] \\centerline{ \\includegraphics[width = 0.73 \\linewidth] {fig1.eps}} \\caption{Temporal evolution of the burning front. \\label{evo_fig}} \\end{figure*}  In the first stages the prevailing densities and temperatures are high enough that the reaction terminates in nuclear statistical equilibrium (NSE), consisting mainly of iron group elements. As the WD expands, density and temperature drop and the thermonuclear reactions produce intermediate mass elements such as Si, S, and Ca.    The light curve of SNe Ia is powered by the radioactive decay of $^{56}$Ni and $^{56}$Co. The peak luminosity is a measure of the $^{56}$Ni produced by the explosion (``Arnett's law'', \\citealt{arnett1982a}). Parameters that have frequently been suggested to affect the amount of $^{56}$Ni synthesized are the carbon-to-oxygen (C/O) ratio of the WD, its central density prior to ignition, and its metallicity. Of course, other parameters like rotation and asphericity of the explosion may play a role, too. A thorough investigation of the impact of these parameters is mandatory for the validation of cosmological measurements, since here possible evolution effects with cosmic age are critical.  The objective of the present study is to explore the effect of the progenitor's C/O ratio on the supernova explosion by means of three-dimensional hydrodynamical models. Since we are mainly interested in the explosion energy and the amount of synthesized $^{56}$Ni, it is justified to focus on the deflagration stage, leaving aside a possible delayed detonation. Only in this first stage of the SN Ia explosion the prevailing densities and temperatures are sufficient for burning to iron group elements, and therefore the main part of the $^{56}$Ni and probably also of the explosion energy is produced here. ",
        "conclusions": "In this letter we described a mechanism that is responsible for the somewhat surprising fact that although the progenitor's C/O ratio affects the energy release of SN Ia explosions, it has little effect on the peak luminosity determined by the $^{56}$Ni mass. This is found with help of three-dimensional simulations and disagrees with the results reported from one-dimensional models. \\citet{hoeflich1998a} find a 14\\% decrease in $M_\\mathrm{Ni}$ when reducing the C/O ratio from 1/1 to 2/3; the fact that they considered a delayed detonation model should not greatly affect the comparability for reasons given in Sect.~\\ref{intro_sect}. The ``working hypothesis'' of an increased $^{56}$Ni production with higher $X(^{12}\\mathrm{C})$ established by \\citet{umeda1999b} cannot be confirmed. The reason for this discrepancy is a complex interplay between nucleosynthesis and nonlinear flame evolution which can be modeled appropriately only in three dimensions. The models presented here are based on a rather simplistic description of the nuclear reactions, which will be improved in forthcoming models. Nevertheless, the features that are important for the explosion dynamics are taken into account and trends can be revealed with this approach.  An interesting question is how our results affect the peak luminosity--light curve shape relation of SNe Ia. Since no light curves were calculated from our models we can only speculate on the effects based on the trends found by \\citet{arnett1982a} and \\citet{pinto2000a} by means  of analytic studies. According to ``Arnett's law'' the peak luminosity should reflect the amount of radioactive $^{56}$Ni synthesized in the explosion. Since in our models the evolution of the densities of the ashes are very similar, the little variation in the iron group elements will result in little variation of the $^{56}$Ni mass. Thus the peak luminosity should be roughly the same for our three models. The explosion energy, however, varies significantly and a greater expansion at late stages of the explosion associated with a larger C/O ratio would result in a more rapid decline of the light curve but also in a somewhat higher luminosity. This is in contradiction to the empirically established relation that for brighter SNe Ia the light curve declines slower, which has been extensively applied to calibrate cosmological distance measurements.  \\citet{timmes2003a} analytically found a linear dependence of the amount of produced $^{56}$Ni on the progenitor's metallicity. This in turn is, however, not likely to affect the explosion energy and dynamics significantly. Of course, from the viewpoint of stellar evolution there is a interrelation between the progenitor's metallicity and other parameters, including its carbon mass fraction. Therefore it is plausible that the ``luminosity-width relation'' is caused by a combination of different parameters and that a direct relation between the $^{56}$Ni mass and the peak luminosity might be oversimplified. It seems well possible that the distribution of the $^{56}$Ni due to three-dimensional effects plays an important role here. In a subsequent publication we will report on a systematic parameter study addressing these questions. Ultimately, SN Ia explosion models will have to be coupled to realistic stellar evolution of the progenitor system and the calculation of synthetic light curves becomes mandatory."
    },
    "0403/astro-ph0403215_arXiv.txt": {
        "abstract": "We probe the relationship between surface magnetic fields and the X-ray emitting corona in the rapidly rotating star AB Dor. Circularly polarised spectra have been inverted to produce a surface (photospheric) magnetic field map. This surface map has been extrapolated to model AB Dor's coronal field topology and  X-ray light curve. {\\em Chandra}/LETG light curves of AB Dor from the same epoch show intrinsic variability at the 30\\% level. Period analysis indicates a fraction of this is due to rotational modulation. We measure  velocity shifts in emission line centroids as a function of ${ P}_{\\rm rot}$ and find evidence of rotational modulation (max.~vel. $\\sim 40 \\pm 13$\\,km\\,s$^{-1}$). This modulation may indicate the presence of a  localised X-ray emitting region at mid-to-high latitudes. ",
        "introduction": "We present the first results from a co-ordinated optical and X-ray campaign to probe the corona of a rapidly rotating, active cool star. New X-ray and EUV observations reveal a picture of stellar coronae unlike anything on the Sun. Emission measures of about $10^{52}$\\,cm$^{-3}$ are found in plasma at temperatures of 8\\,MK and beyond, even in relatively quiescent coronae (e.g. Sanz-Forcada et al. 2003). Emission measures peak at $10^{50}$\\,cm$^{-3}$ at 2.5\\,MK on the Sun  {\\em in active regions} (Drake et al. 2000). While the thermal properties of  active stellar coronae are increasingly well-determined, we have yet to establish where this emission originates and to confront observations with theoretical models.  Previous studies provide evidence for both extended and compact coronae in active cool stars. The extended coronal model is supported by observations of fast moving cool ($10^4$\\,K) H$\\alpha$ absorption transients in young stars. These are caused by prominences rigidly corotating with the star at heights beyond the Keplerian corotation radius (Collier Cameron et al. 2001). Further support for extended coronae comes from time-resolved HST spectroscopy of the eclipsing binary, V471 Tau (K2V+DA). The hot white dwarf acts as a probe of circumstellar material around the K2V enabling Walter ({\\em this volume}) to detect transition region temperature plasma ($10^4<T<10^5$\\,K) above  2\\,$R_*$.   \\begin{figure} \\plotone{hussainfig1.eps} \\caption{[{\\em bottom row}] Surface magnetic field maps  from 2002~Dec~11 are extrapolated to model AB Dor's 3-D coronal field,  generate ``X-ray snapshots''   [{\\em middle row}], and  predict its X-ray light curve  [{\\em top row}]. } \\end{figure}  Support for compact coronae comes from X-ray light curves of the eclipsing binary, Algol (B8V+K2III). Flaring regions are found to have heights, $0.1<H<0.6$\\,$R_*$ above the K star (Schmitt \\& Favata 1999, Schmitt et al. 2003). Maggio et al. (2000) find evidence for a small flaring region ($H<0.3$\\,$R_*$) on the single, K0V, AB Dor. likely to be located near its pole. Electron density diagnostics measured in $10^7$\\,K plasma in active stars are typically very high ($n_e>10^{12}$\\,cm$^{-3}$). Densities this high are more easily explained in terms of compact X-ray emitting regions (e.g. Sanz-Forcada et al. 2003).  Brickhouse et al. (2001) find evidence for compact, quiescent coronae in active stars in a {\\em Chandra}/HETG study of the contact binary, 44i~Boo (G0V~\\&~G0V), spanning over 2.5\\,${P}_{\\rm rot}$. Flaring emission is filtered out and velocity shifts in  line centroids of emission line profiles are measured as a function of rotation phase. Rotational modulation in the quiescent corona is detected, indicating  a compact emitting area  near the pole of the primary. We  want to follow up on this latter study in order to learn whether  the picture of compact stellar coronae is applicable to other active cool stars or whether it is exclusive to contact binaries.   \\subsection{AB Doradus}  AB Doradus (K0V, $m_V$=6.9, ${P}_{\\rm rot}$=0.51\\,d, v$_e \\sin i$=90\\,km\\,s$^{-1}$) is an active, rapidly rotating single star that has recently arrived onto the main sequence. It is an ideal candidate for spectroscopic mapping techniques such as Doppler imaging as its 12.4\\,hour rotation period means that a large fraction of its surface can be tracked in one night, and its rotationally broadened spectral line profiles enable detailed surface maps to be reconstructed ($\\sim 3$\\deg\\ longitude resolution at the equator). Spot maps of AB Dor's surface typically show a dark spot covering its pole (polar cap) with smaller spots co-existing at lower latitudes  (e.g. Donati et al. 1999). AB Dor is also bright in X-ray wavelengths, $L_{x}/L_{bol}~10^{-3}$,  and shows strong signs of coronal variability. ROSAT X-ray light curves of AB Dor indicate rotational modulation at the 5-13\\% level (Kuerster et al. 1997). ",
        "conclusions": "Fig.~3 shows the X-ray light curve for AB Dor. We find that the light curve varies by 30\\%. Initial periodogram analyses indicate that a fraction of this is indeed due to rotational modulation. We will use the method used by K\\\"urster et al. (1997) to filter this X-ray light curve in order to distinguish the rotationally modulated component. This rotationally modulated component of the light curve will be compared with the predicted light curve in Fig.~1. Our 3-D coronal model of AB Dor can also be refined, we can compensate for missing magnetic field information (e.g. in the polar cap). However, only a unipolar spot should significantly modify the coronal field distribution (McIvor et al. 2003). We will also evaluate the effect on the predicted X-ray lightcurve if the global magnetic field is non-potential using the technique described by Hussain et al. (2002).  Measurements of the centroids of the best resolved spectral lines (the summed profile derived using the above method and  O\\,{\\sc viii}\\,18\\,\\AA\\ line) show evidence of rotational modulation that repeats from one rotation cycle to the next. Even though both the O\\,{\\sc viii}\\,18\\,\\AA\\ and the summed line profile show the same trend, the uncertainty associated with our velocity shift measurements makes this result somewhat inconclusive. This pattern of rotational modulation would indicate that the bulk of quiescent X-ray emission in AB Dor (at this epoch) originated in localised  regions, most likely at high latitudes ($\\sim$60$^{\\circ}$). Emission originating at the  pole would show no rotational modulation. We will add more lines from the longer wavelengths  ($>88$\\,\\AA) to reduce the error bars on this measurement."
    },
    "0403/astro-ph0403486_arXiv.txt": {
        "abstract": "We present a spectroscopic analysis of nearly 8000 late-type dwarfs in the Sloan Digital Sky Survey.  Using the H$\\alpha$ emission line as an activity indicator, we investigate the fraction of active stars as a function of spectral type and find a peak near type M8, confirming previous results.  In contrast to past findings, we find that not all M7-M8 stars are active.  We show that this may be a selection effect of the distance distributions of previous samples, as the active stars appear to be concentrated near the Galactic Plane. We also examine the activity strength (ratio of the luminosity emitted in H$\\alpha$ to the bolometric luminosity) for each star, and find that the mean activity strength is constant over the range M0-M5 and declines at later types.  The decline begins at a slightly earlier spectral type than previously found.  We explore the effect that activity has on the broadband photometric colors and find no significant differences between active and inactive stars. We also carry out a search for subdwarfs using spectroscopic metallicity indicators, and find 60 subdwarf candidates. Several of these candidates are near the extreme subdwarf boundary.  The spectroscopic subdwarf candidates are redder by $\\sim 0.2$ magnitudes in $g-r$ compared to disk dwarfs at the same $r-i$ color. ",
        "introduction": "Low mass stars with late spectral types (M,L) are the majority constituent of the Galaxy by number.  Their main-sequence lifetimes are much greater than the current age of the Universe and they therefore serve as useful probes of Galactic star formation history in the local solar neighborhood (Gizis, Reid \\& Hawley 2002). They also encompass many important regions of stellar parameter space, including the onset of complete convection in the stellar interior, the onset of significant electron degeneracy in the core, and the formation of dust and subsequent depletion of metals onto dust grains in the stellar atmosphere.  Of particular interest is the fact that many late-type stars have strong surface magnetic fields (Johns-Krull \\& Valenti 1996) that heat the outer atmosphere above the photosphere, and lead to observable emission from the chromosphere (e.g. Ca II and H Balmer series lines), the transition region (e.g resonance lines of abundant ions such as C IV), and the corona (e.g. thermal soft X-rays).  The physics that controls the production of magnetic fields in low mass stars is not well understood, as the lack of a radiative-convective boundary layer precludes storing large scale fields as in a solar-type dynamo.  However, recent work on turbulent dynamo mechanisms (Bercik et al. 2004) may soon provide new insight into the formation and properties of magnetic fields in low mass stars.  The present study seeks to place strong empirical constraints on the magnetic activity in these stars, by measuring optical H$\\alpha$ emission in a sample of nearly 8000 low mass dwarfs. Our results will aid in the interpretation of the models and provide a connection to the physical processes occurring in the atmospheres of the stars.  The H$\\alpha$ emission line is produced by collisional excitation in the relatively dense chromospheres of these low mass, high gravity stars. Located in the red region of the optical spectrum, it is the strongest and best-studied indicator of magnetic activity in late type stars (in contrast to solar type stars which are usually studied in the Ca II H and K resonance lines).  Previous results indicate that the fraction of M dwarfs with H$\\alpha$ emission increases from early to mid M spectral types (Joy \\& Abt 1974; Hawley, Gizis \\& Reid 1996), reaches a peak near type M7 where essentially all stars are active, and declines toward later types (Gizis et al. 2000a).  The increasing fraction in the early to mid M types may be an age effect such that the activity lasts longer in the mid M types (Hawley, Reid \\& Tourtellot 2000a; Gizis et al. 2002), while the decline toward late M types may reflect difficulty in producing and/or maintaining surface magnetic fields due to the physics of turbulent dynamos and/or the increasingly neutral atmosphere (Hawley, Reid \\& Gizis 2000b; Fleming, Giampapa \\& Schmitt 2000; Mohanty et al. 2002).  It does not appear that the atmosphere changes character such that H$\\alpha$ is no longer produced when magnetic activity is present in the later type stars. Indeed the ratio of H$\\alpha$ to transition region (C IV) and soft X-ray emission appears to be similar throughout the M dwarf spectral sequence (Hawley \\& Johns-Krull 2003).  Previous work also shows that the activity strength, measured by the ratio of the luminosity in H$\\alpha$ to the bolometric luminosity (L$_{\\rm{H}_\\alpha}$/L$_{bol}$), is nearly constant (with large scatter) through the M0-M6 range (Hawley et al. 1996) and declines at later types (Burgasser et al. 2002, Cruz \\& Reid 2002). However, the sample of stars on which these conclusions are based numbers less than 100 for types M7 and later, due to the difficulty of obtaining spectra for these faint objects.  Our sample in this crucial spectral type range is an order of magnitude larger, comprising more than 1000 stars with types between M7 and L0, and nearly 8000 stars between M0 and L0. This uniformly acquired and reduced sample from the Sloan Digital Sky Survey (SDSS), with well-characterized uncertainties, now allows us to obtain statistically significant magnetic activity results for the entire M dwarf spectral type range.  In addition, we compare the photometric colors of the active and inactive stars to investigate whether the active stars are typically bluer at the same spectral type, an effect attributed to an increase in plage areas on active stars by Amado \\& Byrne (1997).  Alternatively, such a correlation might indicate that micro-flaring activity plays a significant role in heating the outer atmospheres of active M dwarfs (cf. G\\\"{u}del et al. 2003).  Our SDSS spectral sample is also well-suited for investigation of the incidence of low mass subdwarfs in the local neighborhood.  Gizis (1997) showed that M subdwarfs could be identified by comparing the metallicity and gravity sensitive molecular bands of CaH and TiO.  Using his technique, we identify dozens of new M subdwarf and extreme subdwarf candidates, and examine their colors in the SDSS $ugriz$ filters. ",
        "conclusions": "Using a sample of more than 22000 candidates from the SDSS spectroscopic database, 7840 late type stars were identified and used to investigate the magnetic activity in M and early L dwarfs.  This sample is much larger than all previous samples used for this purpose.  Our results show that: \\begin{enumerate}  \\item{The fraction of active stars rises monotonically from spectral type M0 to M8 and then declines monotonically to the latest types we measured (L3-L4).}   \\item{Only $\\sim 70\\%$ of the M7 stars in our sample are active, which differs from previous studies showing 100\\% activity at spectral type M7. We suggest that the answer to this discrepancy may lie in the distance distributions of the samples used.  Our results indicate the possibility that stars near the Galactic Plane are more likely to be active, as expected if activity depends on age.}   \\item{The activity strength, as measured by ${\\rm L}_{\\rm{H}\\alpha}/{\\rm L}_{bol}$, is approximately constant from types M0-M5, and declines at later types. The width of the distribution is relatively narrow at early types but broadens considerably at type M6 and later.}  \\item{The average $(u-g)$ color is slightly bluer for active stars compared to inactive ones, while the average $(g-r)$ color is slightly redder. However, these differences are not statistically significant.}  \\item{60 new subdwarf candidates are identified spectroscopically by comparing the strengths of the TiO and CaH molecular bands.  Several of the candidates lie near the extreme subdwarf boundary and may be interesting targets for future observations.}  \\item{The subdwarf candidates have similar colors to the disk dwarfs except in the $(g-r)$ color, where the median subdwarf $(g-r)$ is almost 0.2 magnitudes redder than disk dwarfs at the same $(r-i)$ color.  This effect allows subdwarfs to be chosen photometrically from the SDSS database.}  \\end{enumerate}"
    },
    "0403/astro-ph0403165_arXiv.txt": {
        "abstract": "{{ We present a combined radio, optical, and X-ray study of the nearby LINER galaxy NGC\\,1052. Data from a short (2.3\\,ksec) {\\it CHANDRA} observation of NGC\\,1052 reveal the presence of various jet-related X-ray emitting regions, a bright compact core and unresolved knots in the jet structure as well as an extended emitting region inside the galaxy well aligned with the radio synchrotron jet-emission. The spectrum of the extended X-ray emission can best be fitted with a thermal model with $kT = (0.4-0.5)$\\,keV, while the compact core exhibits a very flat spectrum, best approximated by an absorbed power-law with $N_{\\rm H} = (0.6-0.8) \\times 10^{22}\\,{\\rm cm^{-2}}$. We compare the radio structure to an optical ``structure map'' from a {\\it Hubble Space Telescope} ({\\it HST}) observation and find a good positional correlation between the radio jet and the optical emission cone. Bright, compact knots in the jet structure are visible in all three frequency bands whose spectrum is inconsistent with synchrotron emission.} ",
        "introduction": "NGC\\,1052 is a nearby\\footnote{$D = $ 22.6\\,Mpc (assuming $z =$ 0.0049 (Knapp et al. \\cite{Kna78}) and $H_0 = 65$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$). At this distance 1\\,arcsec corresponds to $\\sim$110\\,pc.} elliptical galaxy which harbors a low-luminosity active galactic nucleus (LLAGN) in its very center (L$_{1-100{\\rm GHz}}=4.4 \\times 10^{40}$\\,erg\\,s$^{-1}$; Wrobel \\cite{Wro84}). { It hosts} a two-sided radio jet emanating from the nucleus and reaching out to { kiloparsec-scales} which is, however, still fully enclosed within the stellar body of the optical galaxy. In the optical, the spectrum of NGC\\,1052 is characterized by strong forbidden lines from low-ionization states which has made NGC\\,1052 the prototypical LINER (low-ionization nuclear emission line region; Heckman \\cite{Hec80}) galaxy. As for LINERs in general, it has long been argued whether these low-ionization lines in NGC\\,1052 are excited by a central photo-ionizing source (e.g., Gabel et al. \\cite{Gab00}) or if shock heating is the dominant mechanism (e.g., Sugai \\& Malkan \\cite{Sug00}). While there is overwhelming evidence for the presence of an active galactic nucleus (AGN) in NGC\\,1052, the role of shocks in this galaxy is still unclear. The improved angular resolution in X-rays offered by {\\it CHANDRA} makes it possible for the first time to image the distribution of X-ray emission on the same scales as accomplished by connected radio interferometers, e.g., MERLIN (Multi-Element Radio-Linked Interferometer Network). Disentangling the contributions of compact (i.e., $<$1\\,arcsec) nuclear and extended (i.e., $>$1\\,arcsec) X-ray emitting regions to the total amount of X-ray emission of NGC\\,1052 can serve as an important tool to study the interaction between the radio jet plasma and the ambient interstellar medium.  NGC\\,1052 has been observed by all major X-ray missions of the pre--{\\it CHANDRA} era, like {\\it Einstein} (Mc Dowell \\cite{McD94}), {\\it ASCA} and {\\it ROSAT} (Weaver et al. \\cite{Wea99}), and {\\it Beppo Sax} (Guainazzi \\& Antonelli \\cite{Gua99}). For these X-ray missions, NGC\\,1052 appeared as a point-like X-ray source. { The X-ray spectrum of NGC\\,1052 is extremely flat, a finding that led to the proposal of an advection dominated accretion flow (ADAF) as the origin of the observed X-ray emission (Guainazzi et al. \\cite{Gua00}).} To model the AGN X-ray spectrum above $E \\simeq $ 2\\,keV absorbing column densities in excess of $10^{23}$\\,cm$^{-2}$ have been discussed, supporting the idea of a high density obscuring torus, comparable to column densities found in other AGN (e.g., Malizia et al. \\cite{mali97}, Risaliti et al. \\cite{risa02}). Independent evidence for the existence of an obscuring torus at the center of NGC\\,1052 is obtained from Very Long Baseline Interferometry (VLBI) observations in the radio regime: on { parsec-scales} NGC\\,1052 exhibits a twin jet structure with a prominent emission gap between both jets (see e.g., Kadler et al. \\cite{Kad03b}). The inner part of the western jet shows a strongly inverted radio spectrum, which was first discovered by Kellermann et al. (\\cite{Kel99}) (see also Kameno et al. \\cite{Kam01}). The cm-wavelength spectral index in this central region is larger than 2.5, exceeding the theoretical limit for synchrotron self-absorption.  Combined studies of the core region of NGC\\,1052 in the radio and X-ray regime are of essential importance { for constraining the physical properties of the { parsec-scale} radio jet and the obscuring torus as well as to determine the nature of the nuclear X-ray emission.} In this paper we present a combined radio, optical, and X-ray study of the jet-related emission in NGC\\,1052 on arcsecond scales. In particular, we focus on the soft X-ray excess in the source-spectrum below $E = $ 2\\,keV. This soft component was identified first by Weaver et al. (\\cite{Wea99}) based on {\\it ROSAT} PSPC data. {\\it CHANDRAs} superior angular resolution makes it possible to present evidence that this soft excess emission is associated with the well known radio jet.  In Sect.~\\ref{obs} we present the {\\it CHANDRA}, MERLIN, and {\\it HST} data as well as their reduction. In Sect.~\\ref{multi-structure} we discuss the arcsecond-scale morphology of NGC\\,1052 in the radio, optical, and X-ray regime and the correlations between the different wave bands. In Sect.~\\ref{spectroscopy} we derive models for the nuclear and extended X-ray emission and Sect.~\\ref{summ} summarizes our conclusions. ",
        "conclusions": "\\label{summ} The {\\it CHANDRA} data provide for the first time direct evidence for { jet-associated X-ray emission in NGC\\,1052.} The diffuse, extended X-ray emission can be best approximated with a thermal plasma model with { $kT \\sim 0.4-0.5$\\,keV}. This temperature is consistent with the thermal component found earlier by Weaver et al. (1999) using {\\it ASCA} and {\\it ROSAT} data. Its absorbed flux is only { $\\sim 3$\\%} of the nuclear X-ray emission { but the intrinsic (absorption corrected) extended emission might contribute up to 14\\% to the total unabsorbed X-ray flux of NGC\\,1052.} Because of the { considerable} pile-up degradation of the {\\it CHANDRA} data, no firm conclusions on the photon index of the nucleus spectrum can be deduced. The derived column density of hydrogen towards the compact X-ray core (depending on the applied model) { of $0.5-0.8 \\times 10^{22}$\\,cm$^{-2}$ is in good agreement with the absorbing column density of ionized material towards the VLBI-jet derived by Kadler et al. (\\cite{Kad02}) and Kadler et al. (in prep.). This suggests that the nuclear X-ray emission of NGC\\,1052 might be produced on the same scales as the parsec-scale structures imaged by VLBI at high frequencies.}  The detection of a diffuse region of X-ray emitting gas with a thermal spectrum and the same extent as the kiloparsec-scale radio jet suggests that jet-triggered shocks might play an important role in NGC\\,1052. In such a model the kinetic power of the radio jet is partially converted into X-ray emission. The optical morphology in the H$\\alpha$ filter substantiates this picture as was noted earlier by Allen et al. (\\cite{All99}). The alignment of the radio jet and the optical emission cone visible in Fig.~\\ref{fig:structure_map} implies that the ionization cone might be drilled out by the radio jets,  resulting in a predominantly shock-excited, conical narrow-line region (see e.g., Dopita \\cite{Dop02}). Shocks might occur also on larger scales giving rise to the soft thermal X-ray emission associated with the radio jet/lobe structure in NGC\\,1052. A rough estimate (see Kadler et al. \\cite{Kad03a}) shows that the soft thermal X-ray spectrum associated with the radio jet of NGC\\,1052 can be explained in terms of the kinetic jet power being partially converted into X-ray emission originating in shocks driven into the ambient medium. (A more detailed model of the relation between jet-driven shock-activity and the spectral shape of the extended X-ray emission in NGC\\,1052 will be discussed in a forthcoming paper.) The comparison of the large-scale distribution of radio emission in NGC\\,1052 between two epochs separated by $\\sim$ 15 years indeed shows activity on { kiloparsec-scales}. This substantiates the idea that shocks in the interstellar medium form at the working surfaces of active regions (hotspots and knots). Moreover, recent numerical simulations (e.g., Zanni et al. \\cite{Zan03}) show that jets in radio galaxies can inflate over-pressured cocoons that drive shocks into the ambient gas resulting in morphologies (in the case of weak shocks) very similar to what is observed in NGC\\,1052: a cavity of hot X-ray emitting gas in conjunction with a local deficit of X-ray emission around the hotspots. A deeper {\\it CHANDRA} observation with an improved  photon statistic { compared to} the observation discussed here { would} provide both a higher sensitivity to the weak diffuse emission and a higher resolution. Additionally, the full resolution of {\\it CHANDRA} of $\\sim 0.5$\\,arcsec would allow one to study in more detail the connection between the knots in the diffuse X-ray emission and the optical emission knots."
    },
    "0403/astro-ph0403353_arXiv.txt": {
        "abstract": "The large-angle (low-$\\ell$) correlations of the Cosmic Microwave Background exhibit several statistically significant anomalies compared to the standard inflationary cosmology. We show that the quadrupole plane and the three octopole planes are far more aligned than previously thought ($99.9$\\% C.L.). Three of these planes are orthogonal to the ecliptic at $99.1$\\% C.L., and the normals to these planes are aligned at $99.6$\\% C.L.~with the direction of the cosmological dipole and with the equinoxes. The remaining octopole plane is orthogonal to the supergalactic plane at $99.6$\\% C.L. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403679_arXiv.txt": {
        "abstract": "s{We are undertaking the first systematic study of the prevalence of AGN activity in a large sample of high redshift galaxy clusters. Local clusters contain mainly red elliptical galaxies, and have little or no luminous AGN activity. However, recent studies of some moderate to high redshift clusters have revealed significant numbers of luminous AGN within the cluster. This effect may parallel the Butcher-Oemler effect --  the increase in the fraction of blue galaxies in distant clusters compared to local clusters. Our aim is to verify and quantify recent evidence that AGN activity in dense environments increases with redshift, and to evaluate the significance of this effect. As cluster AGN are far less prevalent than field sources, a large sample of over 120 cluster fields at $z > 0.1$ has been selected from the Chandra archives, and is being analysed for excess point sources. The size of the excess, the radial distribution and flux of the sources and the dependence of these on cluster redshift and luminosity will reveal important information about the triggering and fueling of AGN. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/nucl-th0403066_arXiv.txt": {
        "abstract": "We use Landau's theory of a normal Fermi liquid to derive expressions for the static response of a system with a general tensor interaction that conserves the total spin and the total angular momentum of the quasiparticle-quasihole pair.  The magnetic susceptibility is calculated in detail, with the inclusion of the center of mass tensor and cross vector terms in addition to the exchange tensor one.  We also introduce a new parametrization of the tensor Landau parameters which significantly reduces the importance of high angular harmonic contributions.  For nuclear matter and neutron matter we find that the two most important effects of the tensor interaction are to give a contribution from multipair states and to renormalize the magnetic moments. Response to a weak probe may be calculated using similar methods, replacing the magnetic moments with the matrix elements of the weak charges. ",
        "introduction": "For systems with central interactions, Landau's theory of normal Fermi liquids provides an economical way of characterizing many low-temperature properties.  The theory was applied to atomic nuclei by Migdal and collaborators, and in that work it was generally assumed that the non-central contributions to the effective nucleon-nucleon interaction were small \\cite{migdal}.  The generalization of Landau theory to include effects of the tensor force was made by D\\c{a}browski and Haensel \\cite{dabrowski2,dabrowski1,haensel2}.  Subsequently, estimates of tensor contributions to the effective interaction were made in Refs.\\ \\cite{baeckman,schwenk,schwenk2,dickhoff}.  The stimulus for the present work arose in the context of astrophysics. In the physics of collapse and the subsequent evolution of a neutron star, the properties of neutrinos in dense matter are a key ingredient \\cite{janka}. As demonstrated in Refs.\\ \\cite{sawyer,iwamoto,raffelt1,raffelt2,raffelt3,burrows,prakash,prakashreddy}, neutrino scattering and absorption rates are sensitive to nucleon-nucleon interactions, especially their spin-dependent parts. Direct calculation of effective interactions is difficult, but for systems with only central interactions, information about the effective interaction at long wavelengths may be obtained directly from a knowledge of static properties of the matter. For example, the spin susceptibility of a single-component Fermi liquid with two spin states is given by \\beq \\chi=\\frac{\\mu_0^2 N(0)}{1+G_0}, \\label{Landauchi} \\eeq where $\\mu_0$ is the magnetic moment of a particle in free space, $N(0)=m^* p_{\\rm F}/\\pi^2\\hbar^3$ is the density of states per unit volume at the Fermi surface, $m^*$ is the quasiparticle effective mass, $p_{\\rm F}$ is the Fermi momentum and $G_0$ is the Landau parameter describing the isotropic part of the spin dependent contribution to the quasiparticle interaction.  Since calculations of the magnetic susceptibility of neutron matter exist \\cite{Fantoni}, it is relevant to ask to what extent it is possible to deduce properties of the effective interaction from such data. A related question is the extent to which tensor contributions to quasiparticle energies and interactions alter neutrino scattering rates, which were previously calculated neglecting tensor effects in Ref. \\cite{iwamoto}.   As shown in earlier work \\cite{olsson}, the tensor force influences the static response of a system in a number of different ways. One is that the magnetic moment of a quasiparticle is different from its value for a particle in vacuum. In particular, the magnetic moment is not a scalar, as it is for systems with central forces only. As long ago as 1951, Miyazawa \\cite{miyazawa} calculated explicitly how the magnetic moment of a nucleon would be modified by the tensor interaction due to one pion exchange and more recent discussions may be found in e.g. the review \\cite{arima} and Ref. \\cite{riska}. A second effect is that the quasiparticle interaction contains explicit tensor contributions. Following Refs. \\cite{dabrowski1,dabrowski2,haensel2}, it has generally been assumed that these have an exchange-tensor structure similar to that of the one-pion exchange interaction. However, Schwenk and Friman \\cite{schwenk2} have pointed out recently that the one-pion exchange interaction, when acting in second order, can give rise to contributions to the effective interaction which have a different structure. In their paper they evaluated these induced interaction contributions to the quasiparticle interaction and found that the exchange-tensor term in the quasiparticle interaction is much reduced, and that terms of a different structure can be of a comparable magnitude to the exchange-tensor ones. A third effect of the tensor interaction is that there are multipair contributions to the magnetic susceptibility.    In this paper we begin by deriving a general expression for the magnetic susceptibility in Sec.\\ \\ref{Sec:response}, taking into account tensor contributions to the magnetic moment, and tensor contributions to the effective interaction which are completely general for an interaction which conserves the total angular momentum and the total spin of the quasiparticle-quasihole pair. This represents a generalization of the earlier calculation of the magnetic susceptibility by Haensel and Jerzak \\cite{haensel}, who took into account the effect of the exchange-tensor contribution to the quasiparticle interaction. Another issue that we address is how to parametrize tensor contributions to the quasiparticle interaction (Sec.\\ \\ref{Sec:tensor}). The scheme usually employed in the past suffers from the disadvantage that it is generally necessary to take into account high angular harmonic contributions to the interaction. We present an alternative parametrization for which higher harmonic terms play little role. In Sec.\\ \\ref{Sec:ME} we evaluate the matrix elements needed for the calculation of the susceptibility and estimate the magnitudes of the different contributions. In Sec.\\ \\ref{Sec:SNM} we extend the result to symmetric nuclear matter and to responses with other spin and isospin properties, for example that to weak interactions.   Sec.\\ \\ref{Sec:conclusions} contains concluding remarks. ",
        "conclusions": "\\label{Sec:conclusions} In this paper we have derived an expression for the magnetic susceptibility of a Fermi liquid for a general interaction which conserves the total angular momentum and the total spin of the quasiparticle-quasihole pair. Apart from the contribution from single quasiparticle-quasihole pairs, which may be calculated using Landau theory, the susceptibility also contains a contributions from excitation of multipair states.  In addition to the exchange tensor term usually included, we have taken into account the new tensor terms found in Ref. \\cite{schwenk2}.  We have also introduced an alternative parametrization of the exchange tensor interaction which has the advantage of reducing the importance of high angular harmonic contributions.  For neutron matter we find that the most important effects of the tensor force are to renormalize the magnetic moment and to give a contribution from multipair states. For symmetric nuclear matter one can draw the same conclusion, the bound on the multipair contribution is, however, much smaller, but it is not clear how close the actual contribution to the susceptibility, from multipairs, is to this lower bound.  The formalism may easily be extended to calculate the response to weak probes: the only difference being that the magnetic moments must be replaced by the corresponding matrix elements of the weak charges."
    },
    "0403/astro-ph0403308_arXiv.txt": {
        "abstract": "We build template maps for the polarized Galactic--synchrotron emission on large angular scales (FWHM~=~7$^\\circ$), in the 20-90~GHz microwave range, by using WMAP data. The method, presented in a recent work, requires a synchrotron total intensity survey and the {\\it polarization horizon} to model the polarized intensity and a starlight polarization map to model polarization angles. The basic template is obtained directly at 23~GHz with about 94\\% sky--coverage by using the synchrotron map released by the WMAP team. Extrapolations to 32, 60 and 90~GHz are performed by computing a synchrotron spectral index map, which strongly reduces previous uncertainties in passing from low (1.4~GHz) to microwave frequencies. Differing from low frequency data, none of our templates presents relevant structures out of the Galactic Plane. Our map at 90~GHz suggests that the synchrotron emission at high Galactic latitudes is low enough to allow a robust detection of the $E$--mode component of the cosmological signal on large--scale, even in models with low--reionization ($\\tau = 0.05$). Detection of the weaker $B$--mode on the largest scales ($\\ell < 10$) might be jeopardized unless the value $\\tau = 0.17$ found by WMAP is confirmed, and  $T/S > 0.1$. For lower levels of the gravitational--wave background the $B$--mode seems to be accessible only at the $\\ell \\sim 100$ peak and in selected low--synchrotron emission areas. ",
        "introduction": "\\label{intro}  Recent results by  DASI (Kovac et al. 2002) and WMAP (Bennett et al 2003a) have opened a new window in studying both the Cosmic Microwave Background (CMB) and the Galaxy.  The polarized signal detected with the DASI experiment, as well as the WMAP measurement of the temperature--polarization cross--spectrum $C^{TE}_\\ell$, emphasized the importance of studying CMB Polarization (CMBP). In particular, the discovery of a relevant signal on large angular scales performed by WMAP has provided evidence of an unexpectedly early reionization era (Kogut et al. 2003).  Furthermore, WMAP provided the first all--sky total--intensity maps of the microwave Galactic emission, allowing an insight into synchrotron, dust and free-free components of the Galaxy (Bennett et al. 2003b).  A sound measurement of CMBP requires good knowledge of polarized foregrounds (both Galactic and extragalactic), which are potentially more dangerous than in total intensity. This is even more true for the $B$--mode, retaining information on the gravitational--wave background (Kamionkowski \\& Kosowski, 1998), whose emission level is expected to be orders of magnitude below that of the $E$-mode.  Among all foreground components, the Galactic synchrotron radiation is expected to be the most relevant up to 100~GHz. The lack of data in both the CMB frequency range and at high Galactic latitudes makes the building of templates necessary (Kogut \\& Hinshaw 2000, Giardino et al. 2002, Bernardi et al. 2003a, hereafter B03). These can substantially help in studying how to extract CMB maps from contaminated signals (e.g. see Bennett et al. 2003b).  In B03 we presented a method modelling the Galactic polarized--synchrotron emission by using the radio--continuum total intensity surveys available at low frequency ($< 1.4$~GHz) to model the polarized intensity, and starlight polarization optical data as a template for polarization angles.  Template maps obtained with this method are virtually free of Faraday rotation. However, when extrapolating low frequency data (the only available before the latest WMAP release) to the cosmological window, uncertainties arise on the spectral index to be used. In B03 we used the mean spectral index derived by Platania et al. (1998) up to 19~GHz, since no information were available at higher frequencies. In addition, variations across the sky were too poorly known to be taken into account.  The WMAP release of total--intensity synchrotron maps in the cosmological window gives us the possibility to significatively improve our templates. Actually, the application of our method directly at frequencies of interest for CMBP fully avoids extrapolation uncertainties.  Our templates are built with FWHM~$=7^{\\circ}$ at the three frequencies of 23, 32 and 90~GHz interesting for the SPOrt experiment (Cortiglioni et al. 2004), one of whose aims is studying the CMBP foregrounds. We also generate a template at 60~GHz, this frequency being of interest for other CMB experiments (WMAP, Planck). All of the templates cover large fractions of the sky, namely 94\\% at 23~GHz and 92\\% at 32, 60 and 90~GHz. In addition, we calculate polarized--syncrotron angular power spectra both for the full-sky templates and for the high Galactic latitude emission, discussing possible effects on the measurement of CMBP $E$ and $B$--mode spectra.  The paper is organized as follows: in Section~\\ref{model} we apply our method to WMAP  data at 23~GHz and provide a Galactic polarized synchrotron template at the same frequency; in Section~\\ref{extr} we extrapolate our results up to 90~GHz; in Section~\\ref{tot_aps} we describe the angular power spectra of our full--sky templates at the different frequencies, in Section~\\ref{for_aps} we compute the angular power spectra for the high Galactic latitude emission and compare synchrotron to CMBP and, finally, conclusions are presented in Section~\\ref{conc}. ",
        "conclusions": "\\label{conc} In this work we generate template maps of polarized Galactic synchrotron emission at 23, 32, 60 and 90~GHz using WMAP data.  We apply the method developed in B03 to the 23~GHz total intensity synchrotron map released by the WMAP team. This greatly improves the B03 templates avoiding the uncertainties in the extrapolation from low frequency. As for the previous work, the angular resolution is limited to FWHM=$7^\\circ$ mainly due to the starlight polarization data sampling.  The basic template (23--GHz--W) is constructed at 23~GHz applying the method directly to the WMAP data and it covers almost all the sky (94\\%). Differing from the 1.4~GHz B03 template where the Northern Galactic Spur represents an important feature, 23--GHz--W does not present relevant structures out of the Galactic Plane. This replicates what already occurs in total intensity and leaves most part of the sky with low synchrotron emission and useful for CMBP investigations.  To achieve the templates at the other frequencies a spectral index map is required to extrapolate the 23--GHz--W result. Differing from Bennett et al. (2003b), which derived a spectral index map from 0.408 to 23~GHz with no free--free subtraction in the 0.408~GHz map, we computed the frequency behaviour using {\\it pure} synchrotron maps at 1.4 and 23~GHz. The 23~GHz one is that released by the WMAP team. The 1.4~GHz map has been achieved from low frequency surveys by applying a Dodelson-like component separation. We use the Haslam et al. 0.408~GHz and the Reich 1.4~GHz surveys for the Northern sky and the Haslam et al. and the Jonas et al. 2.3~GHz ones for the Southern Hemisphere. The results is the first almost all-sky spectral index map of the Galactic synchrotron emission in the 1.4--23~GHz range. Furthermore, we add an extra 0.5 to the indeces to account for the steepening found by Bennett et al. (2003b) in the 23--41~GHz range and this spectral index map is used to compute template maps at 32, 60 and 90~GHz covering 92\\% of the sky.  The spectral index map provides steeper indexes at high Galactic latitudes so that the resulting templates have the Galactic Plane emission more and more dominant from 23~GHz up to 90~GHz. This difference also explains the disappearing of the Northern Galactic Spur in our templates. Unclear, instead, it is the reason for such a flat frequency spectrum found in the Galactic Plane, it being either intrinsic to the synchrotron emission (for instance due to a young population of relativistic electrons) or due to a residual thermal component.  Beyond the lack of relevant structures at high Galactic latitudes, the differences among the templates of this work and the B03 one also concern the APS. APS of our full--sky templates are dominated by the emission from the Galactic Plane and flatter slopes, namely in the 1.4--1.6 range for $C^E_\\ell$, 0.5--0.8 for $C^B_\\ell$ and 1.2--1.5 for $C^P_\\ell$. This general flattening is still not clear.  A further unclear point is that the synchrotron APS slopes (in particular $C^B_\\ell$ and $C^P_\\ell$) seem to flatten with increasing frequency. If confirmed, these two aspects would confirm that low frequency APS cannot be straightforwardly extrapolated from low to microwave frequencies.  The situation is less complex at high Galactic latitudes. In fact, limiting the analysis out of the Galactic Plane, a steepening of the slopes occurs, which become more similar to the B03 ones, and no significative dependency from the frequency is found.  Interesting conclusions are drawn by the comparison of the high Galactic latitude portion of our templates to CMBP as from the {\\it concordance model} of the WMAP first--year data. Indeed, the CMBP $E$--mode spectrum is about 2 orders of magnitude above the synchrotron signal in our template at 90~GHz. Synchrotron is thus unlikely to contaminate the cosmological signal. Furthermore, even a less re-ionized Universe ($\\tau = 0.05$) should be dominant over the Galactic emission.  The situation looks different for the CMBP $B$--mode. Our template predicts that this cosmological signal might be accessible at largest scales ($\\ell < 10$) only for a Tensor to Scalar fluctuation power ratio $T/S > 0.1$. To investigate models with lower values of $T/S$ one should restrict to selected low synchrotron emission areas large enough to detect the $\\ell \\sim 100$ peak. An example is the sky patch observed by the BOOMERanG experiment where models with $T/S > 0.01$ seem to be accessible (Bernardi et al. 2003b).  Our templates obviously rely on the assumption that WMAP's dust-correlated signal at 22~GHz is genuine synchrotron, as supported by Bennet et al. (2003b). However several authors (de Oliveira--Costa et al. 2003, Banday et al. 2003, Lagache 2003) recently support the opposite view, that WMAP's data may be more consistent with anomalous dust emission. For instance, while the  spinning-dust model originally proposed by Draine \\& Lazarian (1998) may meet some difficulty, emission by very small grains is supported by Lagache (2003).  Considering that dust polarization level is likely less than 5\\% rather than the 15-30\\% of the synchrotron, a measurement of polarization at 22-32~GHz could help to decide on this issue: should the polarization level be found a factor $\\sim 5$ below the predictions of our template, the role of anomalous dust emission would be supported.  Also, in such a case, the contamination of CMBP maps at 90~GHz would be even smaller than considered in our discussion above, and prospects for the detection of a cosmological B mode would be slighly better."
    },
    "0403/astro-ph0403414_arXiv.txt": {
        "abstract": "We report the search for intracluster light in four Abell Type II/III (non-cD) galaxy clusters: Abell 801, 1234, 1553, \\& 1914.  We find on average that these clusters contain $\\sim$ 10\\% of their detected stellar luminosity in a diffuse component.  We show that for two of the clusters the intracluster light closely follows the galaxy distribution, but in the other two cases, there are noticeable differences between the spatial distribution of the galaxies and the intracluster light.  We report the results of a search for intracluster tidal debris in each cluster, and note that Abell~1914 in particular has a number of strong tidal features likely due to its status as a recent cluster merger.  One of the Abell~1914 features appears to be spatially coincident with an extension seen in weak lensing maps, implying the feature traces a large amount of mass.  We compare these results to numerical simulations of hierarchically-formed galaxy clusters, and find good general agreement between the observed and simulated images, although we also find that our observations sample only the brightest features of the intracluster light.  Together, these results suggest that intracluster light can be a valuable tool in determining the evolutionary state of galaxy clusters. ",
        "introduction": "As the most massive gravitationally bound structures in the universe, galaxy clusters stand to teach us much about the hierarchical assembly of matter in the universe. Clusters exhibit a wide variety of structural properties, from massive, X-ray luminous clusters dominated by early type galaxies (i.e., Coma) to irregular, spiral-rich clusters like the Ursa Major cluster, down to poor clusters and loose groups. The fact that many clusters are still obviously in the process of assembly can be seen via many tracers of substructure, such as X-ray isophotes, gravitational lensing maps, and kinematic and spatial substructure in the galaxy populations (see, e.g., ~reviews by Girardi \\& Biviano 2002 and Buote 2002)  A potentially powerful new tracer of the assembly history of clusters is intracluster light, the diffuse starlight which permeates many galaxy clusters. Once simply another curiosity of Zwicky (1951), individual intracluster stars have been clearly detected in several nearby galaxy clusters \\citep{1996arna,ftv1998,durr2002,ipn3}, and through deep imaging this diffuse light has been detected in many more distant clusters \\citep{uson1991b,vg1994,bern1995,gon2000}. From the results to date, it is clear that intracluster light (ICL) is a common component of galaxy clusters and contains between 10\\% and 50\\% of the total stellar luminosity of the cluster, albeit with large uncertainties due to the intrinsic low surface brightness of the component (less than 1\\% of the night sky background).  The study of intracluster light is now entering a new phase, focusing on what can be learned about the evolution of galaxy clusters and their member galaxies using the ICL. Since galaxy clusters form hierarchically, and since the bulk of ICL production is believed to occur due to tidal-stripping from galaxy interactions and from the mean tidal field of the cluster \\citep{rich1983,miller1983,mer1984,gnedin2003}, the properties of the ICL should be intimately tied to the dynamical evolution of clusters (\\eg ~Merritt 1984; Moore \\etal 1996; Dubinski 1998; Mihos 2003; Napolitano \\etal 2003; Willman \\etal 2004). As clusters dynamically evolve, the fractional amount of ICL should increase if the cluster is isolated, and its spatial and dynamical structure should become well-mixed in the cluster potential well.  As additional groups of galaxies enter the cluster, the fractional amount of ICL will briefly decrease, as the newer galaxies are initially unstripped, and then the fraction will increase as the additional galaxies suffer the effects of the cluster environment. The exact details of this overall evolution depend on the mass and accretion history of the cluster, tying the properties of the ICL directly to the cosmic history of cluster assembly.  The exact mechanisms of ICL production have implications for a number of other galaxy cluster studies.  Ultra-compact dwarf galaxies (Drinkwater \\etal 2003), the formation of S0 galaxies (Quilis, Moore, \\& Bower 2000, and references therein) and tidal debris in nearby clusters (Trentham \\& Mobasher 1998; Gregg \\& West 1998; Calcaneo-Roldan \\etal 2000) may all be closely related to the ICL phenomenon.  Since searches for individual intracluster stars in nearby galaxy clusters can be influenced by metallicity effects (Durrell \\etal 2002; Feldmeier \\etal 2004), knowing the dominant progenitor population is critical to avoid underestimating the true fraction of intracluster light. However, observational constraints on the ICL are still extremely poor, due largely to the scarcity of quantitative measurements of the ICL in clusters, especially over a range of cluster properties.  To address these questions, we have recently begun deep imaging of a sample of galaxy clusters to quantify the structure of ICL as a function of galaxy cluster properties (Feldmeier \\etal 2002; hereafter Paper~I).  We have observed two cD-dominated (Bautz-Morgan Type I) galaxy clusters thus far (Abell~1413 and MKW~7); in each we quantify the extended cD envelope and find an excess of luminosity over a pure \\rquart law. However, this extended light is very smooth, and we find relatively little substructure in the ICL of either Abell~1413 or MKW~7.  Again, however, both of these clusters are cD-dominated, and there are a number of reasons to expand our sample beyond cD-dominated galaxy clusters.  First, less than twenty percent of the total number of Abell clusters have a Bautz-Morgan type of I \\citep{leir1977}.  Observations of ICL in Type~I clusters may therefore not be representative of ICL structure in galaxy clusters as a whole.  Second, the relation between intracluster light and cD envelopes is still unclear.  Does the presence of a cD galaxy in a cluster always imply a large amount of intracluster starlight, or are ICL properties less correlated to the precise details of the cluster core?  If intracluster stars are predominantly removed early in the cluster's dynamical history (Merritt 1984), we would expect the properties of ICL to be intimately tied to the process of cD formation.  If instead the bulk of the intracluster stars are removed from their parent galaxies by late tidal-stripping and related ``galaxy harassment'' scenarios operating on low luminosity galaxies (Richstone \\& Malumuth 1983; Moore \\etal 1996), then the ICL properties may be less sensitive to the presence or absence of a cD galaxy in the cluster core.  Finally, given the expected evolution of galaxy clusters under hierarchical structure formation models (\\eg\\ Merritt 1985; Dubinski 1998), it is expected that most rich galaxy clusters will eventually form cD galaxies, though the timescales involved could be quite long (up to thousands of Gyr; Adams \\& Laughlin 1999).  Therefore Bautz-Morgan type II or III clusters may be less dynamically evolved proxies of the more evolved type I clusters, and ICL studies of these clusters may give us insight into to the evolution of galaxy clusters at higher redshift and younger dynamical ages. ",
        "conclusions": "We have surveyed four Abell type II-III (non-cD) galaxy clusters, and find a significant amount of intracluster light in all of them.  The amount of intracluster starlight is intermediate, approximately 10--20\\% of the galaxy light.  This may reflect the dynamical adolescence of these galaxy clusters: as the clusters continue to evolve the mean intracluster fraction should increase.  However, it is likely that cluster mass (or richness) also plays a role: both our own measurements of the rich cluster Abell~1413 (Paper~I), and other studies of richer clusters such as Coma (Bernstein \\etal 1995) find significantly higher intracluster star fractions than the clusters studied in this paper.  Clearly, more observations of clusters over a range of cluster richnesses will be important to separate the two related effects.  We have searched for signs of asymmetries and tidal debris in the intracluster light, and we find multiple examples of these features, though in some clusters we found signs of a regular symmetric structure in the intracluster light and few tidal features.  The number of tidal features in Abell~1914 is clearly unusual, in both number and surface brightness.  It is worth reiterating here that Abell~1914 has a radio halo\\citep{gio1999,kemp2001} and is believed to be undergoing a cluster merger (Jones \\etal 2001).  Therefore, the wealth of high surface brightness tidal features in the ICL in Abell~1914 is likely to be due to the ongoing cluster merger.  We may be observing the cluster in {\\it flagrante delicto}: had we observed the cluster a few crossing times later, the evidence of such a strong merger would be much less apparent.  Comparing our observational results to modern numerical simulations of galaxy clusters, we find good overall qualitative agreement in both the amount and distribution of the intracluster light, although the simulations show that significant ICL substructure may exist well below out current detection levels.  From the results found here, we suggest that intracluster light may act as a dynamical ``clock'' of galaxy clusters, one that is complementary to other studies of galaxy clusters such as X-ray substructure, weak lensing, and galaxy radial velocities.  However, the comparison of ICL properties and models of structure formation and evolution is still in its infancy.  A detailed comparison of the two will require much deeper ICL imaging and measurements of the kinematics of the ICL (though planetary nebulae velocities, for example), as well as through numerical simulations that have larger dynamical range, and better descriptions of galaxy formation processes.  Such studies are now underway."
    },
    "0403/astro-ph0403287_arXiv.txt": {
        "abstract": "We present spectral analysis of the Crab Nebula obtained with the {\\it Chandra} X-ray observatory.  The X-ray spectrum is characterized by a power-law whose index varies across the nebula. The variation can be discussed in terms of the particle injection from the pulsar in two different directions: the equatorial plane containing the torus and the symmetry axis along the jet. In the equatorial plane, spectra within the torus are the hardest, with photon index $\\alpha \\approx 1.9$, and are almost independent of the surface brightness. At the periphery of the torus, the spectrum gradually softens in the outer, lower surface brightness regions, up to $\\alpha \\approx 3.0$. This indicates that synchrotron losses become significant to X-ray emitting particles at the outer boundary of the torus. We discuss the nature of the torus, incorporating information from observations at other wavelengths. Spectral variations are also seen within the southern jet.  The core of the jet is the hardest with $\\alpha \\approx$ 2.0, and the outer sheath surrounding the core becomes softer with $\\alpha$ up to 2.5 at the outermost part.  Based on the similarity between the spectra of the jet core and the torus, we suggest that the electron spectra of the particles injected from the pulsar are also similar in these two different directions.  The brightness ratio between the near and far sides of the torus can be explained by Doppler boosting and relativistic aberration; however, the observed ratio cannot be derived from the standard weakly magnetized pulsar wind model. We also found a site where an optical filament comprised of supernova ejecta is absorbing the soft X-ray emission ($<$ 2 keV). ",
        "introduction": "OBSERVATION AND ANALYSIS}  Owing to the X-ray brilliance of the Crab Nebula, observations with the {\\it Chandra} ACIS instrument are technically challenging. The high count rate generally causes two problems in the data acquisition: telemetry saturation and event pile-up (caused by more than one photon depositing charge in a given pixel during a CCD exposure, thereby compromising our ability to measure the energy of each individual photon). Although telemetry saturation results in loss of observing efficiency, it does not lead to biased estimates of physical parameters (e.g., count rate) because events are discarded in a unit of frames, not event by event. On the other hand, pile-up makes interpretation of spectral results problematic, since it converts two photons into a single ``event'' which has apparent energy equal to the sum of the photon energies.  This results in underestimated count rate and hardens the detected spectrum, as discussed in Weisskopf et al.\\ (2000).  We observed the Crab Nebula eight times (Table~\\ref{tbl:log}), using an instrumental mode of ACIS-S different from that of Weisskopf et al.\\ (2000) in order to reduce pile-up. We adopted a frame time of 0.2 sec instead of the standard frame time of 3.2 seconds, but we did not utilize the grating to reduce count rate. This reduced the counts pixel$^{-1}$ frame$^{-1}$ by a factor of 2.3 with respect to the data of Weisskopf et al.\\ (2000), while improving statistics by increasing the total number of counts by one order of magnitude for a similar effective exposure time.  Since the Crab Nebula is a diffuse source, the event pile-up was suppressed by more than a factor of 2.3.  The observational parameters (e.g., telemetry format of graded mode, no dithering) were the same through the eight observations. A combined image from the last seven observations is shown in Figure~\\ref{fig:index_map}a (the first observation did not include the whole remnant due to an instrument setup error).  This image is an order of magnitude deeper than the image of Weisskopf et al.\\ (2000).  In order to study X-ray spectral variations within the Crab Nebula, we divided the whole image into $2.^{\\!\\!\\prime\\prime}5 \\times 2.^{\\!\\!\\prime\\prime}5$ square regions, from which spectra were extracted.  The eight data sets were combined in the analysis in order to provide each region with enough statistics. We excluded regions dominated by scattered photons (less than 2000 counts per analysis region) and/or by trailing (or ``out of time'') events. Trailing events, which are caused by photons detected during the CCD readout, can be seen not only as the line emanating from the pulsar along the detector readout direction (which is almost east-west for all of our observations), but also as diffuse emission which dominates the faint regions to the east and west of the nebula due to the enormous brilliance of the Crab PWN. The number of the trailing events in a given pixel can be estimated by multiplying the total number of events integrated along the readout direction by the ratio of the charge transfer row period to the frame exposure time (this ratio is $1.89 \\times 10^{-4}$). We exclude regions where more than 20\\% of the total events recorded in that region are due to the trailing events.  We fitted the 2074 spectra with a power-law model in the 0.5--8 keV energy band. We included galactic absorption with a fixed $N_{\\rm H}$ of $3.2 \\times 10^{21} $cm$^{-2}$, which is derived from a fit of the faintest region where pile-up is negligible. The recently discovered quantum efficiency decay of ACIS\\footnote{See http://asc.harvard.edu/cal/Acis/Cal\\_prods/qeDeg/index.html} is taken into account\\footnote{See http://www.astro.psu.edu/users/chartas/xcontdir/xcont.html}. Figure~\\ref{fig:correlation_plot}a shows a correlation plot of apparent photon index versus observed surface brightness.  The statistical errors of the apparent photon indices are quite small. They can be roughly expressed empirically as $0.01 \\times I^{-0.5}$ where $I$ is the observed surface brightness in unit of counts s$^{-1}$ arcsecond$^{-2}$ (90\\% confidence level).  In spite of the reduction in counts per frame, many of the spectra are still distorted by event pile-up. The spectral parameters derived from a fit of the whole nebula spectrum with a power-law plus galactic absorption model are summarized in Table~\\ref{tbl:whole_fit} (uncorrected values), as are the canonical values for comparison. The pile-up results in a lower apparent spectral index and lower apparent intensity than the canonical values. The observed intensity normalization is about 84\\% of that estimated from the canonical values with the instrumental response.  Pile-up has the effect of reducing photon index in regions of increased surface brightness, which is the same general trend seen in the data of Figure~\\ref{fig:correlation_plot}a.  The pile-up effects must be corrected in order to determine the nature of any real spectral variations across the Crab nebula.  In order to estimate how pile-up distorts the incident spectrum, we utilized the spectral simulator tool LYNX (Chartas et al.\\ 2000). LYNX traces the propagation of individual photons through the mirror and ACIS via the raytrace tool MARX (Wise et al.\\ 1997) and the PSU ACIS Monte Carlo CCD simulator (Townsley et al.\\ 2001), respectively. Since LYNX takes into account the possible overlap of the resulting charge clouds within each CCD exposure frame, spectra affected by pile-up can be simulated for both point sources and diffuse sources. We simulated the observed Crab spectrum for several different values of true surface brightness and for true photon indices of 1.5, 2.0, and 2.5. The results are superposed on the data in Figure~\\ref{fig:correlation_plot}a.  The expected reduction of apparent photon index in higher surface brightness regions is evident. However, it is also clear that the variations in observed photon indices greatly exceed those expected from pile-up, and there are certainly real spectral variations across the Crab Nebula, as described by Weisskopf et al.\\ (2000).  We used the LYNX simulation results to correct the data points in Figure~\\ref{fig:correlation_plot}a for the photon index and surface brightness errors induced by the pile-up effects.  For each data point, we interpolated the simulated data points to the observed surface brightness and photon index and calculated the corresponding true surface brightness and photon index. Figure~\\ref{fig:correlation_plot}b presents a correlation plot of the corrected data points. This simulation was intended to give corrections on the data points in Figure~\\ref{fig:correlation_plot}a and was not adjusted to correct the whole nebula spectrum in order to match to the canonical spectral parameters. However, we used the corrected photon indices and surface brightnesses of each data point in Figure~\\ref{fig:correlation_plot}b to calculate a corrected spectrum for the whole nebula and we present the results in Table~\\ref{tbl:whole_fit} (corrected values). The photon index and the normalization factor of the corrected spectrum agree well with the canonical values within the errors, which strongly supports the success of our pile-up corrections.  Hereafter, we use the terms ``photon index'' and ``surface brightness'' to mean these corrected values.  We should note a limitation of our simulation. Our corrections are based on modeling of a uniform diffuse source. We apply them to each $2.^{\\!\\!\\prime\\prime}5 \\times 2.^{\\!\\!\\prime\\prime}5$ square region under the reasonable assumption that the surface brightness within each region is uniform. Therefore, application of this technique to small complicated structures like the inner ring or the relativistically expanding wisps requires more complex modeling which is beyond the scope of this work. The systematic errors of the photon index mainly come from this assumption.  Judging from the results of the simulation shown in Figure~\\ref{fig:correlation_plot}a and typical surface brightness fluctuations of $\\sim 10\\%$ within an analysis region, the systematic errors in photon index are $\\sim$ 0.05 and dominate the statistical errors.  Figure~\\ref{fig:index_map}b shows a map of the photon index. The structure in the torus is more symmetrical about the pulsar in the photon index map than in the broad-band image (Fig.~\\ref{fig:index_map}a).  While Doppler boosting and relativistic aberration brighten the northwestern portion of the torus (Pelling et al.\\ 1987) in the Figure~\\ref{fig:index_map}a image, by contrast there is relatively little variation of the photon index within the torus. The hardest structures in the nebula, with photons indices as low as 1.8, are the inner ring and portions of the torus, including the circular structures seen at each extremity of the torus discovered in the first {\\it Chandra} observation (Weisskopf et al.\\ 2000). However, these regions contain small complex structures and the hardest photon index should be re-examined using data set free from pile-up.  The entire torus is quite hard ($\\alpha \\approx 1.9$) compared with the outer portions of the nebula.  The southern jet is also relatively hard ($\\alpha \\approx 2.0$), whereas the northern counter jet is significantly softer ($\\alpha \\approx 2.25$). The bright region around the counter jet to the northwest of the torus, hereafter called the ``umbrella'' because of its shape, has even softer emission ($\\alpha \\approx 2.5$).  Photon indices as large as 3.0 are found in the outer peripheral portions of the nebula. ",
        "conclusions": "SUMMARY}  We have shown spatial variations of the X-ray spectrum of the Crab Nebula in terms of photon index, at an angular scale of arcseconds. The variations can be viewed in two different directions of the particle injection from the pulsar.  Across the equatorial plane, the spectrum is almost constant to the outer boundary of the torus, with photon index $\\alpha$ $\\approx$ 1.9 regardless of the surface brightness.  It softens significantly, up to $\\alpha$ $\\approx$ 3.0, in the outer, fainter peripheral region.  This seems qualitatively consistent with the previous suggestions that the outer boundary of the torus is interpreted as a synchrotron burn-off boundary where the synchrotron losses become significant to X-ray emitting particles. However, the fact that structures similar to the torus are seen at other wavelengths indicates that the torus is not a simple result of synchrotron burn-off.  Within the southern jet, photon index variations are also seen: the spectral softening takes place from the central core to the outer sheath. The photon index at the central core is almost the same as that of the torus. This indicates that the electron spectra are similar in the two different directions of the particle injection from the pulsar. We also found that the volume emissivities of the jet and the torus are similar.  Assuming that the brightness difference between the near and far sides of the torus is caused by Doppler boosting and relativistic aberration, the ratio can be explained by the observed speed of the downstream flow at the torus. However it cannot be obtained from the so-called weakly magnetized pulsar wind with $\\sigma \\approx 0.003$, as suggested by Kennel \\& Coroniti (1984).  Finally, we found that an optical filament comprised of supernova ejecta surrounding the pulsar wind nebula is absorbing the soft X-ray emission from a small portion of the X-ray nebula."
    },
    "0403/astro-ph0403552_arXiv.txt": {
        "abstract": " ",
        "introduction": "Chandra observatory, thanks to its sub-arcsec angular resolution, opened a new era in studying X-ray binary populations in nearby galaxies. For the first time an opportunity was presented to observe compact sources in a nearly confusion free regime. The long suspected fact has been proved, that X-ray binaries are an important, if not dominant, contributor to the X-ray emission of the normal galaxies,\\cite{fabbiano2003} as illustrated by the example of our Galaxy.\\cite{grimm02}  Depending on the mass of the optical companion, X-ray binaries are subdivided in to two classes -- low  and high mass X-ray binaries, having significantly different evolutionary time scale, $\\sim 10^{6-7}$ and $\\sim 10^{9-10}$ years respectively \\cite{xrbrev}. The nearly prompt emission of HMXBs  makes them a potentially  good tracer of the recent star formation activity in the host galaxy \\cite{rs78}.  The LMXBs, on the other hand, have no relation to the present star formation, but, rather, are related to the stellar content of the host galaxy \\cite{lmxb}. Chandra observations of the nearby galaxies presented a possibility to confirm this simple picture and to calibrate the HMXB--SFR\\cite{grimm03,lx-sfr,ranalli} and LMXB$-{\\rm M_*}$\\cite{colbert04,lmxb,kim2004} relations.   An unusual class of compact sources -- ultraluminous X-ray sources, has been discovered in nearby galaxies more than a decade ago \\cite{colbert99,fab89}. Although bright, $L_X>10^{39}$ erg/s, point-like sources are found both in young star forming galaxies and in old stellar population of elliptical and S0 galaxies, the most luminous and exotic objects are associated with actively star forming galaxies. Their nature and relation to more ordinary X-ray binaries is still a matter of a significant debate. Based on a simple Eddington luminosity argument, they appear to be powered by accretion onto an intermediate mass object -- a black hole with the mass in the hundreds-thousands solar masses range\\cite{fabbiano2003,miller2003}. However, a number of alternative models have been considered as well -- from collimated radiation\\cite{koerding2002} to  $\\sim$stellar mass black holes, representing the high mass tail of the standard stellar evolution sequence and accreting in the near- or slightly super-Eddington regime\\cite{king2001}.    \\begin{figure} \\centerline{\\hbox{ \\resizebox{0.45\\hsize}{!}{\\includegraphics{lf_all.ps}} \\resizebox{0.45\\hsize}{!}{\\includegraphics{lf_all_corr.ps}} }} \\caption{{\\em Left:} The XLFs of compact X-ray sources in nearby star forming galaxies. {\\em Right:} The XLFs of the same galaxies normalized to the star formation rates (from Grimm et al.\\cite{grimm03}). The XLF of the Milky Way includes only HMXBs. } \\label{fig:lf_hmxb} \\end{figure}   \\begin{figure} \\centerline{\\hbox{ \\resizebox{0.45\\hsize}{!}{\\includegraphics{clf_all_annulus_unsc.ps}} \\resizebox{0.45\\hsize}{!}{\\includegraphics{clf_all_annulus_scaled_mass.ps}} }} \\caption{{\\em Left:} The XLFs of compact X-ray sources in nearby elliptical, S0 galaxies and bulges of spiral galaxies. {\\em Right:} The same XLFs normalized to the stellar mass of the host galaxy (from Gilfanov\\cite{lmxb}). The XLF of the Milky Way includes only LMXBs. } \\label{fig:lf_lmxb} \\end{figure} ",
        "conclusions": ""
    },
    "0403/astro-ph0403078_arXiv.txt": {
        "abstract": "A population of black holes (BHs) at high redshifts ($z\\gsim 6$) that contributes significantly to the ionization of the intergalactic medium (IGM) would be accompanied by the copious production of hard ($\\gsim 10$ keV) X-ray photons.  The resulting hard X--ray background would redshift and be observed as a present--day soft X--ray background (SXB).  Under the hypothesis that BHs are the main producers of reionizing photons in the high--redshift universe, we calculate their contribution to the present--day SXB.  Our results, when compared to the unresolved component of the SXB in the range 0.5-2 keV, suggest that accreting BHs (be it luminous quasars or their lower--mass ``miniquasar'' counterparts) did not dominate reionization.  Distant miniquasars that produce enough X--rays to only partially ionize the IGM to a level of at most $x_e\\sim 50\\%$ are still allowed, but could be severely constrained by improved determinations of the unresolved component of the SXB. ",
        "introduction": "\\label{sec:intro} The recent discovery of the Gunn--Peterson (GP) troughs in the spectra of $z>6$ quasars in the Sloan Digital Sky Survey (SDSS; White et al. 2003; Wyithe \\& Loeb 2004) suggests that the end of the reionization process occurs at a redshift near $z\\sim 6$.  At this epoch, the ionizing sources drive a strong evolution of the neutral fraction of the intergalactic medium (IGM) from values near unity down to $x_{\\rm HI}\\sim 10^{-3}$ (e.g., Fan et al. 2002).  On the other hand, the high electron scattering optical depth, $\\tau_e=0.17 \\pm 0.04$, measured recently by the {\\it Wilkinson Microwave Anisotropy Probe (WMAP)} experiment (Spergel et al. 2003) suggests that ionizing sources were abundant at a much higher redshift, $z\\sim 15$.  These data imply that the reionization process is extended and complex, and is probably driven by more than one population of ionizing sources (see, e.g., Haiman 2003 for a post-{\\it WMAP} review).  The exact nature of these ionizing sources remains unknown.  Natural candidates to account for the onset of reionization at $z\\sim 15$ are massive, metal--free stars that form in the shallow potential wells of the first collapsed dark matter halos (Wyithe \\& Loeb 2003a; Cen 2003a; Haiman \\& Holder 2003).  The completion of reionization at $z\\sim 6$ could then be accounted for by a normal population of less massive stars that form from the metal--enriched gas in larger dark matter halos present at $z\\sim 6$.  The most natural alternative cause for reionization is the ionizing radiation produced by gas accretion onto an early population of black holes (``miniquasars''; see Haiman \\& Loeb 1998, Wyithe \\& Loeb 2003c, Bromm \\& Loeb 2003).  The ionizing emissivity of the known population of quasars diminishes rapidly beyond $z\\gsim 3$, and bright quasars are unlikely to contribute significantly to the ionizing background at $z\\gsim 5$ (Shapiro, Giroux \\& Babul 1994; Haiman, Abel \\& Madau 2001; Wyithe \\& Loeb 2003a).  However, if low--luminosity, yet undetected miniquasars are present in large numbers, they could dominate the total ionizing background at $z\\sim 6$ (Haiman \\& Loeb 1998). Recent work, motivated by the {\\it WMAP} results, has emphasized the potential significant contribution to the ionizing background at the earliest epochs ($z\\sim 15$) from accretion onto the seeds of would--be supermassive black holes (Madau et al. 2003; Ricotti \\& Ostriker 2003).  The soft X--rays emitted by these sources can partially ionize the IGM early on (Oh 2001; Venkatesan \\& Shull 2001).  A population of miniquasars at $z\\gsim 6$ would be accompanied by the presence of an early X-ray background.  Since the IGM is optically thick to photons with energies $E$ below $E_{\\rm max} = 1.8 [(1+z)/15)]^{0.5} x_{\\rm HI}^{1/3}\\hs {\\rm keV}$, the soft X-rays with $E\\lsim E_{\\rm max}$ would be consumed by neutral hydrogen atoms and contribute to reionization.  However, the background of harder X-rays would redshift without absorption and would be observed as a present--day soft X--ray background (SXB).  In this paper, {\\it we examine the hypothesis that accreting BHs are the main producers of reionizing photons in the high--redshift universe, and calculate their contribution to the present--day SXB in this case}.  Our results, when compared to the unresolved component of the SXB, suggest that accreting BHs cannot contribute significantly either to the completion of reionization at $z > 6$, or to the partial ionization of the IGM at $z\\gsim 15$ to ionized fractions of $x_e\\gsim 0.5$.  The outline of the paper is as follows: In \\S~\\ref{sec:method}, we describe the method to calculate the SXB from quasars that contribute to reionizing the universe. In \\S~\\ref{sec:sxrb}, we critically discuss current X-ray observations, focusing on the unresolved fraction of the SXB that could be attributed to distant quasars. In \\S~\\ref{sec:fullion}, we calculate the expected contribution to the SXB from hypothetical quasars and their lower--mass miniquasar counterparts that fully ionized the universe. In \\S~\\ref{sec:preion}, we repeat our analysis for a putative miniquasar population that partially ionizes the IGM at high redshifts. In \\S~\\ref{sec:discuss}, we discuss how various simplifications made in our analysis influence our final results. Finally, in \\S\\ref{sec:conclusions} we summarize our results and the implications of this work. Throughout this paper, we adopt the background cosmological parameters as measured by the {\\it WMAP} experiment, $\\Omega_m=0.27$, $\\Omega_{\\Lambda}=0.73$, $\\Omega_b=0.044$, and $h=0.71$ (Spergel et al. 2003) and set the mass fraction of helium to $Y_{\\rm He}=0.24$ ( e.g. Burles, Nollett, \\& Turner, 2001). In the rest of the paper `accreting BHs' will refer to both quasars and their lower--mass ``miniquasar'' counterparts.  \\newpage ",
        "conclusions": "\\label{sec:conclusions}  Our main results, obtained in \\S~\\ref{sec:quasars} and \\S~\\ref{sec:miniquasars}, are summarized in Table~\\ref{table2}.  The third column in this table denotes the percentage of the total SXB contributed by accreting BHs (which we used to refer to both quasars and their lower--mass ``miniquasar'' counterparts) for models in which high redshift accreting BHs contribute significantly to full or partial reionization.  These percentages can be compared directly with the percentages given by M03. The second column gives the total intensity of the quasar populations, to be compared to the allowed range of the mean or maximum unresolved SXB flux of $0.35 \\times 10^{-12}$ \\unitE to $1.23 \\times 10^{-12}$ \\unitE, respectively.  As Table~\\ref{table2} shows, models in which $z>6$ accreting BHs contribute to reionization overproduce the SXB.  Pre-ionization by miniquasars requires fewer ionizing photons, both because only a fraction of the H atoms need to be ionized, and because hard X--rays can produce multiple secondary ionizations.  We find that models in which X--rays are assumed to partially ionize the IGM up to $x_e \\sim 0.5$ at $6\\la z\\la 20$ are still allowed, but could be severely constrained by improved determinations of the unresolved component of the SXB.  We emphasize that our constraints derive from the total number of ionizing photons that the population as a whole needs to produce.  Therefore, our conclusions depend mostly on the assumed spectral shape, and are independent of the details of the population, such as the luminosity function and its evolution with redshift.  Future improvements in resolving the SXB, improving the limits on the unresolved component by a factor of a few, would place stringent constraints on the contribution of $z\\sim 15$ accreting BHs to the scattering optical depth measured by \\wmap."
    },
    "0403/astro-ph0403234_arXiv.txt": {
        "abstract": "We explore the self-enrichment hypothesis for globular cluster formation with respect to the star formation aspect. Following this scenario, the massive stars of a first stellar generation chemically enrich the globular progenitor cloud up to Galactic halo metallicities and sweep it into an expanding spherical shell of gas.  This paper investigates the ability of this swept proto-globular cloud to become gravitationally unstable and, therefore, to seed the formation of second generation stars which may later on form a globular cluster.  We use a simple model based on a linear perturbation theory for transverse motions in a shell of gas to demonstrate that the pressures by which the progenitor clouds are bound and the supernova numbers required to achieve Galactic halo metallicities support the successful development of the shell transverse collapse. Interestingly, the two parameters controling the metallicity achieved through self-enrichment, namely the number of supernovae and the external pressure, also rule the surface density of the shell and thus its ability to undergo a transverse collapse.  Such a supernova-induced origin for the globular cluster stars opens therefore the way to the understanding of the halo metallicity distributions.  This model is also able to explain the lower limit of the halo globular cluster metallicity. ",
        "introduction": "Globular clusters (GC) are dense, massive and round-shaped groups of stars present in the vast majority of galaxies. In our Galaxy, the halo GCs were among the very first bound structures to form and their study provides therefore valuable information about the early Galactic evolution.  Their formation is an exciting but yet unsolved problem.  For instance, it is still an open question whether Galactic halo GCs formed out of gas already chemically enriched ({\\it pre}-enrichment models, e.g. Harris \\& Pudritz 1994) or whether they produced their own heavy elements through an earlier generation of stars within the GC progenitor itself ({\\it self}-enrichment models).  In the second class of models, the issue of their formation is directly related to the origin of their metal content (Cayrel 1986; Brown, Burkert \\& Truran 1995; Parmentier et al.~1999). Such a feature makes the self-enrichment scenario especially appealing if we hypothesize that the GC progenitor clouds are made of primordial (i.e. metal-free) gas, a reasonable assumption for the Old Halo GCs, i.e. the population of old and coeval (Rosenberg et al.~1999) halo GCs.  Regarding the origin of the proto-globular cluster clouds (PGCC), Fall \\& Rees (1985) suggested that they formed out of the collapsing protoGalaxy, as cold and dense clouds in pressure equilibrium with a hot and diffuse background.  In the frame of this theory, the GC progenitor clouds are thermally supported (i.e. no additional support against gravitation due to magnetic fields and turbulence) and made of primordial gas.  In order to explain the metallicities of halo GCs, Parmentier et al.~(1999, hereafter Paper I) expanded the Fall \\& Rees (1985) model for GC formation by the self-enrichment picture. According to this one, a first stellar generation forms in the central regions of each PGCC.  When the massive stars explode as Type II supernovae (SNeII), they chemically enrich the surrounding gas and sweep the cloud, turning it in an expanding shell of gas in which the formation of a second, chemically enriched, stellar generation may be triggered.  These second generation stars form the proto-globular cluster. \\\\  The sweeping and compression of the interstellar medium by massive star explosions is not the sole mechanism invoked to account for star formation during the earliest stages of the Galactic evolution. Some other models assume a different origin for the trigger. For instance, following Vietri \\& Pesce (1995), the high pressure confining the GC gaseous progenitor leads to the propagation of a strong shock inwards the cloud, stimulating thereby the formation of new stars.  On another side, Murray \\& Lin (1992) and Dinge (1997) suggested that the propagation of shock waves is promoted by cloud-cloud collisions.  In a similar way, star forming clouds may have coalesced into larger units until they reach a density and/or accumulated mass high enough to enable the formation of bound globular clusters (Larson 1988, Smith 1999).  Obviously, several processes are able to collect ambient gas into dense layers/clouds where star formation can thereafter take place.  Several of them may have been at work at the same time.  Our interest being in understanding the origin of the metal content of GCs, in what follows, we investigate the hypothesis of star formation in gas layers swept by the explosions of PopIII massive stars located at the center of the PGCCs. \\\\  The debate of how PopIII stars looked like has been raging for many years. Numerous studies have addressed the issue of the collapse and fragmentation of primordial gas clouds in order to estimate the masses of PopIII stars. The primordial gas being deficient in heavy elements (i.e., the most efficient coolants in present-day star forming clouds), all of them have emphasized the importance of cooling by H$_2$ molecules. Inspite of this, the achieved conclusions do not necessarily converge. Recent numerical simulations (Abel, Bryan \\& Norman 2002) suggest that metal-free stars form in isolation and are massive (30 $\\lesssim$ M $\\lesssim$ 100~M$_{\\odot}$) objects.  From their own simulations, Bromm, Coppi \\& Larsson (1999) quoted a characteristic mass even larger than 100~M$_{\\odot}$.  These results are in marked contrast with the early study performed by Palla, Salpeter \\& Stahler (1983) following which primordial gas clouds should be capable of fragmenting into low-mass stars (i.e., down to $\\sim$ 0.1~M$_{\\odot}$). Nakamura \\& Umemura (1999) reached a somewhat intermediate conclusion. According to them, the mass range of the first stars is similar to its present value (i.e. no star with M $\\gtrsim$ 100~M$_{\\odot}$) but nevertheless excludes low-mass long-lived stars, that is, the lowest mass star allowed to form in a metal-free medium is a $\\simeq$ 3~M$_{\\odot}$ star.  In a more recent model, Nakamura \\& Umemura (2001) predict a bimodal initial mass function for PopIII stars.  In fact, the initial mass function of metal-free stars would show two distinct peaks at $\\simeq$ 1~M$_{\\odot}$ and $\\simeq$ 100~M$_{\\odot}$, that is, it would include intermediate mass (1 $\\lesssim$ M $\\lesssim$ 10~M$_{\\odot}$) stars, massive (10 $\\lesssim$ M $\\lesssim$ 100~M$_{\\odot}$) stars as well as very massive (M $\\gtrsim$ 100~M$_{\\odot}$) ones. As quoted by Christlieb et al.~(2003), the recent discovery of HE 0107-5240, the most metal-poor ([Fe/H]=$-$5.3) star ever discovered in the Galactic halo, could be a challenge to these models precluding the formation of stars with mass low enough for their life duration to exceed a Hubble time.  While there is currently a wide consensus that present star formation operates predominantly in a clustered mode (Lada, Strom \\& Myers 1993), the question whether clusters of metal-free stars managed to form in the first (proto-) galaxies is still open.  In this paper, we assume that at least some of the primordial star formation sites produced stellar {\\it clusters} whose most massive stars (10 $\\lesssim$ M $\\lesssim$ 50\\,M$_{\\odot}$) end their life as canonical SNeII.  Therefore, our study does not include the possibility of a prompt enrichment of the primordial interstellar medium by a population of very massive (i.e., $\\gtrsim$ 100~M$_{\\odot}$) stars as suggested by, e.g., Wasserburg \\& Qian (2000).  The chemical enrichment provided by isolated (very) massive objects will be qualitatively discussed in Sect.~3 where we will see that they may be appropriate to account for the formation of very metal-poor stars, i.e., stars more metal-poor than the most metal-deficient halo GCs ([Fe/H] $\\lesssim$ $-$2.5).  \\\\  Supernovae having long been thought to disrupt the cloud of gas out of which they have formed, Parmentier et al.~(1999) studied the ability of pressure-truncated clouds to retain SNII ejecta.  This ability will rule the metal content of the proto-cluster.  In fact, in this class of models, the final metallicity is determined by the number of supernovae ($N$, which, assuming a given initial mass function and given SNII yields, determines the amount of metals dispersed within the PGCC) and the background pressure ($P_h$, which determines the mass of the cloud, that is, the mass of primordial gas to be chemically enriched). Comparing the gravitational energy of the PGCC with the kinetic energy of the supershell resulting from the SNII explosions, Parmentier et al.(1999) showed that the gaseous GC progenitors can sustain up to $\\sim$ 200 supernovae (the disruption criterion, Eq.~14, Paper I).  This is a number high enough for the PGCCs to achieve halo metallicities.  Furthermore, for a given number of exploding massive stars, a self-enrichment episode in pressure-bound clouds lead to a metallicity gradient throughout the resulting system of GCs and to a correlation between the mass and the achieved metallicity of the gaseous progenitors in the sense that the least massive clouds are the most metal-rich.  Such a trend emerges because if the bound pressure is higher, the mass of the pressure-truncated cloud will be lower, and its ability to retain supernova ejecta will be greater.  These trends, i.e. a metallicity gradient and a mass-metallicity relation, are indeed statistically present in the Old Halo, that is, the halo from which the presumably accreted component has been removed (Parmentier et al.~2000, Parmentier \\& Gilmore 2001). \\\\  After having analysed to which extent the Galactic halo GC data fit the correlations induced by the self-enrichment process, the next step is to wonder whether there are some stars forming out of the chemically enriched supershell, i.e. whether there is a second stellar generation tracing these correlations. The idea of star formation induced by supernova explosions dates back at least to Opik (1953).  Numerous cases of distant star forming loops, connected with shells resulting from supernova explosions and located in the Galactic disc as well as in dwarf galaxies, are detailed in the literature (e.g. Comeron \\& Torra 1994, Walter et al. 1998, Efremov \\& Elmegreen 1998, Rubio et al. 1998).  In this paper, we address the specific case of supershells made of swept PGCCs and within which the formation of the stars of future GCs may be triggered. In a first step, we limit our study to the propagation of the shell throughout the hot protogalactic background in which the PGCCs are initially embedded, this part of the shell propagation being much longer than the propagation throughout the cloud (see Fig.~\\ref{fig:RsHPB}).  \\\\  The outline of the paper is as follows.  In Sect.~2, we solve the perturbed equations of continuity and motion for transverse flows within the shell (i.e. the swept cloud) in order to identify the conditions supporting a successful shell transverse collapse.  We also discuss in turn the impact of the different parameters acting upon the shell collapse and, thus, upon the temporal growth of the shell fragments in which further star formation may be stimulated.  In Sect.~3, we show how the conditions required to stimulate a star formation episode within the shell provides a natural explanation to the observed metallicity range of halo GCs and how the shape of their metallicity spectrum constitutes the next step to work on.  Sect.~4 describes some effects which our forthcoming computations should take into account in order to refine the present model.  Finally, our conclusions are presented in Sect.~5.  \\section[]{Stimulated Star Formation in Proto-Globular Cluster Clouds} \\label{sec:stimSF} The shell will in general contain perturbed (transverse) velocity components and perturbations in the column density whose development leads to the transverse collapse of the swept PGCC and, thereby, to the formation of a second stellar generation.  We now derive the conditions to get such a collapsing shell.  The elementary method described below is adopted. \\subsection{Modelling the transverse collapse of the shell} The computations are based on the linear perturbed equations of continuity and motion for transverse flows in the shell (e.g. Elmegreen 1994).  \\\\ The perturbed equation of continuity (mass conservation) is  \\begin{equation} \\frac{\\partial \\sigma _1}{\\partial t} = -2 \\frac{V_s}{R_s} \\sigma _1 - \\sigma _0 \\nabla _T . v\\,, \\label{eq:per_cont} \\end{equation}  where subscript T means that the gradient component under consideration is the transverse one, and the perturbed equation of motion (momentum conservation) is \\begin{equation} \\sigma _0 \\frac{\\partial v}{\\partial t} = - \\sigma _0 \\frac{V_s}{R_s} v - {c_s}^2~ \\nabla \\sigma _1 + \\sigma _0 g_1\\;. \\label{eq:per_motion} \\end{equation}  In these equations, $R_s$ and $V_s$ are respectively the radius and the velocity of the shell, $\\sigma _0$ is the unperturbed surface density, $\\sigma _1$ is the perturbed surface density, $v$ is the perturbed (transverse) velocity, $c_s$ is the velocity dispersion of the material inside the shell, $g_1$ is the perturbed gravity, this one being related to the surface density through (Elmegreen 1994): \\begin{equation} g _1 = -2 \\pi i G \\sigma _1\\;. \\label{g1_sigma1} \\end{equation} As mentioned above, the evolution of the shell is studied while propagating through the hot background, i.e. when the whole cloud has been swept inside the shell. The hot background being a diffuse medium, the shell mass $M_s(t)$ does not increase any longer at this stage of its propagation and is given by the mass $M$ of the progenitor cloud. The unperturbed surface density of the shell is thus given by \\begin{equation} \\sigma _0 = \\frac{1}{4 \\pi} \\frac{M_s(t)}{R_s ^2} = \\frac{1}{4 \\pi} \\frac{M}{R_s ^2}\\;. \\label{sigma_0} \\end{equation}   Equation \\ref{eq:per_cont} shows that the development with time of any perturbation of the shell surface density ($\\partial \\sigma _1 / \\partial t  > 0$) is inhibited by the stretching of the perturbed region due to the shell expansion (i.e. $V_s > 0$, first term on the right hand-side, hereafter rhs) while the convergence of the perturbed flows supports the growth of the perturbation (second term on the rhs). Equation \\ref{eq:per_motion} shows that an initial transverse flow of material along the shell develops ($\\partial v/ \\partial t  > 0$) only if the self-gravity (third term on the rhs) overcomes the stabilizing effects of the stretching (first term on the rhs) and of the internal pressure (second term on the rhs), here represented by $c_s^2$, the shell sound speed squared. \\\\  In order to solve Eqs.\\ref{eq:per_cont} and \\ref{eq:per_motion} properly, we now turn to the determination of the expansion law of the shell, i.e. $R_s(t)$ and $V_s(t)$, while it propagates throughout the hot protogalactic background.  \\subsection{Supershell propagation throughout the hot background} \\label{sub:propa_HPB}  \\begin{figure} \\begin{center} \\epsfig{figure=RsHPB.eps, width=\\linewidth} \\caption{Evolution with time of the shell radius inferred from Eqs.~(\\ref{eq:bubble_nrj_HPB} - \\ref{eq:mass_sh_HPB}) assuming that 200\\,SNeII explode at a constant rate during 30 million years and for three different hot protogalactic background pressures, from top to bottom 5$\\times 10^{-11}$ (plain curve), 10$^{-10}$ (dashed-dotted curve) and 5$\\times 10^{-10}$\\,dyne.cm$^{-2}$ (dotted curve).  In each case, an arrow indicates the time $t_{em}$ at which the shell crosses the interface between the cold and hot phases} \\label{fig:RsHPB} \\end{center} \\end{figure}  The equations describing the propagation of a supernova-driven shell are as follow (Castor et al.~1975). \\begin{enumerate} \\item The supernova explosions add energy to the bubble at a constant rate $\\dot{E_o}$ and the dominant energy loss of the bubble comes from the work against the dense shell, hence the variation with time of the energy $E_b$ of the bubble obeys \\begin{equation} \\dot{E_b}=\\dot{E_o}-4\\pi {R_s}^2 P_b \\dot{R_s}\\,. \\label{eq:bubble_nrj_HPB} \\end{equation}  We assume that the kinetic energy of every SN is 10$^{51}$~ergs and that the SN phase lasts about thirty million years, i.e. $\\dot{E_o}=N 10^{51} {\\rm ergs}/30 {\\rm Myr}$.  \\item The internal energy $E_b$ and the pressure $P_b$ of the bubble are related through \\begin{equation} \\frac{4\\pi}{3} {R_s}^3 P_b = \\frac{2}{3} E_b\\,, \\label{eq:bubble_P_HPB} \\end{equation}  \\item The shell motion obeys Newton's second law \\begin{equation} \\frac{d}{dt} [M_s(t) \\dot{R_s}(t)] = 4\\pi {R_s}^2 (P_b-P_{ext}) -\\frac{{GM_s}^2(t)}{2{R_s}^2(t)}\\,, \\label{eq:Newton2_HPB} \\end{equation} where $P_{ext}$ is the pressure of the medium just outside the shell. \\item Considering the case of swept PGCCs propagating through the hot protogalactic background, the mass of the shell is constant in time and is given by the mass of the cloud (see Sect.~2.1): \\begin{equation} M_s(t) = M\\;. \\\\ \\label{eq:mass_sh_HPB} \\end{equation} \\end{enumerate}  The pressure external to the shell is exerted by the surrounding hot background (i.e. P$_{ext}$ = P$_h$ in Eq.~\\ref{eq:Newton2_HPB}) and is therefore assumed to be constant in time.  This pressure will strongly decelerate the shell as illustrated below. Numerical resolution of Eqs.~(\\ref{eq:bubble_nrj_HPB} - \\ref{eq:mass_sh_HPB}) provides $R_s(t)$ and thereby $V_s(t)$ and $\\sigma _0(t)$.  The initial conditions are those at the time $t_{em}$, i.e. when the shell crosses the interface between the cloud and the surrounding background: the mass and radius of the shell are those of the pressure-bound cloud ($R_s(t_{em})=R$, $M_s(t_{em})=M$), the velocity of the shell is determined by its former propagation at constant speed $V$ through the cloud (see Paper I, Eq.~13), that is $V_s=V$ and $A_s=0$. \\\\  \\begin{figure} \\begin{center} \\epsfig{figure=RsHPBt_3.eps, width=\\linewidth} \\caption{Comparison between the evolution with time of the shell radius computed from Eqs.~(\\ref{eq:bubble_nrj_HPB} - \\ref{eq:mass_sh_HPB})  and the average radius given by Eq.~9 (i.e. $<R_s(t)> \\propto t^{1/3}$) for the quoted numbers $N$ of SNeII and background pressures $P_h$} \\label{fig:RsHPBt1/3} \\end{center} \\end{figure}  Figure \\ref{fig:RsHPB} shows the evolution with time of the shell radius for 200\\,SNeII, i.e. the maximum number of supernovae that a PGCC can sustain (disruption criterion, Paper I), and for 3 different values of the hot protogalactic background pressure, i.e. $P_h$ = 5$\\times 10^{-11}$, 10$^{-10}$ and 5$\\times 10^{-10}$\\,dyne.cm$^{-2}$. The SN rate being the same in the three cases, the shells propagate at the same velocity in the cold phase (Eq.~13 in Paper I).  Their expansions begin to differ once they have crossed the interface between their respective PGCC and the hot background in which the latter is embedded, this time being indicated by an arrow in Fig.~\\ref{fig:RsHPB}.  Obviously, the propagation of the shell through the cloud is much shorter than its propagation through the background.   A significant part of the shell expansion through the background takes place at early time.  Indeed, after a transient phase during which the velocity of the shell does not differ markedly from its velocity inside the cloud (Fig.~\\ref{fig:RsHPB}), the overall expansion slows down and the radius of the shell scales roughly as $t^{1/3}$. This average expansion can be obtained from Eqs.~(\\ref{eq:bubble_nrj_HPB} - \\ref{eq:mass_sh_HPB}) assuming that $M_s(t)=0$, reflecting thereby that the temporal evolution of the shell radius does not depend strongly on the mass (Brown et al.~1995):  \\begin{equation} <R_s(t)> = \\left (\\frac{3}{10\\pi} \\frac{\\dot{E_o}}{P_h} \\right )^{1/3} t^{1/3}. \\label{eq:RsHPBt1/3} \\end{equation}  Figure \\ref{fig:RsHPBt1/3} shows the good agreement between Eq.~\\ref{eq:RsHPBt1/3} and the result of the numerical integration over time of Eqs.~(\\ref{eq:bubble_nrj_HPB} - \\ref{eq:mass_sh_HPB}). Equation \\ref{eq:RsHPBt1/3} shows that, during the long-term evolution, the expansion rates of two shells having the same $(N/P_h)$ ratio are similar. Figure \\ref{fig:RsHPBt1/3} illustrates this effect for two sets of values, namely $N$=100 and $P_h$=5$\\times 10^{-11}$\\,dyne.cm$^{-2}$ (plain curve) and $N$=200 and $P_h$=$10^{-10}$\\,dyne.cm$^{-2}$ (dashed-dotted curve).  \\subsection{Growth with time of the shell fragments} \\label{sec:frag_num}  In order to assess whether the shell transverse collapse proceeds successfully or not, we now numerically integrate over time Eqs.\\ref{eq:per_cont} and \\ref{eq:per_motion} in order to derive the temporal evolutions of the perturbed and unperturbed surface densities, $\\tilde{\\sigma} _1(t)$ and $\\sigma _0(t)$, respectively.  Assuming that the perturbed quantities follow a complex exponential of the angular position $\\phi$ along the shell, we get: \\begin{equation} \\sigma _1 (t,\\phi) = \\tilde{\\sigma} _1 (t) ~e^{-i \\eta \\phi} \\label{sigma1_exp} \\end{equation} and \\begin{equation} v (t,\\phi) = \\tilde{v} (t)  ~e^{-i \\eta \\phi} ~e^{i \\Delta \\phi}\\,. \\label{v_exp} \\end{equation}  In these equations, $\\Delta \\phi$ represents the phase difference between the perturbed surface density $\\sigma _1$ and the perturbed velocity $v$. $\\eta$ is the angular wavenumber and is related to the spatial wavenumber $k$ by: \\begin{equation} \\eta = k R_s = \\frac{2 \\pi}{\\lambda}~R_s \\label{eta_lambda} \\end{equation} where $\\lambda$ is the wavelength of the perturbation, namely the average distance between forming fragments, the sites of future star formation.  Therefore, $\\eta$ is the number of forming clumps along a shell circumference and, as such, $\\eta$ must be an integer. Moreover, any realistic perturbation must fit inside a fraction of the shell circumference, say, $\\lambda \\leq R_s$ or $\\eta \\geq 6$.  \\\\  At a time $t$ and an angular position $\\phi$ along the shell, the shell surface density $\\sigma _s$ obeys \\begin{equation} \\sigma _s(t,\\phi )=\\sigma _0(t)+\\sigma _1(t,\\phi ) =\\sigma _0(t)+\\tilde{\\sigma _1}(t)~cos(\\eta \\phi)\\;. \\label{eq:sigt_t_anlt} \\end{equation}  The fragmentation of the shell, i.e.~its fully developed transverse collapse, occurs when \\begin{equation} \\tilde{\\sigma} _1(t) = \\sigma _0(t). \\end{equation}  Writing $- i \\eta /R_s$ for the gradient and using Eqs.~\\ref{g1_sigma1} and \\ref{sigma_0}, Eqs.~\\ref{eq:per_cont} and \\ref{eq:per_motion} successively become:  \\begin{itemize} \\item[$\\triangleright$] Perturbed equation of continuity:  \\begin{eqnarray} \\frac{\\partial \\tilde{\\sigma _1}}{\\partial t} & = & -2\\frac{V_s}{R_s} ~\\tilde {\\sigma _1} +  \\sigma _0 ~ \\frac{i~\\eta}{R_s} ~ \\tilde{v}~e^{i\\Delta \\phi} \\nonumber \\\\ \\frac{\\partial \\tilde{\\sigma _1}}{\\partial t} & = & -2\\frac{V_s}{R_s} ~\\tilde{\\sigma _1} +  \\frac{i~\\eta~M}{4 \\pi} \\frac{\\tilde{v}}{{R_s}^3} ~e^{i\\Delta \\phi} \\label{eq:per_cont_num} \\end{eqnarray}  \\item[$\\triangleright$] Perturbed equation of motion:  \\begin{eqnarray} \\sigma _0 ~e^{i\\Delta \\phi} ~\\frac{\\partial \\tilde{v}}{\\partial t} & = & -\\sigma _0 ~ \\frac{V_s}{R_s} ~ \\tilde{v}~e^{i\\Delta \\phi} + {c_s}^2 ~ \\frac{i~\\eta}{R_s} ~ \\tilde{\\sigma _1} \\nonumber \\\\ & & \\,-  2 \\pi i G ~ \\sigma _0 ~ \\tilde{\\sigma _1} ~~~ \\nonumber \\\\ ~\\frac{\\partial \\tilde{v}}{\\partial t} & = & - \\frac{V_s}{R_s} ~ \\tilde{v} \\hspace{1.3cm} \\,+ \\frac{4~\\pi ~i ~\\eta ~c_s^2}{M} R_s ~\\tilde{\\sigma _1} ~ e^{- i\\Delta \\phi} \\nonumber \\\\ & & -  2 \\pi i G \\tilde{\\sigma _1} e^{- i\\Delta \\phi} ~~~ \\label{eq:per_motion_num} \\vspace*{3pt} \\end{eqnarray} \\end{itemize}  Seven parameters intervene in the ability of the shell to undergo a transverse collapse: \\begin{itemize} \\item[- ] the supernova number $N$ and the pressure external to the shell $P_h$, which determine the shell expansion law ($R_s$ depends on $N$ and $P_h$) and mass ($M$ depends on $P_h$), \\item[- ] the initial values of the perturbed quantities, $\\tilde \\sigma _1(t_{em})$ and $\\tilde v (t_{em})$, and the associated phase difference $\\Delta \\phi$, \\item[- ] the number $\\eta$ of forming clumps embedded along a shell circumference, \\item[- ] the velocity dispersion $c_s$ of the shell material which is related to the internal pressure $P_s$ of the shell, supporting it against transverse collapse, through $P_s={c_s}^2 \\rho _s$ ($\\rho _s$ is the volumic mass density of the shell material). \\end{itemize}  In what follows, the influence of each of the parameters involved in the fragmentation process is investigated, the final aim being to check whether some reasonable sets of conditions can lead to the shell fragmentation.  \\subsubsection{The phase difference between $v$ and $\\sigma_1$: $\\Delta \\phi$} \\label{sub:delta_phi}  \\begin{figure} \\hspace*{5mm} \\epsfig{figure=frag_shell.eps, width=11cm} \\caption{Schematic description of the sinusoidal (i.e. complex exponential) behaviour with $\\phi$ of the perturbed quantities $v$ (lower curve) and $\\sigma _1$ (upper curve) along a fraction of the shell circumference at a given time assuming that the number of clumps $\\eta$ is 15 and the phase difference $\\Delta \\phi$ between $v$ and $\\sigma _1$ is $-\\frac{\\pi}{2}$.  Such a phase difference corresponds to the convergence of the transverse (perturbed) flows (indicated by the thick arrows) towards the initial clumps (filled circles). The open circles represent the shell regions which are progressively depleted by the transverse flows} \\label{fig:fragshell} \\end{figure}  The amplitude of the perturbed surface density at a time $t$, $\\tilde \\sigma_1(t)$, results from its initial value $\\tilde \\sigma_1(t_{em})$ and from the transverse flows which redistribute the shell mass accumulated while the shell was propagating throughout the progenitor cloud.  The fragmentation is therefore favoured if these transverse flows converge towards the clumps initially present within the shell. Such a situation corresponds to a phase difference $\\Delta \\phi$ = $-\\frac{\\pi}{2}$ between the two perturbed quantities $\\sigma _1(t,\\phi)$ and $v(t,\\phi)$, that is,  the transverse velocity exhibits a phase delay of a quarter of a wavelength with respect to the perturbed surface density (see Fig.~\\ref{fig:fragshell}). Replacing $\\Delta \\phi$ by this value, Eqs.~\\ref{eq:per_cont_num} and \\ref{eq:per_motion_num} respectively become:  \\begin{equation} \\frac{\\partial \\tilde{\\sigma _1}}{\\partial t}=-2\\frac{V_s}{R_s} ~\\tilde {\\sigma _1} + \\frac{\\eta~M}{4 \\pi} \\frac{\\tilde{v}}{{R_s}^3} \\label{eq:per_cont_num_phi}\\,, \\end{equation}  \\begin{equation} \\frac{\\partial \\tilde{v}}{\\partial t} =-~\\frac{V_s}{R_s} ~ \\tilde{v} - \\frac{4~\\pi  ~\\eta ~c_s^2}{M} R_s ~\\tilde{\\sigma _1} + 2\\pi G ~ \\tilde{\\sigma _1}\\;.  \\\\ \\label{eq:per_motion_num_phi} \\end{equation}  In what follows, the roles of the other parameters is studied assuming that $\\Delta \\phi = - \\pi $/2, i.e. through the numerical integration of Eqs.~\\ref{eq:per_cont_num_phi} - \\ref{eq:per_motion_num_phi}.  \\subsubsection{The global parameters: $N$ and $P_h$} \\label{sub:discNPh}  In order to assess whether the shell fragments or not during its propagation through the hot protogalactic background, we compare in Fig.~\\ref{fig:disc_sig1init} the evolutions with time of the perturbed and unperturbed surface densities, namely $\\tilde{\\sigma} _1(t)$ and $\\sigma _0(t)$, for different values of the external pressure $P_h$ (5$\\times 10^{-10}$ and 10$^{-10}$\\,dyne.cm$^{-2}$) and numbers $N$ of SNeII (100 and 200). The corresponding metallicity is indicated in each panel. The other parameters are kept the same in every case: the sound speed $c_s$ of the shell material is 1\\,km.s$^{-1}$, the perturbed velocity $v(t_{em})$ is 0.01\\,$V_s(t_{em})$, the number of forming clumps $\\eta$ is 10 and the initial perturbed surface density is assumed to be of order one per cent of the unperturbed value, namely $\\sigma _1(t_{em}) = 0.01 \\times \\sigma _0(t_{em})$. This choice may seem rather arbitrary but the next paragraph will show that the initial amplitude of  $\\sigma _1$ does not affect the results significantly .  \\\\  While the unperturbed surface density decreases due to the shell expansion at constant mass, the perturbed surface density grows at a rate which depends on the external pressure and on the number of SNeII. The comparison of the different panels in Fig.~\\ref{fig:disc_sig1init} shows that the development of $\\tilde{\\sigma} _1(t)$ with respect to $\\sigma_0(t)$ is favoured by a {\\sl a high background pressure} and a {\\sl low number of SNeII}. Among the three panels presented ($P_h=10^{-10}$\\,dyne.cm$^{-2}$ and $N$=100, $P_h=5\\times 10^{-10}$\\,dyne.cm$^{-2}$ and $N$=100, $P_h=5\\times 10^{-10}$\\,dyne.cm$^{-2}$ and $N$=200), the transverse collapse of the shell is successfully achieved\\footnote{At this point, it should be kept in mind that once the shell has achieved its complete transverse collapse, Eqs.~\\ref{eq:per_cont} and Eqs.~\\ref{eq:per_motion} are no longer valid.  Therefore, once ${\\sigma} _0 \\leq \\tilde\\sigma _1$, the curve $\\tilde \\sigma _1(t)$ has no longer physical meaning.} when the lower number of supernovae is combined with the larger external pressure ($P_h=5\\times 10^{-10}$\\,dyne.cm$^{-2}$ and $N$=100). This effect comes from the dependence of the shell surface density on the external pressure $P_h$ and the supernova number $N$.  Indeed, combining the mass of a pressure-truncated gas cloud (i.e. $M \\propto P_h^{-1/2}$) with Eq.~\\ref{eq:RsHPBt1/3}, the surface density of the shell scales as: \\begin{equation} \\sigma _0  (t) \\propto \\frac{M}{R_s(t)^2} \\propto \\frac{P_h^{1/6}}{N^{2/3}}. \\label{eq:sig0_N_Ph} \\end{equation} Consequently, the larger the pressure and/or the lower the supernova number, the larger the shell surface density and the larger the growth rate of the perturbation (see Eq.~\\ref{eq:per_cont} and Eq.~\\ref{eq:per_motion}, a larger surface density favours the terms promoting an efficient transverse collapse). {\\sl Therefore, although the presence of exploding massive stars is required to stimulate the formation of a second stellar generation, too large a number of SNeII inhibits the ability of the shell to collapse and to form new stars}\\footnote{Unlike the protogalactic shells we study, the collapse of shells expanding in the Galactic disc is made easier by a larger number of SNeII.  In fact, assuming that the surrounding interstellar medium is roughly homogeneous, i.e. the shell radius is smaller than the density scale-height of the Galactic HI layer, one gets by mass conservation, assuming a pre-shell density $\\rho _0$, $\\sigma _0 = \\frac{R_s \\rho _0}{3}$.  In this case, a larger number of SNeII leads to a larger radius at a given time and, therefore, to larger surface density and perturbation growth rate.}.  The influence of the surface density of the shell on its collapse is reminiscent of the star formation law on large scales, i.e. averaged over entire galactic discs.  In that case, providing that the gas surface density is larger than a density threshold, the star formation rate follows a power-law of the gas surface density (the so-called Schmidt law) while it falls sharply below (Kennicutt 1989).  The threshold surface density varies from one galaxy to another but remains nevertheless in the range 10$^{20}$-- 10$^{21}$\\,cm$^{-2}$.  It may not be a coincidence that the shell surface densities obtained in the frame of this model are of the same order of magnitude or even larger (see Sect.~2.3.3).  The dependence of the perturbation growth rate on $N$ and $P_h$ raises a point of interest, worthy of a mention here. It just so happens that the parameters determining the final metallicity of the proto-cluster, namely the number of supernovae ($N$, which determines the amount of metals dispersed within the PGCC) and the background pressure ($P_h$, which determines the mass of the cloud, that is, the mass of primordial gas to be chemically enriched) are also some of those influencing the ability of the shell to form new stars.  {\\sl If some combinations of external pressures and SN numbers support the completion of the shell transverse collapse much more than some others do, then the formation of second stellar generations with the corresponding metallicities will be favoured}.  Therefore, in the frame of the self-enrichment scenario, there is a direct link between the achieved metallicity and the probability of forming halo stars, i.e. {\\sl the study of the fragmentation process may shed light on the metallicity distribution function of Galactic halo field stars and GCs}.  Figure \\ref{fig:disc_sig1init} and Eq.~19 show that the shell fragmentation is favoured by a low number of SNe and a high pressure of the hot protogalactic background.  Now, let us imagine that, on the contrary, the collapse efficiency increases with both decreasing number of SNeII and pressure.  As $P_h$ and $N$ get smaller, the number of successful transverse collapses increases and the newly formed stars are more metal-poor. Thus, the metallicity distribution function of the Galactic halo would exhibit an increasing number of proto-clusters/halo stars with decreasing metallicity.  If, on the other hand, the shell ability to achieve fragmentation increased with both increasing external pressure and SN number, then the larger $P_h$ and $N$, the more numerous the transverse collapses and the more metal-rich the newly formed stars.  As a result, there would be an increasing number of proto-clusters/halo stars with increasing metallicity. Neither an increasing nor a decreasing metallicity distribution function is observed for the Galactic halo.  In contrast, the metallicity distributions, for both halo field stars (Laird et al.~1988) and halo GCs (Zinn 1985), are peaked-shape.  While the finding that the shell transverse collapse is favoured by large external pressures (promoting ``large'', i.e. mildly metal-poor, metallicities) and low SN numbers (supporting low metallicities) is not sufficient in itself to draw some definitive conclusions regarding the shape of the halo metallicity distribution, it appears that it does not contradict it either.  \\subsubsection{The initial perturbed surface density: $\\tilde{\\sigma} _1(t_{em})$} \\label{sub:disc_sig1init}  \\begin{figure} \\epsfig{figure=discNPh_4.eps, width=\\linewidth} \\epsfig{figure=discNPh_5bis.eps, width=\\linewidth} \\epsfig{figure=discNPh_6.eps, width=\\linewidth} \\caption{Temporal evolution of $\\sigma _0$ (plain curves) and $\\tilde \\sigma _1$ (dashed-dotted curves), expressed in units of 1\\,M$_{\\odot}$.pc$^{-2}$, considering two different initial perturbed surface densities (see the keys) and the external pressures and SNII numbers indicated.  In each panel, $c_s$=1\\,km.s$^{-1}$, $\\eta=10$ and $v(t_{em})=0.01\\,V_s(t_{em})$} \\vspace{11pt} \\vspace{11pt} \\label{fig:disc_sig1init} \\end{figure}  Protogalactic shells such as those studied here exhibit number surface densities of order several 10$^{20}$\\,cm$^{-2}$. Indeed, using Eqs.~\\ref{sigma_0} and \\ref{eq:RsHPBt1/3}, one gets: \\begin{eqnarray} \\sigma _0(t, N, P_h) & \\simeq & 5 \\times 10^{-3} P_{h(10)}^{1/6} N_{200}^{-2/3} t_6^{-2/3}\\, {\\rm g.cm^{-2}} \\nonumber \\\\ & \\simeq & 2.5 \\times 10^{21} P_{h(10)}^{1/6} N_{200}^{-2/3} t_6^{-2/3}\\,{\\rm cm^{-2}} \\end{eqnarray}  where the subscript (10) means that the pressure is expressed in units of $10^{-10}\\,\\rm dyne.cm^{-2}$, the subscript 200 means that $N$ is expressed in units of 200\\,SNe and the subscript 6 means that the time is expressed in units of $10^6$\\,years.  Considering the interstellar medium of the Galactic disc with roughly the same surface density (though a bit lower, i.e. $\\simeq$ 10$^{20}$\\,cm$^{-2}$), the relative non-uniformities in the number surface density are typically of order 0.01 - 0.2 (Wunsch \\& Palous 2001).  However, the initial perturbation in the shell surface density is certainly not as large as some of the clumps of this {\\sl quiescent} disc interstellar medium since {\\sl turbulent mixing} is expected to take place within the shell (this is an important requirement to achieve the efficient mixing of the supernova heavy elements with the cloud gas and, therefore, to maintain the chemical homogeneity of the shell, that is, of the proto-globular cluster; see Brown et al.~1991), thus lowering the surface density inhomogeneities. \\\\  Figure 4 also shows the temporal evolution of $\\sigma _0$ and $\\tilde \\sigma _1$ assuming that the initial perturbed surface density is lower than in Sect.~2.3.2, namely 0.25\\%\\,$\\sigma _0(t_{em})$ instead of 1\\%\\,$\\sigma _0(t_{em})$. Despite these different initial values, each panel displays almost identical $\\tilde \\sigma _1(t)$ curves in both cases.  Owing to the initial transverse flows, the initial clumps quickly undergo a replenishment in shell material leading afterwards to very similar temporal evolutions of the perturbed surface density whatever the initial perturbed surface density.  Obviously, this one is not a key parameter of the fragmentation process.  \\subsubsection{The initial perturbed velocity: $v(t_{em})$} \\label{sub:disc_vtrans}  All the results presented here above assume the spherical symmetry of the supershell.  It is clear however that such a system is not expected to remain perfectly spherical.  The deviations from spherical symmetry can arise, for instance, from the non point-like nature of the energy input, namely all the stars of the first generation cluster are not located exactly at the centre of the cloud/shell. A rough estimate of the initial amplitude of the transverse motions within the shell can therefore be derived from the size of this cluster. A star located at a distance $r$ from the shell centre will induce a transverse velocity $v$ such that:  \\begin{equation} v ~\\simeq ~V_s ~ \\frac{r}{R_s}\\;. \\vspace{11pt} \\label{eq:vt_1} \\end{equation}  In order to estimate the size $r$ of the cluster of massive stars hosted by a PGCC, we refer to R\\,136, the dense core of the 30\\,Doradus Nebula located in the Large Magellanic Cloud.  The 30\\,Doradus nebula shows an impressive example of a two-stage stellar formation.  The energetic activity of a very compact bright cluster, R136, which includes several tens of O stars, triggers the formation of a new stellar generation revealed by numerous infrared sources in or near some bright filaments west and northeast of R136 (Rubio et al.~1998).  Based on Hubble Space Telescope photometry, Campbell et al.~(1992) detected about 160 stars more massive than 10\\,M$_{\\odot}$ in R\\,136 which they define as a region of 2.2\\,pc $\\times$ 1.9\\,pc.  This number of massive stars being remarkably similar to the numbers of SNeII used in our self-enrichment model, we adopt R\\,136 as the most similar example in the Local Group of what may have been the first generation cluster. A radius of 1\\,pc appears therefore as a reasonable estimate of the size of this cluster, the source of the energy input.  \\\\  Considering an average background pressure of 10$^{-10}$\\,dyne.cm$^{-2}$ and the corresponding cloud radius ($\\simeq$ 30\\,pc, see Eq.~3, Paper I), the transverse velocity when the shell reaches the cloud boundary is: \\begin{equation} v(t_{em}) ~\\simeq ~\\frac{1\\,pc}{30\\,pc} ~V_s(t_{em}) ~\\simeq 0.03 V_s(t_{em})\\;. \\label{eq:vt_2} \\end{equation}  Figure 5 displays the evolution with time of $\\sigma _0$ and $\\tilde \\sigma _1$ assuming three different values for the initial transverse velocity, namely 0.01\\,$V_s(t_{em})$, 0.02\\,$V_s(t_{em})$ and 0.03\\,$V_s(t_{em})$, while keeping all the other parameters to their previous values.  In sharp contrast with $\\tilde \\sigma _1(t_{em})$, $\\tilde v(t_{em})$ appears to be an important parameter of the shell transverse collapse.  For instance, considering the bottom panel in Fig.~5 ($P_h=5 \\times 10^{-10}$\\,dyne.cm$^{-2}$ and $N$=200), the fragmentation takes place less than 15 million years after the first SN explosion if $v(t_{em})=0.03 \\,V_s(t_{em})$, whereas it is prevented if $v(t_{em})=0.01 \\,V_s(t_{em})$.  \\begin{figure} \\epsfig{figure=discNPh_12.eps, width=\\linewidth} \\epsfig{figure=discNPh_11bis.eps, width=\\linewidth} \\epsfig{figure=discNPh_10.eps, width=\\linewidth} \\caption{Temporal evolution of $\\sigma _0$ (plain curves) and $\\tilde \\sigma _1$ (dashed-dotted curves), expressed in units of 1\\,M$_{\\odot}$.pc$^{-2}$, considering three initial perturbed velocities (see the keys) and the external pressures and SNII numbers indicated.  In each panel, $c_s$=1\\,km.s$^{-1}$, $\\eta$=10 and $\\tilde \\sigma _1(t_{em})=0.01\\,\\sigma _0 (t_{em})$} \\vspace{11pt} \\vspace{11pt} \\label{fig:disc_vtrans} \\end{figure}  \\subsubsection{The number of forming clumps: $\\eta$} \\label{sub:disc_eta}  \\begin{figure} \\epsfig{figure=discNPh_7.eps, width=\\linewidth} \\epsfig{figure=discNPh_8bis.eps, width=\\linewidth} \\epsfig{figure=discNPh_9.eps, width=\\linewidth} \\caption{Temporal evolution of $\\sigma _0$ (plain curves) and $\\tilde \\sigma _1$ (dashed and dotted curves), expressed in units of 1\\,M$_{\\odot}$.pc$^{-2}$, considering the external pressures, SNII numbers and numbers of forming clumps indicated (see the keys).  In each panel, $c_s$=1\\,km.s$^{-1}$, $\\tilde \\sigma _1(t_{em})=0.01\\sigma _0(t_{em})$, $v(t_{em})=0.01\\,V_s(t_{em})$} \\vspace{11pt} \\label{fig:disc_eta} \\end{figure}  In supernova-driven shells of gas, star formation does not take place all along the whole periphery but, instead, takes place in regularly spaced  clumps (i.e., where the transverse collapse brings some gas from depleted adjoining regions in the shell, see Fig.~\\ref{fig:fragshell}). The number $\\eta$ of such clumps along a shell circumference is thus related to the wavelength $\\lambda$ of the perturbation (Eq.~\\ref{eta_lambda}).  This aspect of shell collapse modelling following which stars are formed in discrete stellar subsystems along a shell periphery is supported by several observations of star forming shells, the so-called Sextant being one of the most illustrative. In the Large Magellanic Cloud, a set of five OB associations are located along a HI supershell, sustaining there about 1/6 of a complete circle and being thus called the Sextant (Efremov \\& Elmegreen 1998). Efremov, Ehlerova \\& Palous (1999) showed that the OB associations are regularly spaced, the deprojected average distance between two subsequent stellar groups being $\\sim$ 37~pc.  Furthermore, using numerical simulations of shells expanding in a mass model of the Large Magellanic Cloud, they concluded that the formation of these 5 OB associations is most probably the result of a triggered star formation episode in the supershell created by SNeII located near the Sextant centre. Their simulations also explain how projection effects make the visible star forming regions only a fraction of a total circle.   \\\\  The influence of the number of clumps embedded within the shell on the transverse collapse is illustrated in Fig.~\\ref{fig:disc_eta}. The parameter $\\eta$ is not as negligeable as the initial perturbed surface density.  In some cases, it acts upon the fragmentation process almost as strongly as the initial transverse velocity. It is thus worth estimating the angular wavenumbers which are the most favourable to the growth of an initial perturbation. In order to do so, we now derive an analytical approximation of the instantaneous growth rate of the perturbation.  This is done straightforwardly by assuming that the perturbed quantities vary exponentially with time $t$, in analogy with the exponential growth rate found in other instability problems (e.g. the Jeans mass) even though there is no real exponential growth in this problem owing to the shell expansion and the corresponding time dependence of $\\sigma _0$.  It is thus important to keep in mind that the perturbation growth rate derived below (Eqs.~\\ref{eq:omega_eta} and \\ref{eq:fg_omega}) is {\\it not} used to compute the temporal evolution of $\\sigma _1$ (the cases displayed in Fig.~\\ref{fig:disc_eta} are obtained from solving Eqs~\\ref{eq:per_cont_num_phi}-\\ref{eq:per_motion_num_phi}) but merely to derive $\\eta$ estimates favouring the shell collapse. The perturbed quantities are written: \\begin{equation} \\sigma _1 (t,\\phi) = \\tilde{\\sigma}_1(t_{em}) ~e^{\\omega t} ~e^{-i \\eta \\phi} \\label{sigma1_exp} \\end{equation} and \\begin{equation} v (t,\\phi) = \\tilde{v}(t_{em}) ~e^{\\omega t}  ~e^{-i \\eta \\phi} ~e^{i \\Delta \\phi}, \\label{v_exp} \\end{equation}  where $\\omega$ is the angular frequency of the perturbation. Following Eqs.~\\ref{sigma1_exp} and \\ref{v_exp}, we write $\\omega$ for the time derivatives and $- i \\eta/R_s$ for the transverse gradients. Therefore, Eqs.~\\ref{eq:per_cont} and \\ref{eq:per_motion} become respectively  \\begin{equation} \\omega \\sigma _1 = -2 \\frac{V_s}{R_s} \\sigma _1 + \\sigma _0 \\frac{i \\eta}{R_s} v \\label{eq:per_cont2} \\end{equation}  and  \\begin{equation} \\sigma _0 \\omega v = - \\sigma _0 \\frac{V_s}{R_s} v + {c_s}^2~ \\frac{i \\eta}{R_s}\\sigma _1 - 2 \\pi i G \\sigma _0 \\sigma _1 \\,, \\label{eq:per_motion2} \\end{equation}  using Eq.~\\ref{g1_sigma1} in the latter.  \\\\  The elimination of the perturbed quantities $\\sigma _1$ and $v$ between Eqs.~\\ref{eq:per_cont2} and \\ref{eq:per_motion2} provides the dispersion equation, namely the relation between the angular frequency $\\omega$ (i.e. the instantaneous growth rate) and the angular wavenumber $\\eta$ of the perturbation: \\begin{equation} \\sigma _0 \\left( \\omega + 2 \\frac{V_s}{R_s} \\right) \\left( \\omega + \\frac{V_s}{R_s} \\right) - \\sigma _0 \\frac{i \\eta}{R_s} \\left( 2 \\pi i G \\sigma _0 - c_s^2 \\frac{i \\eta}{R_s} \\right) = 0\\,, \\end{equation} whose solution is given by \\begin{equation} \\omega (\\eta) = -\\frac{3}{2} \\frac{V_s}{R_s} + \\sqrt{\\frac{3}{2} \\frac{V_s ^2}{R_s ^2} + 2 \\pi G \\sigma _0 \\frac{\\eta}{R_s} - c_s^2 \\frac{\\eta ^2}{R_s ^2}}. \\label{eq:omega_eta} \\end{equation}  Let us consider the {\\sl first growing mode}.  This one corresponds to the sequence of values of $\\eta$ which maximises the angular frequency $\\omega$ at each moment of the shell propagation, i.e. $\\eta _{fg} = \\eta (t)$ such that \\begin{equation} \\frac{d \\omega}{d \\eta} = 0\\;. \\label{eq:fg_def} \\end{equation} Equation \\ref{eq:fg_def} indeed corresponds to a maximum since the discriminant of Eq.~\\ref{eq:omega_eta} shows a negative curvature with $\\eta$. The instantaneous angular wavenumber and the instantaneous angular frequency associated to the first growing mode obey respectively \\begin{equation} \\eta _{fg} = \\frac{\\pi G}{{c_s}^2} \\sigma _0 R_s = \\frac{1}{4 {c_s}^2} \\frac{GM}{R_s} \\label{eq:fg_eta} \\end{equation}  and  \\begin{equation} \\omega _{fg} = - \\frac{3}{2} \\frac{V_s}{R_s} + \\sqrt{\\frac{V_s^2}{R_s^2} + \\frac{\\pi ^2 G^2 \\sigma _0^2 }{c_s^2} }\\;. \\label{eq:fg_omega} \\end{equation}  Taking into account the dependence of $M$ and $R_s$ (Eq.~\\ref{eq:RsHPBt1/3}) on $P_h$ and $N$, Eq.~\\ref{eq:fg_eta} shows that $\\eta _{fg}$ depends on the external pressure $P_h$, on the SN number $N$ and on time $t$ as \\begin{equation} \\eta _{fg} \\propto P_h ^{-1/6} N^{-1/3} t^{-1/3}\\;. \\label{eq:eta_fg/NPh} \\end{equation}  A first guess of a favourable angular wavenumber can be estimated from the temporal average of Eq.~\\ref{eq:fg_eta} over the time spent by the supershell in the hot protogalactic background: \\begin{equation} <\\eta _{fg}> = \\frac{1}{\\Delta t - t_{em}} \\int_{t_{em}}^{\\Delta t} \\eta _{fg}(t') ~ dt' \\,, \\label{eta_fg_ave} \\end{equation}  where $\\Delta t$ is the duration of the SN phase.  Among the cases displayed in Fig.~\\ref{fig:disc_eta} ($c_s$=1\\,km.s$^{-1}$, $\\tilde \\sigma _1(t_{em})=0.01\\sigma _0(t_{em})$, $v(t_{em})=0.01\\,V_s(t_{em})$), we see that the shell transverse collapse is achieved if, for instance, $P_h = 10^{-10}$\\,dyne.cm$^{-2}$, $N$=100 and $\\eta$=16 (top panel) or if $P_h = 5 \\times 10^{-10}$\\,dyne.cm$^{-2}$, $N$=100 and $\\eta$=10 (middle panel).  It is interesting to note that these values of $\\eta$ reasonably match those given by Eq.~\\ref{eta_fg_ave}, for the above mentioned cases, i.e., $<\\eta _{fg}>$=12 and $<\\eta _{fg}>$=9, respectively.  As Eq.~\\ref{eta_fg_ave} has been derived under the assumption of an exponential growth rate, this agreement a posteriori justifies its validity as a convenient estimate of the number of forming clumps in collapsing shells.  \\\\  \\subsubsection{The shell sound speed: $c_s$} \\label{sub:disc_cs}  \\begin{figure} \\begin{center} \\epsfig{figure=disc_cs.eps, width=\\linewidth} \\label{fig:disc_cs} \\caption{Temporal evolution of $\\sigma _0$ (plain curve) and $\\tilde \\sigma _1$ (dashed curves), expressed in units of 1\\,M$_{\\odot}$.pc$^{-2}$, considering the sound speeds (see the key), external pressures and SNII numbers indicated.  Other parameters are $\\eta$=10,  $\\tilde \\sigma _1(t_{em})=0.01\\sigma _0(t_{em})$, $v(t_{em})=0.01\\,V_s(t_{em})$} \\end{center} \\end{figure}  The development of a gravitational instability within a supershell also depends on the sound speed $c_s$ of the shell material. Indeed, $c_s$ is directly related to the thermal pressure $P_s$ of the shell gas through \\begin{equation} P_s = \\frac{\\rho _s k T_s}{\\mu _s m_H} = \\rho _s c_s ^2\\,. \\end{equation} In this equation, $k$ and $m_H$ are the Boltzmann constant and the hydrogen mass, respectively, while $\\rho _s$, $T_s$ and $\\mu _s$ are the mass density, the temperature and the mean molecular weight of the shell, respectively. The larger the velocity dispersion, the more the layer resists gravitational collapse. In an H\\,I layer (T $\\simeq$ 100\\,K, $\\mu ~\\simeq$ 1.3\\,m$_H$), the sound speed is $\\simeq$ 0.8\\,km.s$^{-1}$, while it can be as low as $\\simeq$ 0.3\\,km.s$^{-1}$ in an H$_2$ layer (T $\\simeq$ 20\\,K, $\\mu ~\\simeq$ 2.1\\,m$_H$) (McCray \\& Kafatos 1987).  However, the turbulence and magnetic fields of the shell will increase these values. To some extent, they can be represented by an additional pressure term in the expression of $c_s$, i.e. \\begin{equation} c_s = \\left ( \\frac{k T_s}{\\mu _s m_H} + \\frac{B_s ^2}{4 \\pi \\rho _s} + \\mathcal{T} {\\rm _s} \\right)^{1/2} \\label{eq:cs_mag} \\end{equation} where $B_s$ is the magnetic field in the shell and $\\mathcal{T} {\\rm _s}$ is the contribution of turbulence.   \\\\  Figure 7 shows how highly sensible to $c_s$ the fragmentation process is. Considering 100\\,SNeII and a hot background pressure of 5 $\\times$ 10$^{-10}$\\,dyne.cm$^{-2}$, the fragmentation takes place even with a low initial transverse velocity, i.e. $v(t_{em})=0.01 V_s(t_{em})$, if $c_s$=1\\,km.s$^{-1}$.  Increasing the latter by 25\\%, the transverse collapse is very weakened and the fragmentation is prevented.  During the last ten million years, the evolutions with time of $\\sigma _0$ and $\\tilde \\sigma _1$ are similar, that is, the evolution of $\\tilde \\sigma _1 (t)$ is mostly driven by the dilution of $\\sigma _0 (t)$ due to the shell expansion. \\\\  Equation \\ref{eq:cs_mag} illustrates the difficulty of estimating the sound speed of a gas.  It implies the computations of its cooling history, its magnetic fields and turbulence. Moreover, the high sensibility of the fragmentation issue to $c_s$ (see Fig.~7) shows that its value must be estimated with some accuracy.  Such a task is well beyond the scope of the present work and, in what follows, we adopt $c_s=1\\,{\\rm km.s}^{-1}$, in agreement with many studies of supershell fragmentation (e.g. Comeron \\& Torra 1994, Ehlerova \\& Palous 2002). ",
        "conclusions": "\\label{sec:search_para}  \\subsection{Stars with halo GC metallicities} \\label{sec:disc_GC}   The previous section has shown that the shell/swept PGCC may become gravitationally unstable and finally break into fragments providing that some conditions are fulfilled, e.g. $v(t_{em})/V_s(t_{em})  \\simeq $ 0.03, $c_s \\simeq $ 1\\,km.s$^{-1}$, $\\eta \\simeq <\\eta _{fg}>$.  This is {\\sl not} to claim that all supershells will encounter such favourable circumstances, but one may expect that at least {\\sl some} of them will do. \\\\  \\begin{figure} \\begin{minipage}[b]{\\linewidth} \\begin{center} \\epsfig{figure=frag_v1.eps, width=\\linewidth} \\end{center} \\end{minipage} \\vfill \\vspace*{-9mm} \\begin{minipage}[b]{\\linewidth} \\begin{center} \\epsfig{figure=frag_v2.eps, width=\\linewidth} \\end{center} \\end{minipage} \\vfill \\vspace*{-9mm} \\begin{minipage}[b]{\\linewidth} \\begin{center} \\epsfig{figure=frag_v3.eps, width=\\linewidth} \\end{center} \\end{minipage} \\label{fig:frag_prob} \\caption{Summary of the ($N$, $P_h$, $v(t_{em})/V_s(t_{em})$, $\\eta$) values which may (open or filled circles), or may not (no symbol), lead to a swept PGCC fragmentation providing that the transverse flows converge towards the clumps embedded within the shell, i.e. $\\Delta \\phi = -\\pi /2$.  The sound speed and the initial perturbed surface density are $c_s = 1$\\,km.s$^{-1}$ and $\\sigma _1 (t_{em}) = 0.01 ~\\sigma _0 (t_{em})$, respectively. The top, intermediate and bottom panels correspond to initial transverse velocity of 1, 2 and 3 per cent of the shell velocity when it emerges out of the cloud, respectively.  The isometallicity curves corresponding to [Fe/H]=$-$1.2, $-$1.5, $-$2, $-$2.5 are also plotted, giving thus the metallicity achieved through self-enrichment for each parameter combination.  For each ($N$,$P_h$) pair, 54 couples of initial perturbed velocities ($v(t_{em})/V_s(t_{em}$)= 0.01, 0.02, 0.03) and clump numbers ($\\eta$ ranging from 6 to 40 by step of 2) are tested.    In each panel (i.e., for each value of the initial perturbed velocity), a couple ($N$, $P_h$) leading to a successful transverse collapse is marked by an open/filled circle as well as by the range of $\\eta$ values leading to fragmentation.  The larger the range of $\\eta$ values, the bigger the symbol (see text for details)} \\end{figure}  Figure 8 presents the results of shell fragmentation simulations for varying values of $N$, $P_h$, $\\eta$ and $v(t_{em})/V_s(t_{em})$, assuming that $c_s$=1\\,km.s$^{-1}$ and $\\sigma _1 (t_{em}) $= 0.01 $\\sigma _0(t_{em})$. Five hot protogalactic background pressures ($P_h=10^{-11}, 3.2 \\times 10^{-11}, 10^{-10}, 3.2 \\times 10^{-10}, 10^{-9}$\\,dyne.cm$^{-2}$) and 3 SN numbers ($N$=50, 100, 200) are tested.  The upper limit for $N$ is the maximum number of supernovae that the GC gaseous progenitor can sustain, namely $N$=200 (i.e., if $N > 200$, the absolute value of the cloud binding energy is lower than the shell kinetic energy: disruption criterion, Paper I).  The lower limit is imposed by $c^{PGCC}$, the sound speed of the PGCC material.  Indeed, the shell is ``built'' while sweeping the PGCC and such a mass accumulation into a shell requires the velocity of the shell to be larger than the sound speed of the ambiant medium. Therefore, the lower limit to the shell velocity in the PGCC obeys: \\begin{equation} {V_s}^{PGCC} = {c}^{PGCC} = \\sqrt{\\frac{k T}{\\mu m_H}} = 8.3\\,{\\rm km.s^{-1}} \\end{equation} where $T$ and $\\mu$ are the temperature ($\\simeq 10^4$\\,K) and the mean molocular weight ($\\simeq 1.2$) of the PGCC, respectively (Fall \\& Rees 1985). The number of SNeII corresponding to this lower limit of the shell velocity is $N \\simeq 50$ (Eq.~13, Paper I). Regarding the upper value of $P_h$, we refer to Murray \\& Lin (1992) who showed that the hot protogalactic background pressure depends on the galactocentric distance $D$ as \\begin{equation} P_{h} = 1.25\\times 10^{-9} D_{kpc}^{-2}  \\,\\, \\rm dyne\\,cm^{-2}. \\label{Ph2} \\end{equation} Thus, $P_h$ is of order 10$^{-9}$\\,dyne.cm$^{-2}$ in the very inner Galactic regions (i.e., $D \\simeq$ 1\\,kpc).  \\\\ In order to see what is the metallicity achieved in the shells which succeed in forming new stars, iso-metallicity curves corresponding to [Fe/H]=$-$1.2, $-$1.5, $-$2 and $-$2.5 (i.e., metallicities typical of the Galactic halo GCs) are displayed in each panel of Fig.~8. \\\\ For each couple ($N$,$P_h$), 54 combinations of $v(t_{em})/V_s(t_{em})$ and $\\eta$ (i.e. 3 initial perturbed velocities $\\times$ 18 numbers of clumps) have been run.  The initial transverse velocities correspond to 1, 2 and 3 per cent of the velocity of the shell when it enters the hot background and the number of clumps ranges from 6 to 40 by step of two.  It appears that the shell is unable to fragment if $\\eta \\ge 40$.  When a set of conditions leads to a successful fragmentation, the corresponding point in the ($N$, $P_h$) diagram is marked by a circle  as well as by the range of $\\eta$ values leading to successful transverse collapses.  Depending on whether the shell fragmentation takes place for 0 to 25 per cent, 25 to 50 per cent, 50 to 75 per cent, or 75 to 100 per cent of the number of $\\eta$ values tested, the corresponding triplet ($N$, $P_h$, $v(t_{em})/V_s(t_{em}$)) is marked by an open circle, a small, a median-size or a large filled circle, respectively. Figure 8 confirms that the probability of successful transverse collapse fades away with decreasing external pressure and increasing number of SNeII. As mentioned in Sect.~\\ref{sub:discNPh}, a low external pressure and a high number of SNeII leads to a larger radius for the shell, decreasing thereby its surface density (Eq.~\\ref{eq:sig0_N_Ph}) and its ability to collapse.   \\\\  At high pressure, i.e. $P_h \\simeq 10^{-9}$\\,dyne.cm$^{-2}$, the ability of the shell to get fragmented is limited by too large a number of SNeII. Figure 8 shows however that a SN number as large as 200 does not prevent the fragmentation.  Accordingly, the largest metallicity which can be achieved through self-enrichment is [Fe/H]$\\simeq -$1.2. On the other hand, at low pressure, the transverse collapse is supported by a low number of SNeII.  Combining a low background pressure with the lower limit on $N$, Fig.~8 shows that the lowest metallicity which can be achieved is [Fe/H]$\\simeq -$2.8 providing that the relative initial transverse velocity is 3 per cent. While this lower limit in metallicity is a bit uncertain, as it is achieved in the case of the highest intial transverse velocity only, the examination of the three panels in Fig.~8 clearly shows that a metallicity of $\\simeq -2.5$ is actually achievable.  These extreme values ([Fe/H]$\\simeq -$1.2 and $-$2.5) match nicely the metallicities exhibited by the most metal-rich and the most metal-poor Galactic halo GCs, respectively. \\\\  At this stage, it is worth keeping in mind that the question addressed here above concerns the ability of the shell to form stars and not yet the ability of these stars to evolve into a stellar cluster. We will address the ability of the newly formed stars to form a bound cluster in a forthcoming paper and we emphasize here that, among the cases of successful fragmentation displayed by Fig.~8, some of the newly formed stars may not be able to form a bound cluster.  In other words,  the collapsed shells may be a source for both halo GCs and halo field stars.  Section 2.3.2 has discussed how the self-enrichment model for GC formation can shed light on the metallicity distribution of the Galactic halo.  Because the parameters which determine the metallicity, i.e. $N$ and $P_h$, also control the shell surface density, there is a direct link between the metallicity and the probability of star formation\\footnote{This is {\\sl not} to say that the probability of star formation is determined by the metallicity, as it can be assumed in pre-enrichment models in which the metallicity may control the star formation through line cooling processes.  It rather means that the metallicity and the probability of star formation derive from the same parameters, namely the number of SNeII and the pressure of the hot protogalactic background}.  Furthermore, the initial radius and velocity of the second generation stars being the radius and the velocity of the shell at the time of their formation, their binding depends on the shell expansion law and, therefore, again, on $N$ and $P_h$.  As a consequence, the probability of getting a bound cluster from a shell of newly formed stars is also related to their metallicity.  If we assume that a significant fraction of the field content comes from ``failed GCs'' (i.e. shells of stars which did not succeed in forming GCs), then these relations between the metallicity on the one hand and the probabilities of forming a second stellar generation and a bound cluster on the other hand {\\sl open the way to the computation of the metallicity distributions of both halo field stars and halo GCs}.  Since the formation of a GC requires one more condition than the shell transverse collapse (i.e. the binding of the second stellar generation), only a subset of the ($N$,$P_h$) combinations giving rise to shell fragmentation may lead to the formation of a bound cluster.  Such a conclusion does not contradict, and even fits, the fact that the metallicity spectrum of halo field stars is larger than the one of GCs (e.g. Laird et al.~1988).  \\\\  \\subsection{The very metal-poor stars ([Fe/H] $\\lesssim -2.5$)} \\label{sub:vmp_st}  The most striking difference between the halo GC and halo field metallicity distributions resides in the much more extended metal-poor tail of the field compared to the GC system.  In fact, while the most metal-poor GCs show [Fe/H]$\\simeq -$2.5, the Galactic halo hosts field stars whose metallicity is much lower (e.g., CS22876-032, [Fe/H]$\\simeq -$3.7, Norris, Beers \\& Ryan 2000).  The recent discovery of HE0107-5240, the most metal-deficient star ever discovered ([Fe/H]$\\simeq -$5.3, Christlieb et al.~2002) has decreased even more the lower limit of the observed metallicity distribution of halo field stars.  Our model does not seem to be able to explain these very metal-poor stars (see Fig.~8) and we must therefore investigate alternative scenarios for the formation of these stars more metal-deficient than the most metal-poor GCs.  \\subsubsection{Isolated Type II Supernovae} The existence of very metal-poor stars has often been attributed to star formation episodes triggered by {\\it individual} SNeII exploding in primordial gas clouds (e.g., Shigeyama \\& Tsujimoto 1998, Argast et al. 2000, Karlsson \\& Gustafsson 2001), that is, a formation scenario similar to the self-enrichment scenario for GCs except regarding the considered number of massive stars. The metallicity of such stars is determined by the ratio of the mass of metals ejected by the SNII to the mass of hydrogen gathered by the shock wave.  The SNII yields computed by Woosley \\& Weaver (1995) for zero metallicity stars show that a SNII with progenitor mass $m$ ejects a mass of metals $m_z$ such that $m_z \\simeq 0.3m - 3.5$ (in unit of one solar mass).  On the other hand, the mass $M_{sw}$ of interstellar gas collected by the shock wave is related to the explosion energy $E_0$ through \\begin{equation} M_{sw} = 5.1 \\times 10^4 M_{\\odot} \\left( \\frac{E_0}{10^{51}\\,ergs} \\right)^{0.97}\\,, \\label{single_SN} \\end{equation} assuming a sound speed of 10\\,km.s$^{-1}$ (or T $\\sim$ 10$^4$\\,K) for the interstellar gas (Shigeyama \\& Tsujimoto 1998).  Using these relations and assuming $E_0 = 10^{51}$\\,ergs for a canonical SNII, the metallicity achieved in shells of gas driven by isolated SNeII is straightforwardly derived (see the second column of Table 1). The explosion energy being the same whatever the SNII progenitor mass, every SNII sweeps the same amount of primordial circumstellar gas, irrespective of the SNII mass (Eq.~\\ref{single_SN}).  Therefore, the final metallicity is determined solely by the mass of heavy elements released in the interstellar medium by the exploding star and, thus, by the SNII progenitor mass.  As a consequence, the more massive the SNII, the larger the final metallicity. As already noticed by Shigeyama \\& Tsujimoto (1998), the formation of very metal-poor stars with [Fe/H] $\\simeq -4$ and [Fe/H] $\\simeq -2$ may therefore be ascribed to ``low''-mass (i.e., M $\\simeq$ 12\\,M$_{\\odot}$) and high-mass (i.e., M $\\simeq$ 40\\,M$_{\\odot}$) SNII, respectively. In order to dig out the potential link between very metal-poor stars and isolated SNeII, Shigeyama \\& Tsujimoto (1998) also exploited the yields of core-collapse SN predicted by nucleosynthesis calculations (e.g., Woosley \\& Weaver 1995, Tsujimoto et al.~1995). They focused on the abundance patterns of C, Mg, Si and Ca. The yields of these elements being not (or at least less) affected by the mass-cut issue, they should therefore be known with a better accuracy than the yields of heavier elements such as the iron peak ones. They noticed the good agreement between the abundance patterns arising in the remnants of first generation SNeII, as predicted by nucleosynthesis calculations, and those inferred from the spectra of halo stars with [Fe/H] $\\lesssim -$2.5.  As a result, Shigeyama \\& Tsujimoto (1998) suggested that very metal-poor stars actually formed out of interstellar gas which had been swept by isolated SNeII.  They also note that the large observational scatter in the abundance ratios among these stars can be explained by the differences in SNII yields with the mass of the progenitor.  It is worthy of a mention that such isolated SNeII are certainly not able to trigger the formation of GCs owing to the small amount of interstellar gas swept by the blast wave (Eq.~\\ref{single_SN}; also, not all the gas will be converted into stars).  Should such low-mass halo clusters have managed to form however, they will have been quickly destroyed by the Galactic tidal fields, low-mass clusters being among the most vulnerable in this respect (e.g., Gnedin \\& Ostriker 1997). Therefore, this extension of the self-enrichement scenario to single SNII explosions does not contradict the absence of Galactic halo GCs with [Fe/H] $\\lesssim -$2.5.  \\subsubsection{Hypernovae} In fact, even stars as massive as 40\\,M$_{\\odot}$ may trigger the formation of stars with a metallicity as low as $-4$, providing that they explode as {\\it hypernovae}, i.e., supernovae characterized by explosion energies of order $E_0 \\sim 10^{52} - 10^{53}$~ergs.  This class of objects includes two categories depending on the progenitor mass, namely core-collapse hypernovae and pair-instability hypernovae. They both have been invoked as possible explanations to peculiar abundance patterns observed in some very metal-poor field stars.  \\\\  {\\it Core-collapse hypernovae.  } Recent observations suggest that at least some core-collapse supernovae explode with explosion energies ten to one hundred times higher than the energy released by a canonical SNII (e.g., Galama et al.~1998).  They likely originate from relatively massive star (M$\\gtrsim 25$~M$_{\\odot}$). In contrast to pair-instability hypernovae (see below), the mass range of their progenitor is thus similar to the one of canonical SNeII.  If hypernovae occured in the early stage of the Galactic evolution and induced star formation, their abundance pattern may still be observable in the atmospheres of some low-mass halo stars. As for the case of star formation triggered by single SNeII, we can derive an estimate of the metallicity of these halo stars using Eq.~\\ref{single_SN} and hypernova yields.  Nakamura et al. (2001) have investigated in detail the nucleosynthesis of core-collapse hypernovae and compared it with the yields of canonical (i.e., $E_0 \\sim 10^{51}$\\,ergs) core-collapse SNeII with similar progenitor mass. Their Tables 2-5 show that SNeII and hypernovae with similar progenitor mass release much the same amount of metals.  The resulting metallicity of stars formed in hypernova remnants are given in Table 1, following the same of line of reasoning as in the previous section. Hypernovae being much more energetic than SNeII, and the mass of interstellar gas swept by the blast wave being roughly proportional to the explosion energy (Eq.~\\ref{single_SN}), they will collect a larger amount of interstellar gas.  Therefore, should core-collapse hypernovae be able to trigger the formation of new stars, these stars will be more metal-poor than the ones formed in the remnants of canonical SNeII.   Thus, they may be well-suited to explain the formation of stars with a metallicity as low as [Fe/H] $\\simeq -4$ (see Table 1).  \\begin{table} \\caption[]{Dependence of the metallicity [Fe/H] on the progenitor mass $m$ and explosion energy E$_0$ of supernovae (SN; E$_0 = 10^{51}$\\,ergs) and hypernovae (HN; E$_0 = 10 ~~ {\\rm or} ~~ 100 \\times 10^{51}$\\,ergs). Every mass is expressed in units of one solar mass. (1) A 12\\,M$_{\\odot}$ star explodes as a canonical SNII. According to Nakamura et al.~(2001), stars less massive than 25\\,M$_{\\odot}$ do not explode as core-collapse hypernovae and such cases are thus not considered in this Table. (2-3) More massive stars explode either as core-collapse supernovae or core-collapse hypernovae. (4) Very massive objects (i.e., M $>$ 100\\,M$_{\\odot}$) explode as pair-instability hypernovae.  M$_{sw}$, the mass of interstellar gas collected by the explosion blast wave, depends on the explosion energy through Eq.~\\ref{single_SN}.  m$_z$ is the mass of metals released by an exploding massive star with progenitor mass $m$. The metallicity [Fe/H] is derived assuming that the mass m$_z$ of metals is mixed with the mass M$_{sw}$ of collected interstellar gas} \\begin{center} \\begin{tabular}{ c c c c } \\hline E$_{0}/10^{51}$\\,ergs & 1 (SN) & 10 (HN) & 100 (HN)  \\\\ \\hline M$_{sw}$  & 5.1 $\\times$ 10$^4$ & 5.1 $\\times$ 10$^5$ & 5.1 $\\times$ 10$^6$ \\\\ \\hline \\hline (1)  (m,m$_z$)=(12,0.1)~  &   &   &    \\\\ {\\rm [Fe/H]} $\\simeq$ & $-$4 & -- & --  \\\\ \\hline (2)  (m,m$_z$)=(25,4)~~~ &  &   &    \\\\ {\\rm [Fe/H]} $\\simeq$ & $-$2.4 & $-$3.4 & $-$4.4  \\\\ \\hline (3)  (m,m$_z$)=(40,10)~~~ & &   &        \\\\ {\\rm [Fe/H]} $\\simeq$  & $-$2.1 & $-$3.1 & $-$4.1  \\\\ \\hline (4)  (m,m$_z$)=(200,100) & &   &        \\\\ {\\rm [Fe/H]} $\\simeq$  & -- & $-$2.0 & $-$3.0  \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  The most significant feature of hypernova nucleosynthesis is their iron production, this one being larger than in SNeII by a factor 2 to 10 (Nakamura et al.~2001).  This leads to small abundance ratios of $\\alpha$ elements over iron.  Nakamura et al. (2001) thus suggested that the gas out of which the very metal-poor binary CS 22873-139 ([Fe/H] = $-$3.4) formed had been contaminated by the ejecta of an hypernova as this halo star shows almost solar [Mg/Fe] and [Ca/Fe] ratios.  On the other hand, while stars with $-2.5 \\lesssim {\\rm [Fe/H]} \\lesssim 0$ show [Zn/Fe] $\\simeq$ 0 (e.g., Primas et al.~2000), this abundance ratio is steadily increasing towards [Zn/Fe] $\\sim$ 0.5 as the metallicity decreases, for stars more metal-poor than [Fe/H] $\\sim -2.5$ (e.g., Cayrel et al.~2003). Umeda \\& Nomoto (2002) notice that such a large [Zn/Fe] ratio arises naturally in their own hypernova model and thus conclude that core-collapse hypernovae are likely to have contributed to the early Galactic chemical evolution.    \\\\  {\\it Pair instability hypernovae.  } These very massive ($\\simeq$ 140-260\\,M$_{\\odot}$) objects exploding with $E \\gtrsim 10^{52}$\\,ergs, they are also called hypernovae. Their name (i.e., ``pair instability'') refers to the electron-positron pair instability process encountered during the central oxygen-burning stages (see Umeda \\& Nomoto 2002, their appendix for a detailed description).  The main feature of these very massive stars is their ability to ``pass carbon and oxygen from the helium-burning core through the hydrogen-burning shell, in such a way that it is CNO processed to nitrogen before entering the hydrogen envelope'' (Carr, Bond \\& Arnett 1984).  These stars thus produce large supersolar values of N/Fe, an effect not predicted by models of Galactic chemical enrichment based on stars less massive than 100\\,M$_{\\odot}$.  Following their discovery of CS 22949-037, a very metal-poor star ([Fe/H] = $-$3.8) showing extreme nitrogen enhancement ([N/Fe] = 2.3), Norris et al. (2002) suggested that this star may have been formed out of gas polluted and compressed by such a very massive hypernova.  On the other hand, in marked contrast with core-collapse hypernova yields, Umeda \\& Nomoto (2002) quote that pair-instability hypernovae are unlikely to produce [Zn/Fe] ratios as large as in very metal-poor stars where [Zn/Fe] ranges from solar up to $\\simeq$0.5 (Cayrel et al.~2003).  Actually,  the abundance ratio derived from their hypernova yields is of order [Zn/Fe]=$-$1.5, that is, 1 to 2 dex smaller than in very metal-poor stars. Additionally, Heger \\& Woosley (2002) note the absence of the r-process in these massive objects, which conflicts with observations showing appreciable amounts of r-process elements in very metal-poor stars (e.g., Burris et al.~2000).  Based on these two arguments, i.e., the low abundance ratio [Zn/Fe] as well as the absence of r-process elements in pair instability hypernova ejecta, Umeda \\& Nomoto (2002) and Heger \\& Woosley (2002) conclude that the abundances of very metal-poor stars cannot be ascribed to pair instability hypernovae only and must include the contribution of an additional nucleosynthetic component, namely those lower mass stars that make ``regular'' supernovae. We also note that pair-instability hypernovae release fairly large amounts of metals in the interstellar medium, of order 100~M$_{\\odot}$ (Umeda \\& Nomoto 2002, their Tables 15-18).  In fact, these stars disrupt completely when exploding, leaving no compact remnant after the explosion (i.e., no issue of ``mass-cut'' or ``fall-back'', see Heger \\& Woosley 2002). The mixing of so large an amount of metals with a mass of primordial gas of $\\simeq 5 \\times 10^5 {\\rm M}_{\\odot}$ (E$_{0} = 10 \\times 10^{51}$\\,ergs) or $\\simeq 5 \\times 10^6 {\\rm M}_{\\odot}$ (E$_{0} = 100 \\times 10^{51}$\\,ergs) will lead to metallicities of order [Fe/H] $\\simeq -2$ or [Fe/H] $\\simeq -3$, respectively. Therefore, the metallicity of the most metal-poor stars in the Galactic halo cannot be explained by pair-instability hypernovae. \\\\  \\subsubsection{Single short-acting source of energy and triggered star formation} The issue of whether isolated exploding massive or very massive stars can stimulate the formation of new stars in the layers of interstellar gas they have swept is not explicitely addressed in the papers mentioned in Sect.~3.2.1 and 3.2.2 (e.g., Shigeyama \\& Tsujimoto 1998, Argast et al.~2000, Norris et al.~2002). It is merely assumed that such stimulated star formation actually took place during the early Galactic stages.  The model presented in this paper (see sect.~2) is not the most appropriate to address this issue quantitatively as it deals with a continuous input energy and not with a single short-acting source of energy.  Actually, supershells created in connection to single explosion and those created around a cluster of massive stars do not show the same expansion rate with time (see below).  We have seen in Sect.~\\ref{sec:stimSF} how the radius and the velocity of the shell affect  its ability to collapse through the shell surface density and the stretching of the forming perturbation, respectively.  Hence, to solve the issue of whether a single supernova/hypernova triggers the formation of new stars requires to derive a new temporal evolution of the shell radius, in the PGCC and in the protogalactic background, adequate to single explosion.  This is beyond the scope of the present paper whose main goal is to address the formation of stars with $-2.5 \\lesssim {\\rm [Fe/H]} \\lesssim -1$ (i.e., the metallicity range of halo GCs).  Some qualitative inferences can however be made. Efremov, Ehlerova \\& Palous (1999) compared the propagation in an homogeneous medium of both kinds of supershells.  In case of an abrupt energy input, the early expansion of the supershell proceeds at a higher velocity than in the case of a continuous input of the same energy amount. Accordingly, in the former case, the shell of swept gas is more strongly stretched and this hampers the formation of an initial perturbation.  After a few millions years however, the shell created in connection to an abrupt energy input slows down and shows both a radius and a velocity lower than in case of a continuous energy input.  In contrast to the initial shell expansion, this second stage favours the transverse collapse through a larger shell surface density and a weaker stretching of the shell perturbed regions. Whether the collapse will occur is actually not certain.  The source of the energy input being point-like (i.e., its size is much smaller than the one of a cluster of massive stars), the initial amplitude of the transverse motions within the shell will be lowered accordingly, thus reducing the collapse.  Comparing SNeII and hypernovae, hypernovae sweep a larger amount of gas thanks to their larger energy input (Eq.~\\ref{single_SN}).  However, this also leads to a larger shell radius.  Both effects acting on the shell surface density in opposite ways, only detail computations will tell us which class of objects is the most efficient in forming new stars.  A $E_0 \\sim 2 \\times 10^{52}$\\,ergs hypernova being able to collect amount of gas as large as 10$^6$\\,M$_{\\odot}$ (Eq.~$\\ref{single_SN}$), one might think that pair-instability hypernovae could be related to the formation of GCs, at least the most metal-poor ones as such hypernovae can chemically enrich the initially pristine gas up to [Fe/H] $\\simeq$ -2 (see Table 1). However, the [Zn/Fe] abundance ratio predicted for hypernova ejecta (i.e., [Zn/Fe] $<-1.5$, Umeda \\& Nomoto 2002) shows a sharp discrepancy with the roughly solar value observed in halo field stars with [Fe/H] $\\simeq -2$ (Primas et al.~2000) as well as in NGC6397, a halo GC with the same metallicity (Thevenin et al.~2001).  Therefore, the abundance patterns of stars in this metallicity range cannot be ascribed to pair-instability hypernovae.  Obviously, it is not straightforward to conclude whether single massive star explosions are able to trigger the formation of new stars.  More computations are required to answer a question which is certainly worthy of further investigation owing to its potential link with very metal-poor stars.  \\subsubsection{External pollution onto PopIII stars} Shigeyama, Tsujimoto \\& Yoshii (2003) have recently suggested that some very metal-deficient stars may have been born as metal-free (PopIII) stars whose external layers were afterwards polluted by the accretion of chemically enriched interstellar gas. The metallicity achieved by the external layers will depend on the metallicity and mass of the accreted gas, as well as on the stellar mass fraction with which the accreted gas has been mixed.  As a consequence, the chemical enrichment of stellar superficial layers will show up more markedly in main sequence stars than in red giants, the convective envelope (i.e. the mixing zone) being almost two orders of magnitude larger in halo giants (about half the star mass) than in non-evolved metal-poor stars (about one per cent of the star mass).  Shigeyama et al.~(2003) have shown that the extremely low metallicity ([Fe/H]$\\simeq -$5.3) of the giant HE0107-5240 (Christlieb et al.~2002) may originate from such a mechanism, that is, an external pollution while the star was on the main sequence followed by the dilution of the accreted material as the star started ascending the red giant branch.  Owing to their thin convective envelopes, polluted Pop III stars still on the main sequence would exhibit larger metallicities, i.e. of order [Fe/H]$\\simeq -$3, although still lower than the most metal-poor GCs.   \\\\  \\subsubsection{Star formation in dwarf galaxies} Within the frame of a hierarchical model for halo formation, Cote et al.~(1999) proposed that the excess of very metal-poor field stars with respect to GCs was formed in the most metal-deficient dwarf galaxies trapped within the Galactic potential well.  These galaxies are not expected to contribute to the GC population. Indeed, due to the luminosity-metallicity correlation observed among dwarf galaxies (Gilmore 2000, Mateo 2000), the most metal-poor dwarfs are also expected to be the least luminous.  Dwarf galaxies exhibit a luminosity threshold for hosting a GC system.  For instance, the Fornax dwarf spheroidal galaxy is the faintest ($M_v \\simeq -12.3$, Harris 1991) galaxy known to harbour its own GC system.  Therefore, the dimmest dwarf galaxies, which are also the most metal-poor, may have contributed to the halo field and not to the halo GC system.   This paper has presented the results of simulations dedicated to the transverse collapse of shells of gas resulting from the sweeping of gaseous GC progenitors by SNeII.  In these simulations, the growth of an initial perturbation in the shell surface density is followed by solving the linear perturbed equations of continuity and motion for transverse flows in a spherical shell of gas. Such a collapse depends on several parameters, namely the number of SNeII, the background presure, the sound speed of the shell gas and the initial conditions of the perturbation (i.e. the number of clumps, the initial perturbed surface density and velocity and the corresponding phase difference).  All these parameters have been discussed in turn.  The results show that the pressure $P_h$ of the hot protogalactic background ($P_h \\sim 10^{-10}$\\,dyne.cm$^{-2}$, Fall \\& Rees 1985, Murray \\& Lin 1992) and the numbers $N$ of SNeII allowed by the disruption criterion (i.e. smaller than 200, Paper I) can indeed lead to a successful shell transverse collapse, and thereby to the formation of new stars, assuming some reasonable initial conditions for the perturbation (see Fig.~8).  The metallicities achieved in the shells able to collapse agrees with the metallicity range of Galactic halo GCs, namely, $-2.5 \\lesssim {\\rm [Fe/H]} \\lesssim -1$. Furthermore, while $N$ and $P_h$ determine the metallicity achieved through self-enrichment, they also control the probability of triggered star formation and the ability of these second generation stars to form a bound GC.  Such a property is the most interesting since it opens the way to the understanding of the halo metallicity distribution functions, for both stars and clusters."
    },
    "0403/astro-ph0403649.txt": {
        "abstract": "{ Magnetic field behaviour in a spherically-symmetric accretion flow for parameters typical of single black holes in the Galaxy is discussed. It is shown that in the majority of Galaxy volume, accretion onto single stellar-mass black holes will be spherical and have a low accretion rate ($10^{-6} - 10^{-9}$ of the Eddington rate). An analysis of plasma internal energy growth during the infall is performed. Adiabatic heating of collisionless accretion flow due to magnetic adiabatic invariant conservation is $25\\%$ more efficient than in the standard non-magnetized gas case. It is shown that magnetic field line reconnections in discrete current sheets lead to significant nonthermal electron component formation. In a framework of quasi-diffusion acceleration, the \"energy-radius\" electron distribution is computed and the function describing the shape of synchrotron radiation spectrum is constructed. It is shown that nonthermal electron emission leads to formation of a hard (UV, X-ray, up to gamma), highly variable spectral component in addition to the standard synchrotron optical component first derived by Shvartsman generated by thermal electrons in the magnetic field of accretion flow. For typical interstellar medium parameters, a black hole at 100 pc distance will be a 16-25$^{\\rm m}$ optical source coinciding with the highly variable bright X-ray counterpart, while the variable component of optical emission will be about 18-27$^{\\rm m}$. The typical time scale of the variability is $10^{-4}$ sec, with relative flare amplitudes of 0.2-6\\% in various spectral bands. Possible applications of these results to the problem of search for single black holes are discussed. ",
        "introduction": "\\label{sec_introduction}  Even though more than 60 years have passed since the theoretical prediction of black holes as an astrophysical objects (Oppenheimer \\& Snyder \\cite{oppenheimer}) in some sense they have not been discovered yet. To identify an object as a black hole, one needs to show that its mass exceeds $3M_{\\odot}$, its size is close to $r_g = 2G M/c^2$ and it has an event horizon instead of a normal surface -- the distinguishing property of black holes which separates them from massive compact objects of finite size in some theories of gravity (Will \\cite{will}). However, only the two former criteria are used now for selection of black hole candidates of two types: a) with masses of 5-18 $M_{\\odot}$, in X-ray binaries (see, for example, Greiner et al. \\cite{greiner}); and b) supermassive black holes in galaxy nuclei with masses of $10^6 - 10^{10} M_{\\odot}$ (Shields \\cite{shields_1999}). Existence of the event horizon in such objects is usually implied by the absence of periodic pulsations of the X-ray emission from strong regular magnetic fields (the black hole \"no-hair\" theorem) and I type X-ray flares due to thermonuclear bursts of the accreted matter on the surface of the neutron star. At the same time, typical masses of X-ray pulsars and bursters are close to the typical neutron star value of 1.4 $M_{\\odot}$ while black hole candidates, missing pulsations and X-ray flares, have masses of 5-18 $M_{\\odot}$ (Miller et al. \\cite{miller_1998}). The absence of an event horizon in low-mass objects is not a proof of its existence in higher-mass ones.  High accretion rates in X-ray binaries and active galactic nuclei result in the screening of regions close to the event horizon, and the most luminous parts of accretion flow are situated at distances of $10 - 100 r_g$ (Chakrabarti \\cite{chakrabarti}; Cherepashchuk \\cite{cherepashchuk}) where general relativity effects are negligible.  There is a very effective way to get information about the innermost parts of accretion disks in X-ray binaries as well as AGNs -- the investigation of the broad (and sharp) iron $K_{\\alpha}$ fluorescent emission line (see review by Reynolds \\& Nowak \\cite{reynolds_2003}). Its intensity and shape depend on the accreted plasma distribution and behaviour until the last stable orbit (0.62 $r_g$ for a extremely spinning Kerr black hole and 3 $r_g$ for a Schwarzschild one) (Miller et al. \\cite{miller_2004}; Miniutti et al. \\cite{miniutti}). However since the photons generated at the different distances from the horizon are mixed in the line profile it is not possible to extract the manifestations of gravitational fields close to the horizon only. This may be possible by study of variability of the iron line (Reynolds et al. \\cite{reynolds_2004}).  At the same time, single stellar-mass black holes, which accrete interstellar medium of  low density ($10^{-2} - 1 $cm$^{-3}$), are the ideal case for detection and study of the event horizon. Shvartsman (\\cite{shvartsman_1971}) first demonstrated that an emitting halo of accreted matter forms around such objects and generates optical featureless emission. The majority of such emission comes from the regions near the horizon at $(3 -- 5) r_g$. Spherical accretion onto the single stellar-mass black holes has been studied in detail in the works of several authors (Bisnovatyi-Kogan \\& Ruzmaikin \\cite{bisnovatyi_1974}; Meszaros \\cite{meszaros}; Ipser \\& Price \\cite{ipser_1977, ipser_1982}) and the main conclusions of Shvartsman have been confirmed.  The most striking property of the accretion flow onto the single black hole is its inhomogeneity -- the clots of plasma act as a probe testing the space-time properties near the horizon. The characteristic timescale of emission variability is $\\tau_v\\sim r_g/c\\sim 10^{-4} - 10^{-5}$ sec and such short stochastic variability may be considered as a distinctive property of black hole as the smallest possible physical object with a given mass. Its parameters -- spectra, energy distribution and light curves -- carry important information on space-time properties of the horizon (Beskin \\& Shvartsman \\cite{beskin_1976}).  The general observational appearance of a single stellar-mass black hole at typical interstellar medium densities is the same as other optical objects without spectral lines -- DC-dwarfs and ROCOSes (Radio Objects with Continuous Optical Spectra, a subclass of blazars) (Beskin \\& Mitronova \\cite{beskin_1991}, Pustilnik \\cite{pustilnik_1977}, Shvartsman \\cite{shvartsman_1977}; Beskin et al. \\cite{beskin_2000}). The suggestion that isolated BHs can be among them is the basis of the observational programme of search for isolated stellar-mass black holes -- MANIA (Multichannel Analysis of Nanosecond Intensity Alterations). It uses photometric observations of candidate objects with high time resolution, special hardware and data analysis methods (Shvartsman \\cite{shvartsman_1977}; Beskin et al. \\cite{beskin_1997}).  In observations using the 6-meter telescope of the Special Astrophysical Observatory of 40 DC-dwarfs and ROCOSes, only upper limits for variability levels of 20\\% -- 5\\% on the timescales of $10^{-6}$ -- 10 sec, respectively, were obtained, i.e. BHs were not detected (Shvartsman et al. \\cite{shvartsman_1989a}, \\cite{shvartsman_1989b}, Beskin et al. \\cite{beskin_2000}).  Recently, some evidences appeared that single stellar-mass black holes may be found among the stationary unidentified gamma-ray sources (Gehrels et al. 2002), gravitational lenses causing long-lasting MACHO events (Bennett et al. 2001) and white dwarf -- black hole binaries detected by means of self-microlensing  flashes (Beskin \\& Tuntsov \\cite{beskin_2002b}). In the last case, the mass transfer from the white dwarf is absent and a black hole behaves as a single one.  In the present work we study spherical accretion onto a single stellar-mass black hole at low accretion rates of $10^{8} - 10^{13}$ g/s. This corresponds to the range of interstellar medium densities of $0.002 - 0.1$cm$^{-3}$ and a 10 $M_{\\odot}$ object moving with a velocity of 20--40 km/s (Bondi \\& Hoyle \\cite{bondi_1944}). This is true for about 90\\% of the Galaxy volume (McKee \\& Ostriker \\cite{mckee}).  Spherical accretion with equipartition of energies, i.e. with roughly equal densities of the magnetic and kinetic energies of plasma has been considered in many papers (Shvartsman \\cite{shvartsman_1971}; Bisnovatyi-Kogan \\& Ruzmaikin \\cite{bisnovatyi_1974}; Kowalenko \\& Melia \\cite{kowalenko}; Ipser \\& Price \\cite{ipser_1977,ipser_1982}). The uniqueness of our approach is in taking into account the significantly non-thermal nature of electron energy distribution function (its synchrotron emission determines the appearance of a black hole). It may be roughly considered as a superposition of two components (this approach is known as \"hybrid plasma\", see Coppi (\\cite{coppi}) and references therein) -- thermal electrons and accelerated electron beams, formed in current sheets where magnetic energy is dissipated in a way similar to solar flares (Pustilnik \\cite{pustilnik_1978,pustilnik_1997}). The latter process supports the equipartition of energies. As a result, the emission of the accretion flow consists of a quasi-stationary \"thermal\" part with a wide-band spectrum from infrared to ultraviolet, and a highly variable flaring nonthermal component. Each such flare is generated due to the motion of the accelerated electron beam in the magnetic field. Its light curve carries information on the magnetic and gravitational field structure near the black hole horizon. Nonthermal luminosity reaches several percents of the total luminosity and may even exceeds it at low rates, while its spectrum covers spectral bands from optical to hard X-ray. This result leads to possible modifications of the search strategy.  In \\S 2, the main characteristics of accretion onto single black holes in the Galaxy are discussed.  In \\S 3, the electron distribution function in phase space is built, in \\S 4, the thermal and nonthermal component luminosities are determined, and in  \\S 5 the shape of its spectra is studied.  In \\S 6, the temporal behaviour of single electron beam emission is studied and some conclusions on the variability of accretion flow are made.  In \\S 7, the main results of this work are summarized, and in \\S 9, possible directions of future work are discussed. ",
        "conclusions": "During recent years the number of works dealing with single stellar-mass black holes has significantly increased. Some are purely theoretical (Punsly \\cite{punsly_1998a,punsly_1998b}; Gruzinov \\& Quataert \\cite{gruzinov}; Abramowicz et al. \\cite{abramowicz}) and other provide discussions of their observational detection (Heckler \\& Colb \\cite{heckler}; Fujita et al. \\cite{fujita}; Beskin et al. \\cite{beskin_2000}; Agol \\& Kamionkowski \\cite{agol_2002b}; Chisholm et al. \\cite{chisholm}). The importance of experiments in strong gravitational fields has been recently noted by Damour (\\cite{damour}). Kramer et al. (\\cite{kramer}) discussed the new possibilities of black holes metric study by the investigation of radio pulsar -- BH binary systems with new generation of radio telescopes.  In this work we tried to concretize physical properties of plasma accreted onto the black hole within the classical paradigm of equipartition of Shvartsman (\\cite{shvartsman_1971}). Assuming the discrete nature of the magnetic energy dissipation processes in current sheets allows us to clarify the shape of the synchrotron spectrum of the accretion flow. A hard highly non-stationary nonthermal spectral component appears as an emission of accelerated particles.  The beams accelerated in the current sheets can generate very short flares, providing information about the neighborhood of the event horizon (Fig.\\ref{fig_flare_lightcurve}). On the other hand it is clear from Fig.\\ref{fig_luminosity_radial} that at low accretion rates a significant amount of thermal synchrotron radiation is generated inside $3r_g$ -- this means that the behaviour of this component will reflect the properties of space-time in strong gravitational fields too.  It is clear that the search for a black hole strategy may be modified in accordance with such results. Optical high time resolution studies of X-ray sources may be very important. Single black holes may be contained inside the known stationary gamma sources (Gehrels et al. \\cite{gehrels}) as well as objects causing long microlensing events (Paczynski \\cite{paczynski}). Thus it is very important to  look for X-ray emission as well as for fast optical variability of these objects. Sample observations of the longest microlensing event MACHO 1999-BLG-22 (Bennett et al. \\cite{bennett_2001}), a stellar-mass black hole candidate, have been performed at the Special Astrophysical Observatory of RAS in the framework of the MANIA experiment in 2003-2004  (Beskin et al. \\cite{beskin_2005}).  The best evidence will be provided by the synchronous high time resolution observations in optical and X-ray ranges.  Detection of the event horizon signatures cannot result from statistical studies. A detailed study of each object is needed to detect its specific appearance."
    },
    "0403/astro-ph0403528_arXiv.txt": {
        "abstract": "Annihilating dark matter particles produce roughly as much power in electrons and positrons as in gamma ray photons.  The charged particles lose essentially all of their energy to inverse Compton and synchrotron processes in the galactic environment.  We discuss the diffuse signature of dark matter annihilations in satellites of the Milky Way (which may be optically dark with few or no stars), providing a tail of emission trailing the satellite in its orbit.  Inverse Compton processes provide X-rays and gamma rays, and synchrotron emission at radio wavelengths might be seen.  We discuss the possibility of detecting these signals with current and future observations, in particular EGRET and GLAST for the gamma rays. ",
        "introduction": "It is almost universally accepted that most of the matter in the universe is non-baryonic.  This dark matter is the chief constituent of gravitationally bound objects from dwarf galaxy scales and larger.  Identifying the nature of dark matter is one of the most important problems in astrophysics, cosmology, and particle physics.  Perhaps the best motivated candidate for cold dark matter is the lightest of the so--called neutralinos arising in supersymmetric extensions to the standard model \\cite{jkg96}.  These are the spin-1/2 Majorana fermion counterparts of the neutral gauge and Higgs bosons, and are expected to have masses at the weak scale (of order 100 GeV).  This scale is intriguing as the relic density of a stable particle with this mass and corresponding cross section turns out to be of order the critical density, as we observe the matter density to be today. Any stable particles with weak scale masses could thus naturally account for the dark matter.  In this paper, we will focus on supersymmetry, but our conclusions are fairly generic to dark matter candidates at the TeV scale.  We outline a new signature of annihilating dark matter in satellites of the Milky Way galaxy.  Prospects for detecting high energy photons as dark matter annihilation products, primarily from the $\\pi^0$ decays that are generic to hadronization processes, have been discussed for many years (for a sample see Ref.~\\cite{gammas}).  Necessarily coming with these photons are high energy electrons and positrons from the analogous $\\pi^\\pm$ decay chains.  Charged particles suffer complicated motions in the galactic magnetic field, and furthermore they lose energy to synchrotron and inverse Compton processes. Searching for the synchrotron emission from the galactic center \\cite{syncGC} and from galactic satellites \\cite{blasi} has been discussed previously, though neglecting the diffusion of the charged particles.  We will show that the inverse Compton emission, extended over a large area from the charged particle annihilation products may be observable for some models of particle dark matter and of galactic satellites. ",
        "conclusions": "We have calculated the spectrum and spatial extent of diffuse emission from the charged particle products of dark matter annihilations.  In addition to the synchrotron emission discussed previously, we have studied the inverse Compton radiation, primarily on starlight photons.  We have focused on galactic satellites that are currently within the diffusion zone, namely within a few kpc of the stellar disk.  For satellites moving with typical galactic halo velocities of 300 km s$^{-1}$, the crossing time of the diffusion zone is of the same order as the diffusion time, thus an inherently time-dependent treatment is required.  For annihilation sources, e.g.\\ galactic satellites at typical distances of 10 kpc, the diffuse emission in both inverse Compton and synchrotron extends over roughly 300 square degrees.  We have shown that at least in terms of the number of photons, the diffuse inverse Compton emission might be detectable by GLAST, assuming bright enough annihilation sources.  The spatial extent of the emission makes its detection problematic of course.  GLAST will certainly detect a significant number of point sources in a region of this size.  In a future work we will study in detail the feasibility of separating these signals.  As mentioned previously, these results are fairly generic, and do not depend strongly on the particle physics model.  As we are concerned with the electrons and positrons, we do not even require that the dark matter have hadronic interactions.  Leptonically interacting dark matter \\cite{krauss,baltzbergstrom} would still provide photons and electrons, albeit by different processes and with different spectra.  Such photon sources would be even harder to reconcile with the EGRET point sources, but annihilation sources below the EGRET detection limit may be detectable by GLAST in any case."
    },
    "0403/astro-ph0403113.txt": {
        "abstract": "We have surveyed $\\sim 400 $ known large-amplitude variables within 15$'$ of the galactic center in the SiO $J=1$--0 $v=$ 1 and 2 maser lines at 43 GHz, % resulting in 179 detections.  SiO lines were also detected from 16 other resulting in 180 detections.  SiO lines were also detected from 16 other sources, which are located within 20$''$ (the telescope half beamwidth) of the program objects.  The detection rate of 48 percent is comparable to that obtained in Bulge IRAS source surveys. Among the SiO detections, five stars have radial velocities greater than 200 km s$^{-1}$. The SiO detection rate increases steeply with the period of light variation, particularly for stars with $P>500$ d, where it exceeds 80\\%.  We found that, at a given period, the SiO detection rate is approximately three times that for OH.  These facts suggest that the large-amplitude variables in the Nuclear Disk region are AGB stars similar in their overall properties to the inner and outer Bulge IRAS/SiO sources. From the set of radial velocity data, the mass distribution within 30 pc of the galactic center is derived by a new method which is based on the collisionless Boltzmann equation integrated along the line of sight. The mass within 30 pc is about $6.4 [\\pm 0.7] \\times 10^7 $ M$_{\\odot}$ and the mass of the central black hole is $2.7 [\\pm 1.3] \\times 10^6 $ M$_{\\odot}$. Consideration of the line-of-sight velocity of each star and its potential energy leads to the conclusion that the five high-velocity stars come from galactocentric distances as high as 300 pc. The high-velocity subsample of stars with negative radial velocities exhibits a tendency to have brighter $K$ magnitudes than the subsample of stars with positive velocities. The origin of these high-velocity stars is discussed. ",
        "introduction": "Radial-velocity data concerning stellar maser sources are useful for studying the dynamical behavior of the central part of the Galaxy (\\cite{lin92a}; \\cite{izu95}; \\cite{sjo98}; \\cite{deg00}). At visible wavelengths, they are difficult to obtain in the galactic center region because of interstellar extinction. Instead, most information comes from radio or near-infrared observations [for example, \\citet{sel87}]. In particular, observations of SiO and OH masers give radial velocities of stars accurate to within a few km s$^{-1}$. The masers arise in the circumstellar envelopes of mass-losing stars on the Asymptotic Giant Branch (AGB), which are intrinsically bright in the near- and mid-infrared regions, and which can potentially be identified at these wavelengths. Large numbers of candidate stars suitable for pointed maser surveys toward the nuclear disk have been discovered in the near-infrared $K$ band by making use of their characteristic large-amplitude variability (\\cite{gla01}).  The dynamical behavior of the central region of the Galaxy has attracted much attention, especially in relation to the central black hole (for example, \\cite{mor96}).  Proper motions of stars have been measured in the near-infrared $K$ band near the black hole (\\cite{gen00}; \\cite{ghe00}), and proper motions of SiO maser stars have also been measured within the central 15$''$ (\\cite{rei03}). These were used to find the position of Sgr A* in the $K$-band images (\\cite{men97}).  This paper concentrates on the dynamics of stars located towards the outer part of the central star cluster around the black hole, i.e., at about 2--30 pc distance, where the gravitational force of the black hole ceases to influence the stellar motions, and the stellar system is nearly self-gravitating. The dynamical (rotational) time scale in this region is a few $\\times 10^6$ y, while the ages of the AGB stars are $10^7$ -- $10^9$ y. Therefore, these stars are considered to be dynamically well relaxed (e.g., \\cite{hoz00}). Since the bar-like structure of the Galactic bulge was discovered (\\cite{bli91}; \\cite{nak91}; \\cite{dwe95}), it has been recognized that non-circular motions must be taken into account when interpreting observational data such as the CO gas distribution in the central nuclear disk (\\cite{bin91}; \\cite{wei99}).  Because double bars and nuclear rings have been proposed as efficient mechanisms for feeding gas into the centers of galaxies (\\cite{shl89}), it has become additionally important to look for signs of non-circularity in the motions of gas and stars.  In this paper, we report on the results of an SiO maser survey of Large Amplitude Variables (Miras or semiregulars; abbreviated as LAV hereafter) in a $24' \\times 24'$ area of the galactic center (\\cite{gla01}), whose amplitudes and periods are known (\\cite{woo98}; \\cite{gla01}). Because these stars are located at approximately the same distance (about 8 kpc) from the Sun, they constitute an ideal sample for studying the statistical characteristics of AGB stars and their detectability in the maser lines. In addition, surveying these sources gives accurate radial velocities, and provides basic data for investigating the kinematics of the galactic nuclear disk.  Although the mass within this region has been obtained previously by various methods [for example, \\citet{lin92a}], none of them [except \\citet{sah96}] are fully valid for treating the problem; the results obtained are likely to be in error by a factor of a few between the radii of 2 and 30pc. In this paper, we analyze the new SiO radial-velocity data set using the Boltzmann equation integrated along the line of sight. We also consider the origin of the high-velocity stars seen toward the galactic center.  %The position of the Sgr A*, which is considered to be at the dynamical center of the Galaxy, %is at ($RA$, $Dec$, epoch)= (17h45m40.045s,  $-29^{\\circ}00'27.9''$, J2000). Because Sgr A* is %shifted to the origin of  galactic coordinates by about $0.05^{\\circ}$,  it is more convenient %to use the coordinates in which Sgr A* is located at the origin. For this purpose, %we use the symbol,  $l^* (=l+0.056^{\\circ})$, and $b* (=b+0.046^{\\circ}$ in this paper.  %\\newpage ",
        "conclusions": "We have surveyed $\\sim 400$ large-amplitude variables within a $24'\\times 24'$ square about the galactic center, and obtained 180 detections (with additional 16 detections other than LAVs) in the SiO maser lines. The SiO detection rate of $\\sim$48\\% is comparable to that in previous SiO surveys of color-selected Bulge IRAS sources. The SiO detection rate increases with the period of light variation, and is well correlated with the OH detection rate.  The longitude-velocity diagram of the SiO sources has been revealed to be quite similar to the OH $l$--$v$ diagram. These facts suggest that the large-amplitude variables in the galactic nuclear disk are mass-losing stars in the AGB phase, quite similar to the IRAS sources in the inner Galactic bulge.  We also analyzed the SiO radial velocity data and obtained the mass distribution of the galactic center area. The mass of the central black hole that we have deduced, $2.7 (\\pm1.3) \\times 10^6 M_{\\odot}$, and the mass within 30 pc, $6.5 (\\pm 0.7) \\times 10^7 M_{\\odot}$, are more accurate than the previous estimates for these quantities. From analysis of the projected mass vs radius diagram, we found a tendency among the high-velocity sources that the subset with negative line-of-sight velocity is systematically brighter than the subset with positive line-of-sight velocity. This results from the fact that the the subsample with negative line-of-sight velocity is in front of the galactic center and the subsample with positive velocity is behind it. This tendency, which also applies to the Bulge SiO maser sources, strongly suggest the presence of streaming motion in the present nuclear-disk LAV sample.  The authors thank Dr. A. Winnberg for the useful comments. One of authors (I.S.G.) thanks the National Astronomical Observatory for providing him with a visiting fellowship for this work. This research was made use of the SIMBAD database operated at CDS, Strasbourg, France. It was partly supported by Scientific Research Grant (C2) 12640243 of the Japan Society for Promotion of Sciences.  %%%%%%%%%%  Appendix %%%%%%%%%%%%%"
    },
    "0403/astro-ph0403658_arXiv.txt": {
        "abstract": "We analyze the statistical properties and dynamical implications of galaxy distributions in phase space for samples selected from the 2MASS Extended Source Catalog. The galaxy distribution is decomposed into modes $\\delta({\\bf k, x})$ which describe the number density perturbations of galaxies in phase space cell given by scale band $\\bf k$ to ${\\bf k}+\\Delta {\\bf k}$ and spatial range $\\bf x$ to ${\\bf x}+\\Delta {\\bf x}$. In the nonlinear regime, $\\delta({\\bf k, x})$ is highly non-Gaussian. We find, however, that the correlations between $\\delta({\\bf k, x})$ and $\\delta({\\bf k', x'})$ are always very weak if the spatial ranges (${\\bf x}$, ${\\bf x}+\\Delta {\\bf x}$) and (${\\bf x'}$, ${\\bf x'}+\\Delta {\\bf x'}$) don't overlap. This feature is due to the fact that the spatial locality of the initial perturbations is memorized during hierarchical clustering. The highly spatial locality of the 2MASS galaxy correlations is a strong evidence for the initial perturbations of the cosmic mass field being spatially localized, and therefore, consistent with a Gaussian initial perturbations on scales as small as about 0.1 h$^{-1}$ Mpc. Moreover, the 2MASS galaxy spatial locality indicates that the relationship between density perturbations of galaxies and the underlying dark matter should be localized in phase space. That is, for a structure consisting of perturbations on scales from $k$ to $ k+\\Delta { k}$, the nonlocal range in the relation between galaxies and dark matter should {\\it not} be larger than $|{\\Delta {\\bf x}}|=2\\pi/|\\Delta {\\bf k}|$. The stochasticity and nonlocality of the bias relation between galaxies and dark matter fields should be no more than the allowed range given by the uncertainty relation $|{\\Delta {\\bf x}|| \\Delta{\\bf k}}|=2\\pi$. ",
        "introduction": "The large scale structure of the universe developed from initial mass density and velocity fluctuations through gravitational instability. Much of the information of the initial perturbations is ``forgotten\" during the gravitational nonlinear evolution. Yet, some features of the initial perturbations are imprinted in the cosmic mass and velocity fields at present. Measuring these memorized features is crucial in studying the initial states of the universe. A well known example is the self-similarity of gravitational clustering, which imprints the initial power spectrum index and is detectable with the scaling behavior of correlation functions of the present mass field (e.g. Peebles 1980). In this paper, we study the spatial locality of the perturbed mass field in phase space, which is also a feature to be memorized in mass field today.  The physics of spatial locality of correlation function in phase space can be illustrated with a Gaussian mass field $\\rho({\\bf x})$ described by \\begin{equation} \\langle\\hat{\\delta}({\\bf k}) \\hat{\\delta}({\\bf -k'})\\rangle = P(k)\\delta^K_{\\bf k,k'}, \\end{equation} where $\\hat{\\delta}({\\bf k})$ is the Fourier counterpart of the density contrast $\\delta({\\bf x})=[\\rho({\\bf x})-\\overline\\rho]/\\overline{\\rho}$, $P(k)$ the power spectrum of the mass field, and $\\delta^K$ the Kronecker delta function. Eq.(1) says that the Fourier modes with different wavevector ${\\bf k}$ are uncorrelated, or the correlation is localized in $k$-space. This is because the phase of the Fourier modes $\\delta({\\bf k})$ is random. On the other hand, the correlation function of density perturbations in physical ($x$) space generally is non-local. The two-point correlation function $\\langle \\delta({\\bf x})\\delta({\\bf x'})\\rangle$ has non-zero correlation length when the Fourier power spectrum $P(k)$ of eq.(1) is $k$-dependent.  In a phase-space description, the mass field is decomposed into modes $\\delta({\\bf k, x})$, the perturbations in the wavevector(scale) from ${\\bf k}$ to ${\\bf k} + \\Delta {\\bf k}$ and physical range ${\\bf x}$ to ${\\bf x}+\\Delta {\\bf x}$. The volume of the phase space cell referring to a mode is given by the uncertainty relation $|\\Delta {\\bf x}||\\Delta {\\bf k}| =2\\pi$. The correlation function of a Gaussian field generally is localized regardless whether the Fourier power spectrum is colored. That is, \\begin{equation} \\langle \\delta({\\bf k, x})\\delta({\\bf k', x'})\\rangle \\propto \\delta^K_{\\bf x,x'}\\delta^K_{\\bf k,k'}. \\end{equation} The reason for eq.(2) is straightforward. First, the perturbations $\\delta({\\bf k, x})$ and  $\\delta({\\bf k', x'})$ are, respectively, given by linear superposition of the Fourier modes in different wavebands (${\\bf k}$, ${\\bf k}+ \\Delta {\\bf k}$) and (${\\bf k'}$, ${\\bf k'}+ \\Delta {\\bf k'}$). For a Gaussian field, the Fourier modes in different wave band are uncorrelated in general [eq.(1)]. This gives rise to the factor of $\\delta^K_{\\bf k, k'}$ of eq.(2). Second, the phases of the Fourier modes of Gaussian field are random. For a superposition of the Fourier modes in the band (${\\bf k}$, ${\\bf k} +\\Delta {\\bf k}$) with random phase, the coherent length of the phase of $\\delta({\\bf k,x})$ is not larger than $|\\Delta {\\bf x}| \\simeq 2\\pi/|\\Delta {\\bf k}|$. On the other hand, the spatial distance between two cells in the ($k-x$) space is at least $|\\Delta {\\bf x}| \\simeq 2\\pi/|\\Delta {\\bf k}|$. This yields the factor $\\delta^K_{\\bf x,x'}$ of eq.(2). Therefore, the phase space correlation function eq.(2) for Gaussian field generally is local.  Non-linear evolution via gravitational clustering will lead to a non-Gaussian field that will deviate from the locality of eq.(2), even when the initial field is Gaussian. However, turbulence studies have found that if  1.) the initial perturbations are spatially localized, and 2.) the random fields evolve via a self-similar hierarchical cascade process, the phase space correlation function of the evolved field will still be spatially localized (Greiner, Lipa, \\& Carruthers, 1995, Greiner et al. 1996). The perturbations of modes $\\delta({\\bf k, x})$ and $\\delta({\\bf k', x'})$ with ${\\bf x}\\neq {\\bf x'}$ stay statistically uncorrelated or only very weakly correlated during the self-similar hierarchical cascade evolution. That is, the factor $\\delta^K_{\\bf x,x'}$ of eq.(2) is memorized during the dynamical evolution.  Phenomenological models that mimic the hierarchical clustering of the cosmic mass field have similar mathematical structures as the hierarchical cascade models of turbulence. For instance, the fractal hierarchy clustering model of Soneira \\& Peebles (1977) is the same as the $\\beta$ model of turbulence (e.g. Frisch, 1995). The block model of Cole \\& Kaiser (1988) is a special case of the multifractal cascade model (Meneveau \\& Sreenivasan 1987, Pando et al 1998). Hence, these models should also memorize the spatial locality. Recently, the spatial locality has been studied with more realistic dynamical models of gravitational clustering. First, if the weakly nonlinear mass field is given by the Zeldovich approximation, the field is found to be spatially localized if the initial perturbations are Gaussian (Pando, Feng \\& Fang 2001). More recently, this result has been extended to fully nonlinear regime (Feng \\& Fang 2004). Using the halo model of the large scale structure (e.g. Cooray \\& Sheth 2002, and references therein), it has been shown that the evolved mass field is approximately spatially localized if the initial perturbations are Gaussian. Although gravitational coupling is long-term, the spatial locality in the phase space is not disturbed by the non-linear evolution of cosmic mass field. This property essentially is due to the self-similarity of the hierarchical clustering. This result has been tested with high resolution N-body simulation samples (Feng \\& Fang 2004).  In this paper, we investigate the spatial locality of phase space correlations with real sample -- the galaxies selected from the 2MASS Extended Source Catalog (XSC, Jarrett et al. 2000). Our motivation is two-fold. First, the spatial locality provides a test of the Gaussianity of the initial density perturbations on small scales. Although many tests on the Gaussianity of the initial perturbations have been done with the temperature fluctuations of the Cosmic Microwave Background Radiation (e.g. Komatsu, et al. 2003, Pando, Valls-Gabaud \\& Fang, 1998), the comoving scales of these tests are not less than a few Mpc. The spatial locality of the 2MASS samples can test the Gaussianity to scale as small as about 0.04 h$^{-1}$ Mpc. Second, the distribution of galaxies is biased from the mass field of dark matter. Some bias models assume that the relation between the distribution of galaxies and underlying mass field is stochastic and nonlocal (e.g. Dekel \\& Lahav, 1999). This mechanism will lead to nonlocality of the galaxy correlation function in phase space, even when the underlying dark matter mass field is spatially localized. Therefore, the spatial locality should be effective in testing the stochasticity and nonlocality of bias models.  The outline of this paper is as follows. \\S 2 presents the statistics and dynamics of the spatial locality in phase space. \\S 3 describes the basic properties of the 2MASS galaxy samples with a space-scale decomposition. The results of the spatial locality of the 2MASS galaxy distribution are presented in \\S 4. Finally, the conclusions and discussions are be given in \\S 5. ",
        "conclusions": "The 2-D distribution of 2MASS galaxies has been studied using the DWT, in which the SFC $\\epsilon_{\\bf j,l}$ is the count-in-cell in physical space, while the WFC $\\tilde{\\epsilon}_{\\bf j,l}$ is a count-in-cell in phase space. We find that the statistical properties of the SFC and WFC variables are very different. The former is non-Gaussian and nonlocally correlated, while the later is non-Gaussian, but its correlation basically is spatially localized. That is, the sample has the following statistical behavior in the phase space \\begin{itemize} \\item the one-point distribution of $\\tilde{\\epsilon}_{\\bf j,l}$ is highly non-Gaussian on angular scales less than 1$^{\\circ}$, corresponding to $\\simeq 3.5$ h$^{-1}$ Mpc at the median redshift of the sample. \\item The 2nd and 4th order mode-mode correlations with $\\tilde{\\epsilon}_{\\bf j,l}$ are always spatially localized. For highest $j$, the locality holds on angular scales as small as 0.06$^{\\circ}$, or $\\simeq$ 0.2 h$^{-1}$ Mpc at the median redshift of the sample. \\item the local scale-scale correlations are significant on scales less than 0.4$^{\\circ}$, corresponding to $\\simeq 1.5$ h$^{-1}$ Mpc at the median redshift of the sample. The scale-scale correlation is also approximately spatially localized. \\end{itemize} A direct physical meaning of these results is that the cosmic gravitational instability causes only strong interaction between the modes in different wavebands and in the same spatial area, but is weak for modes in different spatial area.  Because the nonlinear evolution of cosmic gravitational clustering has the memory of its initial spatial correlation in the phase space, the observed spatial locality of the 2MASS galaxies provides solid evidence for models assuming that the initial perturbations are spatially uncorrelated among phase space modes. This result is consistent with the assumption that the initial perturbations are Gaussian. Although the Gaussianity of the initial perturbations on large scales has been extensively tested with the temperature fluctuations of the cosmic microwave background radiation, the test given by the spatial locality is effective to comoving scales as small as $\\simeq$ 0.2 h$^{-1}$ Mpc.  The spatial locality has been studied with samples of Ly-alpha forests (Pando, Feng \\& Fang, 2001). Since Ly-alpha forests refer to weakly non-linear clustering, their spatial locality can be used to rule out the models of initially non-Gaussian fields which are non-spatially localized, but cannot do so for models of initial non-Gaussian which are spatially localized. In this paper, the spatial locality is found for sample referring to fully nonlinear regime. The memory of the spatial locality is based on the halo model, for which an initially Gaussian field is necessary. In this sense, the spatial locality of the 2MASS gives stronger support to the initially Gaussian assumption than that that of Ly-alpha forests.  The one-point distribution of 2MASS galaxies (Fig. 3) is very different from that given by N-body simulation sample. The later in nonlinear regime generally is lognormal (see Fig. 6 of Feng \\& Fang, 2004), while the former is more complicated (Fig. 3). This difference indicates that galaxy distribution is biased with respect to the underlying dark matter field. Nevertheless, the galaxy distribution is also highly spatially localized. This result indicates that relationship between $\\rho_g({\\bf x})$ and $\\rho({\\bf x})$ should be localized in phase space. That is, the stochasticity and nonlocality of the relation between $\\rho_g({\\bf x})$ and $\\rho({\\bf x})$ are limited in the cell of phase space. For a structure consisting of perturbations on scales ${\\bf k}$ to ${\\bf k}+\\Delta {\\bf k}$, the nonlocal size in the relation between $\\rho_g({\\bf x})$ and $\\rho({\\bf x})$ should not be larger than $|\\Delta {\\bf x}|=2\\pi/|\\Delta {\\bf k}|$. Otherwise, the galaxy bias will violate the spatial locality of the galaxy distribution even when the underlying field is spatial localized. Hence, from the 2MASS galaxies we can conclude that the stochasticity and nonlocality of the bias relation between $\\rho_g({\\bf x})$ and $\\rho({\\bf x})$ probably are no more than that given by the uncertainty relation $|\\Delta {\\bf x}||\\Delta {\\bf k}|=2\\pi$."
    },
    "0403/astro-ph0403691_arXiv.txt": {
        "abstract": "We have fitted the {\\it Chandra\\/} Low Energy Transmission Grating spectrum of SS~Cygni in outburst with a single temperature blackbody suffering the photoelectric opacity of a neutral column density and the scattering opacity of an outflowing wind. We find that this simple model is capable of reproducing the essential features of the observed spectrum with the blackbody temperature $T_{\\rm bl}\\approx 250\\pm 50$ kK, hydrogen column density $N_{\\rm H}\\approx 5.0^{+2.9}_{-1.5}\\times 10^{19}~\\rm cm^{-2}$, fractional emitting area $f\\approx 5.6^{+60}_{-4.5}\\times 10^{-3}$, boundary layer luminosity $L_{\\rm bl}\\approx 5^{+18}_{-3}\\times 10^{33}~\\rm erg~s^{-1}$, wind velocity $v\\approx  2500~\\rm km~s^{-1}$, wind mass-loss rate $\\Mdot_{\\rm w}\\approx 1.1\\times 10^{16}~\\rm g~s^{-1}$, and arbitrary values of the wind ionization fractions of 20 ions of O, Ne, Mg, Si, S, and Fe. Given that in outburst the accretion disk luminosity $L_{\\rm disk}\\approx 1\\times 10^{35}~\\rm erg~s^{-1}$, $L_{\\rm bl}/L_{\\rm disk}\\approx 0.05^{+0.18}_{-0.03}$, which can be explained if the white dwarf (or an equatorial belt thereon) is rotating with an angular velocity $\\Omega_{\\rm wd}\\approx 0.7^{+0.1}_{-0.2}$ Hz, hence $V_{\\rm rot}\\sin i\\sim 2300~\\rm km~s^{-1}$. ",
        "introduction": "According to simple theory, the boundary layer between the accretion disk and the surface of the white dwarf is the dominant source of high energy radiation in nonmagnetic cataclysmic variables (CVs). Unless the white dwarf is rotating rapidly, the boundary layer luminosity $L_{\\rm bl}\\approx G\\Mwd\\Mdot /2\\Rwd $, where $\\Mwd $ and $\\Rwd $ are respectively the mass and radius of the white dwarf and $\\Mdot $ is the mass-accretion rate. When $\\Mdot $ is low, as in dwarf nova in quiescence, the boundary layer is optically thin to its own radiation and its temperature is of order the virial temperature $T_{\\rm vir} = G\\Mwd m_{\\rm H}/3k\\Rwd\\sim 10$ keV. When $\\Mdot $ is high, as in nova-like variables and dwarf novae in outburst, the boundary layer is optically thick to its own radiation and its temperature is of order the blackbody temperature $T_{\\rm bb} = (G\\Mwd\\Mdot /8\\pi\\sigma f\\Rwd^3)^{1/4}\\sim 100$ kK.  The compact nature of the boundary layer is established by the narrow eclipses of the hard X-ray flux of high inclination dwarf novae in quiescence \\citep{woo95a, muk97, tes97, pra99a, whe03b} and the short periods [$P\\sim 10~{\\rm s}\\sim 2\\pi(\\Rwd ^3/G\\Mwd )^{1/2}$] of the oscillations of the soft X-ray and extreme ultraviolet (EUV) flux of nova-like variables and dwarf novae in outburst \\citep{war04, mau04a}. However, an {\\it extended\\/} source of soft X-ray and EUV flux is required to explain the lack of eclipses in high inclination nova-like variables and dwarf novae in outburst \\citep{nay88, woo95b, pra99b, mau00, pra04}.  Using the {\\it Extreme Ultraviolet Explorer\\/} ({\\it EUVE\\/}) light curve and spectrum of OY Car in superoutburst, \\citet{mau00} argued that the source of the extended EUV flux of high-$\\Mdot $ nonmagnetic CVs is the accretion disk wind. In this picture, in high inclination systems the boundary layer is hidden by the accretion disk for all binary phases, but its flux is scattered into the line of sight by resonance transitions of ions in the wind. Because the scattering optical depths of these transitions are much higher than the Thompson optical depth, the EUV spectra of such systems should be dominated by broad lines.  Consistent with this picture, in outburst the EUV spectra of the high inclination dwarf novae OY~Car ($i\\approx 83^\\circ $; \\citealt{mau00}) and WZ~Sge ($i\\approx 75^\\circ $; Wheatley \\& Mauche 2004, in preparation) are dominated by broad lines, while that of the low inclination dwarf novae VW~Hyi ($i\\sim    60^\\circ $; \\citealt{mau96}) and SS~Cyg ($i\\sim    40^\\circ $; \\citealt{mau95}) are dominated by the continuum. U~Gem ($i\\approx 70^\\circ $; \\citealt{lon96}) is a transitional case, as its EUV flux is partially eclipsed and its EUV spectrum contains of broad emission lines superposed on a bright continuum.  The general case of the radiation transfer of the boundary layer flux through the accretion disk wind is complicated, but schematically we expect that in high inclination systems the EUV spectrum will be of the form $F_\\lambda\\, e^{-\\tau_\\lambda ^\\prime}\\, (1-e^{-\\tau_\\lambda})$, while in low inclination systems it will be of the form $F_\\lambda\\, e^{-\\tau_\\lambda ^\\prime} e^{-\\tau_\\lambda}$, where $F_\\lambda$ is the intrinsic boundary layer spectrum, $\\tau_\\lambda ^\\prime$ is the photoelectric optical depth of the wind and the interstellar medium, and $\\tau_\\lambda$ is the scattering optical depth of the wind. \\cite{mau00} showed that the second version of this model describes the main features of the {\\it EUVE\\/} spectrum of OY~Car in superoutburst, and we show here that the first version of this model describes the main features of the {\\it Chandra\\/} Low Energy Transmission Grating (LETG) spectrum of SS~Cyg in outburst. A preliminary discussion of this spectrum is supplied by \\citet{mau04a}, while the quasi-coherent oscillations of the EUV flux is discussed by \\citet{mau02}. ",
        "conclusions": "We have fitted the {\\it Chandra\\/} LETG spectrum of SS~Cyg in outburst with a single temperature blackbody suffering the photoelectric opacity of a neutral column density and the scattering opacity of an outflowing wind. The best-fit model, shown in Figure~3, has parameters $T_{\\rm bl}=250$ kK, $N_{\\rm H} = 5.0\\times 10^{19}~\\rm cm^{-2}$, $f=5.6\\times 10^{-3}$, $\\Mdot_{\\rm w}=1.1 \\times 10^{16}~\\rm g~s^{-1}$, and the values of $X$ listed in Table~1 ($X=1$ otherwise). Assuming that the wind is optically thin and photoionized by a 250 kK blackbody, the range of ions present in the model requires that the ionization parameter $-1\\lax\\log\\xi\\lax 5$.  Given these results, in outburst the boundary layer luminosity $L_{\\rm bl} = 4.7\\times 10^{33}~\\rm erg~s^{-1}$. A lower limit on the value of this quantity is imposed by the value of the interstellar H column density, which \\citet{mau88} estimated to be $3.5\\times 10^{19}~\\rm cm^{-2}$ based on the curve-of-growth of interstellar absorption lines in high resolution {\\it International Ultraviolet Explorer\\/} spectra of SS~Cyg in outburst. For $-1\\lax\\log\\xi\\lax 5$, H and He are fully ionized, so the wind should have no photoelectric opacity in the EUV bandpass, and $N_{\\rm H}$ should be equal to the interstellar H column density. With $N_{\\rm H}=3.5\\times 10^{19}~\\rm cm^{-2}$ and $T_{\\rm bl}=300$ kK, Figure~2 shows that $f=1.1\\times 10^{-3}$, hence $L_{\\rm bl}=1.7\\times 10^{33}~\\rm erg~s^{-1}$. At the other end of the $\\chi^2$ ellipse, with $T_{\\rm bl}=200$ kK and $N_{\\rm H}=7.9\\times 10^{19}~\\rm cm^{-2}$, $f=6.6\\times 10^{-2}$, hence $L_{\\rm bl}= 2.3\\times 10^{34}~\\rm erg~s^{-1}$. We conclude that a reasonable estimate of the boundary layer luminosity $L_{\\rm bl}\\approx 5^{+18}_{-3}\\times 10^{33}\\, (d/160~{\\rm pc})^2~\\rm erg~s^{-1}$.  How does this compare with the accretion disk luminosity? \\citet{pol84} used optical through far ultraviolet spectra of SS~Cyg in outburst to determine that $L_{\\rm disk}\\approx 1\\times 10^{35}\\, (d/160~{\\rm pc})^2~\\rm erg~s^{-1}$, so $L_{\\rm bl}/L_{\\rm disk}\\approx 0.05^{+0.18}_{-0.03}$. Theoretically, we expect that this ratio is equal to one unless the white dwarf is rotating rapidly, in which case $L_{\\rm bl}/L_{\\rm disk}=(1-\\omega)^2$, where $\\omega =\\Omega_{\\rm wd}/\\Omega_{\\rm K}(\\Rwd )$, $\\Omega_{\\rm wd}=2\\pi/P_{\\rm spin}$ is the angular velocity of the white dwarf, and $\\Omega_{\\rm K}(\\Rwd ) = (G\\Mwd/\\Rwd^3)^{1/2}\\approx 0.9$ Hz is the Keplerian angular velocity of material just above its surface \\citep{pop95}. The white dwarf can be spun up to the required rate by accreting an amount of mass $\\Delta M=\\Mwd - M_{\\rm wd,i}$, where $M_{\\rm wd,i}$ is the initial white dwarf mass, $\\Mwd = M_{\\rm wd,i}\\, (1-4k^2\\omega/3)^{-3/4}$ is the final white dwarf mass, and $k$ is the radius of gyration \\citep{lan00}. With $k\\approx 0.4$ and $\\omega = 1- (L_{\\rm bl}/L_{\\rm disk})^{1/2}\\approx 0.78^{+0.08}_{-0.26}$, $\\Delta M\\approx 0.13\\, \\Mwd\\sim 0.1~\\Msun $. If rapid rotation is the cause of SS~Cyg's low boundary layer luminosity, the white dwarf angular velocity $\\Omega_{\\rm wd} =\\omega\\Omega_{\\rm K}(\\Rwd )\\approx 0.7^{+0.1}_{-0.2}$ Hz, hence the spin period $P_{\\rm spin}\\approx 9^{+4}_{-1}$ s (comparable to the period of the quasi-coherent oscillations), the rotation velocity $V_{\\rm rot} = \\Omega_{\\rm wd}\\Rwd\\approx 3800^{+400}_{-1300}~\\rm km~s^{-1}$, and $V_{\\rm rot}\\sin i\\sim 2300~\\rm km~s^{-1}$. This last quantity is rather large, given that $V_{\\rm rot}\\sin i\\approx 1200~\\rm km~s^{-1}$ for WZ~Sge and $V_{\\rm rot}\\sin i\\lax 400~\\rm km~s^{-1}$ for six other nonmagnetic CVs \\citep{sio99}.  Finally, we note that in the single scattering limit, conservation of momentum limits the mass-loss rate of a radiatively driven wind to $\\Mdot_{\\rm max} =L/V_\\infty c$, where $L$ is the system luminosity and $V_\\infty $ is the wind terminal velocity. For SS~Cyg in outburst, $L=L_{\\rm disk}+L_{\\rm bl}\\approx 1\\times 10^{35}~\\rm erg~s^{-1}$ and $V_\\infty\\approx 3000~\\rm km~s^{-1}$, so $\\Mdot_{\\rm max}\\approx 1\\times 10^{16}~\\rm g~s^{-1}$. We derived a wind mass-loss rate $\\Mdot_{\\rm w}\\approx 1\\times 10^{16}~\\rm g~s^{-1}$ for SS~Cyg in outburst, so $\\Mdot_{\\rm w}/\\Mdot_{\\rm  max}\\approx 1$. In $\\tau $~Sco (B0 V), P~Cyg (B1 Ia$^+$), $\\zeta $~Pup (O4 f), and $\\epsilon $~Ori (B0 Ia) $\\Mdot_{\\rm w}/\\Mdot_{\\rm max}=0.02$, 0.22, 0.33, and 0.65, but in WR1 (WN5) $\\Mdot_{\\rm w}/\\Mdot_{\\rm max}=60$ \\citep{lam99}, so we cannot conclude from this result that a mechanism other than radiation pressure is needed to drive the wind of SS~Cyg."
    },
    "0403/astro-ph0403372_arXiv.txt": {
        "abstract": "In the last few years new evidence has been presented for the presence of ongoing massive star formation in the outer HI disks of galaxies. These discoveries strongly suggest that precursor molecular gas must also be present in some physical state which is escaping detection by the usual means (CO(1-0), IR, etc.).  We present a model for such a gas in a framework which views the HI as the result of an ongoing ``photodissociation $\\Leftrightarrow$ dust grain reformation'' equilibrium in a cold, clumpy molecular medium with a small area filling factor. ",
        "introduction": "Recently Komiyama et al.\\ (2003) and De Blok \\& Walter (2003) have reported the detection of young stars in the outer HI disk of the nearest dwarf irregular galaxy NGC 6822 ($D \\approx 490$ kpc) at galactocentric radii well beyond $R_{25}$.  These authors discuss the results in terms of the commonly-accepted picture, where young massive stars form out of the observed HI in response to some ``triggering'' mechanism such as the passage of a companion galaxy (in this case a massive intergalactic HI cloud). However, an ``interaction'' scenario is problematic for NGC 6822, since as De Blok \\& Walter point out there is little evidence for tidal ``impact trauma'' on the distribution of old stars in the galaxy. ",
        "conclusions": "\\begin{itemize} \\item \\textit{\\textbf{The hypothesis:} An $H_2$ photodissociation $\\Leftrightarrow$ reformation equilibrium scenario provides a simple explanation for the coexistence of young stars, dust, and HI in the outer parts of disk galaxies.} \\item \\textit{\\textbf{The hypothesis is falsifiable:}  If we find dynamically quiescent HI regions in the outer parts of galaxies without dissociating stars ($M(V) \\leq -1$, B5 or brighter, $M \\geq 6 M_{\\sun}$) then this idea is wrong.} \\item  \\textit{\\textbf{A prediction:} Even low surface brightness galaxies such as NGC 2915 will have dissociating B-stars mixed in with their HI.} \\end{itemize}"
    },
    "0403/astro-ph0403144_arXiv.txt": {
        "abstract": "A modified non-linear time series analysis technique, which computes the correlation dimension $D_2$, is used to analyze the X-ray light curves of the black hole system GRS 1915+105 in all twelve temporal classes. For four of these temporal classes $D_2 $ saturates to $\\approx 4-5$ which indicates that the underlying dynamical mechanism is a low dimensional chaotic system. Of the other eight classes, three show stochastic behavior while five show deviation from randomness. The light curves for four classes which depict chaotic behavior have the smallest ratio of the expected Poisson noise to the variability ($ < 0.05$) while those for the three classes which depict stochastic behavior is the highest ($ > 0.2$). This suggests that the temporal behavior of the black hole system is  governed by a low dimensional chaotic system, whose nature is detectable only when the Poisson fluctuations are much smaller than the variability. ",
        "introduction": "\\label{sec: I}   Black hole X-ray binaries are variable on a wide range of timescales ranging from months to milli-seconds. A detailed analysis of their temporal variability is crucial to the understanding of the geometry and structure of these high energy sources. Such studies may eventually be used to test the relativistic nature of these sources and to understand the physics of the accretion process. The variability in different energy bands is generally quantified by computing the power spectrum which is the amplitude squared of the Fourier transform. The power spectra give information about the characteristic frequencies of the system which show up as either breaks or as near Gaussian peaks, i.e Quasi-Periodic Oscillation (QPO) in the spectra (e.g. \\citet{Bel01,Tom01,Rod02}). The shape of the power spectra, combined with the observed frequency dependent time lags between different energy bands, have put constraints on the radiative mechanisms and geometry of emitting regions (e.g. \\citet{Now99,Mis00,Cui99,Pou99,Cha00,Nob01})   These results are based on the response of the system to temporal variations whose origin is not clear. Important insight into the origin can be obtained by the detection and quantification of the possible  non-linear behavior of the fluctuations. For example, the presence of stochastic  fluctuations would favor X-ray variations driven by variations  of some external parameters (like the mass accretion rate), or the possibility that active flares occur randomly. On the other hand, if the fluctuations can be described as a deterministic chaotic system, then  inner disk instability or coherent flaring activity models will be the likely origin. A quantitative description of the temporal behavior can also be compared with time dependent numerical simulations of the accretion process and will help examine the physical relevance of these simulations.  The non-Gaussian and non-zero skewness values of the temporal variation of the black hole system Cygnus X-1 suggested that the variations are non-linear in nature \\citep{Thi01,Tim00,Mac02}. More rigorous tests were applied to the AGN ArK 564 \\citep{Gli02} which also suggested non-linear behavior. Nonlinear time series analysis(NLTS) seems to be the most convenient tool to check if the origin of the  variability is chaotic, stochastic or a mixture of the two and has been adopted in several disciplines  to study complex systems (e.g. human brain, weather) and predict their immediate future \\citep{Sch99}. This technique has also been used earlier to analyze  X-ray data of astrophysical sources. Based on a NLTS analysis of EXOSAT data, \\cite{Vog87} claimed that  the X-ray Pulsar Her X-1  was a low dimensional chaotic system. However, \\cite{Nor89} pointed out problems with that analysis since the source has a strong periodicity and the data analyzed had low signal to noise ratio. \\cite{Leh93} used the NLTS technique to analyze EXOSAT light curves of several AGN,  and found that only one, NGC 4051, showed signs of low dimensional chaos. A similar analysis on the noise filtered{\\it Tenma} satellite data of Cyg X-1, suggested that the source may be a low dimensional chaotic system with large intrinsic noise \\citep{Unn90}. These analysis were hampered by small number of data points ($\\simless 1000$) in the light curve and/or noise. Hence, the reported detection of low dimension chaos was only possible by rather subjective comparison of the results of the data analysis with those from simulated data of chaotic systems with noise.    The Galactic micro quasar GRS 1915+105 is a highly variable black hole system. It shows a wide range of variability \\citep{Che97,Pau97,Bel97a} which required \\citet{Bel00} to classify the behavior in no less than twelve temporal classes. In this work, our motivation is to determine the temporal property of this source by using a modified nonlinear time series analysis for each of these twelve classes. The different kinds of variability and its brightness ( the average  RXTE PCA count rate ranges from $5000-32000$ counts/s) makes this source an ideal one to detect chaotic behavior.  In the next section we describe the technique used to determine the Correlation dimension. The results of the analysis are presented in \\S 3, while in \\S 4 the work is summarized and discussed. ",
        "conclusions": "The saturation of the correlation dimension $D_2 \\approx 4-5$ for four of the temporal classes clearly indicates that the underlying dynamic mechanism that governs the variability of the black hole system is a low dimensional chaotic one. As indicated by simulations of the Lorenz system with noise, the effect of Poisson noise in the data is to increase the $D_2$ values. Hence the real dimension of the system is probably smaller than $D_2 \\approx 4-5$ that is obtained here. In fact it is possible that the the temporal behavior of the black hole system is always governed by a low dimensional chaotic system, but is undetectable when Poisson noise affects the analysis.   Alternatively, there may be a stochastic component to the variability which dominates for certain temporal classes. The two scenarios may be distinguished and better quantitative estimates of the correlation dimension may be obtained by either appropriate noise filtering of the data and/or appropriate averaging of the different light curves. Much longer ($\\approx 30000$ sec long) continuous data streams sampled  at $1$ second resolution, would decrease Poisson noise and hence provide better quantitative measure of $D_2$. However, such long data streams are presently not available and merging non-continuous light curves, will require sophisticated gap filling techniques which might give rise to spurious results.   The  variability of GRS 1915+105 can be interpreted as being transitions between three spectral states \\citep{Bel00}, one of which (the so called soft state) is a long term canonical state observed in other black hole systems like Cygnus X-1 which do not show such high amplitude variability. It is attractive to identify these spectral states as fixed points which for GRS 1915+105 become unstable giving rise to the observed chaotic behavior which may also account for the ring like movement of the system in color-color space \\citep{Vil98}. The above hypothesis may be verified by future characterization of the chaos in GRS 1915+105. Note that GRS15+105 spends most of it's time in the $\\chi$ class whose variability is similar to that observed in other black hole systems like Cygnus X-1. However, as shown in this work, Poisson noise effects the analysis for the $\\chi$ class and the $D_2 (M)$ values reflect stochastic behavior. This may be the reason why earlier different non-linear analysis of Cygnus X-1 data, while showing non-linearity \\citep{Tim00,Thi01} did not conclusively reveal chaotic behavior.  The identification of the temporal behavior of the black hole system as a chaotic one, has opened a new window toward the understanding of the origin and nature of their variability. The present analysis can be extended to characterize the chaotic behavior. Using the minimum required phase space dimension, the data can be projected into different $2$ dimensional planes, which will reveal the structure of the attractor and help to identify any possible centers of instability in the system. Further, dynamical invariants like the full Lyapunov spectrum , multi-fractal dimensions etc. can be also be computed. Recently, \\citet{Win03} have studied and quantified the chaotic flow in magneto-hydrodynamic simulations of the mass accretion processes that is believed to be happening in black hole systems. The measured chaos parameters like the largest Lyapunov exponent for such simulations can be compared with that obtained from the light curve of black hole systems to validate such simulations and enhance our understanding of these systems. Note that such analysis can practically be applied only after the identification of the minimum phase space dimension which in turn usually requires the computation of $D_2 (M)$."
    },
    "0403/astro-ph0403199_arXiv.txt": {
        "abstract": "Mean pulse profiles and polarization properties are presented for nine southern pulsars. The observations were made using the Parkes radio telescope at frequencies near 1330 MHz; three of the nine pulsars were also observed at 660 MHz. A very high degree of circular polarization was observed in PSR J1603$-$7202. Complex position angle variations which are not well described by the rotating-vector model were observed in PSRs J2124$-$3358 and J2145$-$0750, both of which have very wide profiles. Rotation measures were obtained for all nine pulsars, with two implying relatively strong interstellar magnetic fields. ",
        "introduction": "Observations of the polarization of mean pulse profiles have given much information about the pulsar radio emission mechanism. Pulsars often have very high linear polarization and the sweep of position angle across the pulse is generally well described by the rotating-vector model \\citep{rc69a,kom70}. This and the bilateral symmetry of many pulse profiles led to a model in which the emission is beamed into a hollow cone centered on the magnetic axis. With observations of an increasing number of pulsars it soon became clear that the mean pulse profiles of many pulsars were more complex. The frequent occurence of a central emission peak \\citep{bac76} led to the idea of a `core' component in contrast to the outer `conal' components. Further observations showed that many pulsars had multiple conal components leading to ideas of multiple cones \\citep{ran93,kwj+94,md99} or patchy beams \\citep{lm88,hm01}. Core emission tends to have a steeper spectrum than conal emission and a higher degree of circular polarization, often with a reversal of sense near the profile center \\citep{ran83,lm88}, although reversing circular polarization is sometimes observed in conal emission \\citep{hmxq98}. The apparent width of both core and conal beams is a function of pulse period ($P$), with an approximate $P^{-1/2}$ dependence \\citep{ran90,big90b}.  With the discovery of millisecond pulsars (MSPs), the parameter ranges available for study of pulse emission processes was greatly extended. MSPs have much weaker surface magnetic fields than `normal' pulsars, with $B_0$ typically a few by $10^8$~G, compared to $10^{12}$~G for normal pulsars. They also have much smaller magnetospheres, conventionally defined to be limited by the light cylinder at radius $R_{LC} = cP/2\\pi$, where $P$ is the pulsar period and $c$ is the velocity of light. For the fastest MSP, PSR B1937+21, $R_{LC}$ is only 74 km, or about ten times the neutron star radius. One would therefore expect the radio emission from MSPs to be rather different to that from normal pulsars. MSP pulse profiles are generally wider (in pulse phase) than those of normal pulsars, which is not surprising in view of the wider angle subtended by the `open' field lines which penetrate the light cylinder and define the extent of the polar cap. Also, MSP profiles have different frequency evolution characteristics compared to most normal pulsars. In particular, the component separations are largely independent of frequency \\citep{kll+99}. It is not clear that the distinctions and relationships between `core' and `conal' emission seen in longer-period pulsars apply to MSPs. For example, the strong component in PSR J0437$-$4715 is clearly central to the emission beam but has a flatter spectrum than the outer components \\citep{nms+97}.  Despite these differences, there are many similarities in the radio emission properties of MSPs and normal pulsars. If the greater width of MSP profiles is ignored, the shapes of MSP mean pulse profiles are very similar to those of normal pulsars. However, among the MSPs, there is a greater proportion of `interpulse' profiles, with two main regions of pulse emission separated by approximately $180\\degr$ of pulse phase or longitude.  MSP profiles often appear to have more components than normal pulsars, but Kramer et al. (1998)\\nocite{kxl+98} have argued that this is simply because MSP components are more obvious because of their wider spacing. On average, the radio emission of MSPs has a spectral index which is only marginally steeper than that of normal pulsars and may in fact be the same \\citep{tbms98,kxl+98}. MSPs are typically somewhat less luminous than normal pulsars but there is a large overlap in the luminosity distributions. Finally, the polarization properties of MSPs and normal pulsars are remarkably similar, with both often having high fractional linear polarization and generally smaller levels of circular polarization. Orthogonal jumps in position angle are seen in both classes of pulsar (e.g. Stinebring et al. 1984\\nocite{scr+84}; Thorsett \\& Stinebring 1990\\nocite{ts90}).  Variations of position angle (PA) through MSP profiles are often more complex than is typical for normal pulsars. Despite this, in many cases the PA can be approximately fitted by the simple rotating vector model (RVM) which applies to most normal pulsars, at least once orthogonal jumps are taken into account (e.g. Arzoumanian et al. 1996)\\nocite{aptw96}. These similarities leave little doubt that the basic radio emission mechanism is the same in normal pulsars and MSPs.  Extensive studies of the polarization properties of MSPs have previously been made by Thorsett \\& Stinebring (1990)\\nocite{ts90}: PSRs B1937+21, B1953+29 and B1957+20 at several frequencies between 430 and 2830 MHz; Navarro et al. (1997)\\nocite{nms+97}: PSR J0437$-$4715 at 438, 660 and 1512 MHz; Arzoumanian et al. (1996)\\nocite{aptw96}: PSR B1534+12 at 430 MHz; Xilouris et al. (1998)\\nocite{xkj+98}: 24 pulsars at 1410 MHz; Stairs, Thorsett \\& Camilo (1999)\\nocite{stc99}: 19 pulsars at 410, 610 and/or 1414 MHz; and Lommen et al. (2000)\\nocite{lzb+00}: PSR J0030+0451 at 430 MHz. All observations except those of PSR J0437$-$4715 by Navarro et al. were made using northern hemisphere telescopes, so there is little information on the polarization of the southern MSPs. Also, there are significant discrepancies in the results of different groups for some pulsars. For example, results from Xilouris et al. (1998) and Stairs et al. (1999) for PSRs J1022+1001, J1713+0747 and J2145-0750 at frequencies near 1400 MHz differ substantially.\\footnote{Stairs et al. (1999) use a convention for sense of position angle which is opposite to the astronomical convention used by other authors.} It is important to understand whether these differences are due to time variations such as mode changing \\citep{kxc+99} or problems with data processing and/or calibration. Furthermore, in some previous studies (e.g. Stairs et al. 1999) the position angle scale is not absolute. Absolute position angles are important for Faraday rotation studies and in comparisons with other properties. Important examples are the comparison of the pulsar rotation axis direction implied by radio polarization measurements with that deduced from X-ray observations and with the direction of the pulsar proper motion. \\citep{hgh01,rd01}.  With their high linear polarization, it is relatively easy to determine rotation measures (RMs) for most pulsars. Combined with the dispersion measure (DM), pulsars give a direct measure of the mean line-of-sight magnetic field between the Sun and the pulsar, weighted by the local electron density. Since the DM also provides an estimate of the pulsar distance, pulsars are a valuable probe of the Galactic magnetic field (e.g. Han et al. 2002)\\nocite{hmlq02}. Most MSPs are relatively close to the Sun, so they provide information on conditions in the local region of the Galaxy.  We have used the Parkes 64-m telescope of the Australia Telescope National Facility and the Caltech pulsar correlator to make observations of nine southern MSPs at 20~cm and, for three of these, at 50~cm. Five of these pulsars have no previously published polarization observations and none have previously published rotation measures. Details of the observing system and observational parameters are given in \\S~2. The polarization and RM results are presented in \\S~3 and \\S~4 respectively, and in \\S~5 we discuss some implications of our polarization measurements. ",
        "conclusions": "In this paper we have presented pulse profiles and polarization parameters for nine southern MSPs, five of which had no previously published polarization data. New RM values are given for all nine pulsars. Notable results include the very high fractional circular polarization ($\\sim 65$\\% at 1327 MHz) in the trailing component of PSR J1603$-$7202, and the extremely wide and complex pulse profile and polarization properties of PSR J2124$-$3358. \\citet{kxl+98} have argued that, on average, MSPs have profiles which are of similar complexity (quantified by the number of recognizable components) to those of normal pulsars. However, there is a tail to the distribution in MSPs exemplified by PSR J0473$-$4715 \\citep{nms+97}, which has at least 10 components, and PSR J2124$-$3358 (Figure~\\ref{fg:2124}; at least 12 components) which is not present in normal pulsars; no normal pulsars have profiles as complex as these, even in data with high signal-to-noise ratio.  While many MSP profiles have the appearance of `stretched' normal pulsar profiles, there are significant differences. One of the most important is the lack of frequency evolution in the spacing of pulse components \\citep{kll+99}. In most normal pulsars, the pulse components are more widely separated at lower radio frequencies, an observation often interpreted in terms of `radius-to-frequency mapping' \\citep[e.g.][]{cor78,tho91a}. The results presented here confirm this lack of frequency evolution in MSP profiles. For example, Figure~\\ref{fg:2124_alpha} shows that for PSR J2124$-$3358, although components are mostly overlapping, there is a clear correspondence of component locations at the two frequencies right across the profile. It is important to note however that this property is shared by normal pulsars with so-called `interpulse' emission, including such well-known examples as the Crab pulsar \\citep{mh99} and PSR B0950+08 \\citep{hc81}, suggesting a close link between the emission processes in MSPs and `normal' interpulse pulsars.  Most MSP profiles do not fit easily into the `core-cone' classification scheme \\citep{ran83,lm88} or its extensions to multiple cones \\citep[e.g.][]{md99}. Even when the profile has a shape similar to a classical `triple', for example PSR J2145$-$0750 (Figure~\\ref{fg:2145}), the spectral properties of the central and outer components do not follow the patterns established for normal pulsars. For PSR J2124$-$3358 there is no identifiable `core' component and there is no pattern suggesting multiple nested conal emission. While there do appear to be discrete quasi-gaussian components within the observed profile, a very large number of them would be required to model it and their amplitude and phase (longitude) shows no clear pattern. Rather, the profile seems more consistent with a random distribution of emission patches as advocated by \\citet{lm88}. The spectral index variations shown in Figure~\\ref{fg:2124_alpha} are consistent with different pulse components originating in spatially distinct emission regions.  The wide observed profiles of MSPs give a good opportunity to fit the RVM and hence to determine both the magnetic inclination $\\alpha$ and the line-of-sight inclination $\\zeta$. Most normal pulsars have narrow profiles and only the impact parameter $\\beta = \\zeta - \\alpha$ can be measured \\citep[cf.][]{lm88}. In some MSPs, the observed PA variations follow the RVM quite well, especially at higher frequencies \\citep[e.g.][]{stc99,lzb+00}. However, in others, there are clear and sometimes large deviations \\citep[e.g.][]{aptw96,nms+97}. Two pulsars in the present sample, PSRs J2124$-$3358 and 2145$-$0750, have sufficiently wide profiles to justify an attempt to fit the RVM. The results are shown in Figure~\\ref{fg:pafits}. \\begin{figure*} \\psfig{file=f11.eps,width=160mm,angle=270} \\caption{Rotating-vector model fits to observed position-angle variations for two pulsars.} \\label{fg:pafits} \\end{figure*}  While the overall PA variation for PSR J2124$-$3358 is approximately described by the RVM fit, it is clear that there are large and systematic deviations from the model fit. The parameters of this fit are $\\alpha = 48\\degr\\pm 3\\degr$ and $\\zeta = 67\\degr\\pm 5\\degr$, suggesting that emission is seen from both poles with impact parameters $\\beta = 19\\degr$ and $65\\degr$ respectively. However, there is a large covariance between $\\alpha$ and $\\zeta$ with only a small increase in $\\chi^2$ for $\\alpha$ values between $20\\degr$ and $60\\degr$. At these limits, the corresponding $\\zeta$ values are $27\\degr$ and $80\\degr$ respectively. For small $\\alpha$, a one-pole interpretation is favoured. Given the poor quality of the fit, it is not at all clear what the true inclination angles for the two axes are. Certainly, the fact that emission is seen over most if not all of the pulse period suggests a one-pole, almost aligned model.  For PSR J2145$-$0750, if we assume there is an orthogonal flip near longitude $120\\degr$ but not at $195\\degr$, we get the fit shown in Figure~\\ref{fg:pafits}. This has $\\alpha = 145\\degr\\pm 12\\degr$ and $\\zeta = 148\\degr\\pm 12\\degr$. Again, there are systematic deviations from the fit and a large covariance between $\\alpha$ and $\\zeta$, but the impact parameter $\\beta \\simeq 3\\degr$ is relatively well determined. This PA fit, the sense reversal of the circular polarization and the spectral behaviour all suggest that, despite the profile morphology, the trailing component is in fact central to the emission beam \\citep[cf.][]{stc99}. If we assume a second orthogonal jump at $195\\degr$ the impact parameter increases to about $15\\degr$. However, this introduces a significant discontinuity in PA at this longitude, suggesting that this is not the correct interpretation.  If we accept that the trailing component is in fact central to the emission beam, the interpretation of the weak leading component becomes something of a problem. \\citet{xkj+98} suggest that it is a `precursor', analogous to that in the Crab pulsar. This is supported by its nearly 100\\% linear polarization. The emission bridge between the leading and strongest components suggests that all components are closely related and probably that all the emission orginates from one pole. This implies a very large effective pulse width, $2\\Delta\\phi \\sim 300\\degr$. For small $\\beta$, $\\sin(\\rho/2) \\simeq \\sin(\\Delta\\phi/2)\\sin(\\alpha/2)$, where $\\rho$ is the true beam radius \\citep[cf.][]{lm88}, giving $\\rho \\sim 68\\degr$. This compares with a predicted value of $49\\degr$ from the relation $\\rho = 6\\fdg2\\;P^{-1/2}$ \\citep{big90b}.  For other pulsars in the sample, and indeed for most MSPs \\citep{xkj+98,stc99}, there is very little swing in PA across the the observed pulse components.  Taken at face value, this implies large impact parameters and/or very small magnetic inclinations. Yet many of these pulsars have relatively narrow pulses, e.g. PSRs J0621+1002, J1518+4904 and J1713+0747 \\citep{stc99}. This implies a beam radius in the longitude direction much less than the apparent impact parameter, that is, a beam effectively elongated in the latitude direction \\citep[cf.][]{nv83}. If these pulse components are part of a wide beam from magnetic field lines associated with one pole, as suggested for young `interpulse' pulsars by \\citet{man96}, intermediate magnetic inclinations and large impact parameters are possible.  In this case, the implied emission altitudes relative to the light cylinder are much higher than those implied by observed pulse widths for normal pulsars. At these altitudes, there will be significant deviations of the magnetic field from a pure dipole form due to displacement and physical currents, providing a possible reason for the observed large PA deviations from RVM fits seen in millisecond pulsars.  It is also possible that these irregular PA variations result from very different magnetic field structures in MSPs. For example, \\citet{rud91b} suggests that, during the spindown and subsequent recycling process, crustal plate movements result in highly distorted field structures. Accretion-induced field decay \\citep[e.g.][]{rom90} would also result in highly non-dipolar fields. Such field structures could also disguise or destroy the relatively consistent patterns of core and conal emission seen in most longer-period pulsars.  In general, our polarization results agree well with those of \\citet{stc99} when account is taken of the reversed PA convention in that paper. However, there are many discrepancies with results in the \\citet{xkj+98} paper. It appears that their sign of Stokes $V$ is systematically reversed. It is also probable that the higher and approximately constant fractional linear polarization with nearly constant PA seen in a number of their figures results from errors in calibration rather than time variations in the polarization properties \\citep[cf.][]{stc99}."
    },
    "0403/astro-ph0403685_arXiv.txt": {
        "abstract": "More than half of all low-redshift AGN exhibit UV and X-ray absorption by highly ionized gas. The observed UV and X-ray absorption lines are almost always blue-shifted at velocities of hundreds of km/s, indicating that the absorbing gas is outflowing from the active nucleus. In some cases the inferred mass flux rivals the Eddington limit of the central black hole, an indication that these outflows are intimately related to the mass accretion and energy generation mechanism in AGN. The ejected material can also have an affect on the interstellar medium of the host galaxy and the surrounding intergalactic medium. Over the past several years, coordinated UV and X-ray observations of several bright AGN at high spectral resolution using {\\it HST}, {\\it FUSE}, {\\it Chandra}, and {\\it XMM-Newton} have contributed greatly to our understanding of these outflows. I will give an overview of these recent observations, summarize our {\\it FUSE} survey of low-redshift AGN, and interpret the results in the context of models of winds from accretion disks and thermally driven winds from the obscuring torus. ",
        "introduction": "Mass outflows from active galactic nuclei (AGN) can profoundly affect the evolution of the central engine \\citep{BB99}, the host galaxy and its interstellar medium \\citep{SR98, WL03} and also the surrounding intergalactic medium (IGM) \\citep{Cavaliere02, Granato04, SO04}.  Winds from the high metal-abundance nuclear regions may be a significant source for enriching the IGM \\citep{Adelberger03}. Absorption by the outflow can also collimate the ionizing radiation \\citep{Kriss97} and thereby influence the ionization structure of the host galaxy and the surrounding IGM.  The outflowing gas in AGN is sometimes visible as extended, bi-conical emission at visible or X-ray wavelengths (e.g., NGC 4151, \\citealt{Evans93}, \\citealt{Hutchings98}; NGC 1068, \\citealt{Ogle03}), but it most frequently manifests itself as blue-shifted absorption features in their UV and X-ray spectra. About half of all low-redshift AGN show X-ray absorption by highly ionized gas \\citep{Reynolds97, George98}, and a similar fraction show associated UV absorption in ionized species such as {\\sc C~iv} \\citep{Crenshaw99} and {\\sc O~vi} \\citep{Kriss01}. In more luminous quasars, the fraction of objects in the Sloan Digital Sky Survey that shows broad {\\sc C~iv} absorption troughs rises steeply as the troughs become narrower \\citep{Tolea02,Reichard03}, comprising over 30\\% of the quasar population at widths narrower than 1000 \\kms. For AGN that have been observed in both the X-ray and the UV, there is a one-to-one correspondence between objects showing X-ray and UV absorption, implying that the phenomena are related in some way (Crenshaw et al. 1999). The high frequency of occurrence of UV and X-ray absorption suggests that the absorbing gas has a high covering fraction, and that it is present in all AGN. The gas has a total mass exceeding $\\sim 10^3~\\rm M_\\odot$ (greater than the broad-line region, or BLR), and is outflowing at a rate $>0.1~\\rm M_\\odot~yr^{-1}$ ($10 \\times$ the accretion rate in some objects) \\citep{Reynolds97}.  Key questions for understanding the outflowing, absorbing gas in AGN are:  \\vspace{-12pt} \\begin{list}{} % {\\listparindent 0pt \\leftmargin 0pt \\rightmargin 0pt \\itemindent -0pt\\parsep 0pt\\labelsep 0pt\\itemsep 0pt} \\begin{sloppy}  \\smallskip \\item {$\\bullet~~~$} What are the column densities and ionic abundances in the absorbing gas?\\\\ Even such a basic question is uncertain since the absorption line profiles are complex. Doublet ratios show that the absorbers can be optically thick, but they are not black at line center. Thus column densities are frequently underestimated, sometimes by as much as an order of magnitude \\citep{Arav02, Arav03}.  \\smallskip \\item {$\\bullet~~~$} Where is the absorbing gas located?\\\\ The location of the absorbers is a vital clue to the process producing the outflow. Winds arising from an accretion disk \\citep{KK94, Murray95, Elvis00, Proga00} will have material at a broad range of radii reaching from near the disk itself to beyond the BLR. Thermally driven winds arising from the obscuring torus \\citep{KK95, KK01} will lie at much larger radii, typically at distances of $\\sim 1$ pc. The radial location also determines the mass flux---the larger the distance, the higher the total mass and the mass flux in the wind.  \\smallskip \\item {$\\bullet~~~$} Do variations reflect an ionization response, or are they due to bulk motion?\\\\ In the limited number of monitoring campaigns carried out thus far, examples of both have been seen. The neutral hydrogen and the {\\sc C~iii} absorption in NGC~4151 responds quite clearly to continuum variations \\citep{Kriss97, Espey98}. Some absorbing clouds in NGC~3783 have shown changes consistent with an ionization response, while others appear to be due to bulk motion \\citep{Crenshaw99, Gabel03b}.  \\smallskip \\item {$\\bullet~~~$} How are the X-ray and UV absorption related?\\\\ In some cases, UV absorbing gas may be directly associated with the X-ray warm absorber (3C351: \\citealt{Mathur94}; NGC~5548: \\citealt{Mathur95}; NGC 3516: \\citealt{Kraemer02}). In the extensive recent {\\it Chandra/FUSE/HST} campaign on NGC 3783 \\citep{Kaspi02, Gabel03a}, the kinematics of the X-ray absorbing gas are also a good match to the UV-absorbing gas. In other cases, however, the UV gas appears to be in an even lower ionization state, and there is no direct relation between the X-ray absorption and the multiple kinematic components seen in the UV (NGC~4151: \\citealt{Kriss95}, \\citealt{Kraemer01}; NGC~3516: \\citealt{Kriss96a, Kriss96b}; NGC~5548: \\citealt{Mathur99}, \\citealt{Crenshaw03}; Mrk~509: \\citealt{Kriss00b}, \\citealt{Yaqoob03}; NGC~7469: \\citealt{Kriss00a, Kriss03, Blustin03}). The X-ray absorbing gas itself contains material spanning a large range of ionization parameters \\citep{Lee02, Sako03, Netzer03}, and it is likely that this broad range of physical conditions can also include the UV-absorbing ions. This is a natural prediction of the thermally driven wind model of \\citet{KK95, KK01}, and would also be likely in disk-driven winds.  \\end{sloppy} \\end{list}  \\vspace{-8pt} To address some of these questions, we have been conducting a survey of the $\\sim100$ brightest AGN using the {\\it Far Ultraviolet Spectroscopic Explorer (FUSE)}. The short wavelength response (912--1187 \\AA) of {\\it FUSE} \\citep{Moos00} enables us to make high-resolution spectral measurements ($R \\sim 20,000$) of the high-ionization ion {\\sc O~vi} and the high-order Lyman lines of neutral hydrogen. As of November 1, 2002, we have observed a total of 87 AGN. Of these, 57 have $z < 0.15$, so that the {\\sc O~vi} doublet is visible in the {\\it FUSE} band. The \\ovi\\ doublet is a crucial link between the higher ionization absorption edges seen in the X-ray and the lower ionization absorption lines seen in earlier UV observations. The high-order Lyman lines provide a better constraint on the total neutral hydrogen column density than Ly$\\alpha$ alone. Lower ionization species such as {\\sc C~iii} and {\\sc N~iii} also have strong resonance lines in the \\FUSE\\ band, and these often are useful for setting constraints on the ionization level of any detected absorption. The Lyman and Werner bands of molecular hydrogen also fall in the \\FUSE\\ band, and we have searched for (but not found) intrinsic $\\rm H_2$ absorption that may be associated with the obscuring torus.  \\begin{figure}[h] \\plottwo{gkrissf1a.eps}{gkrissf1b.eps} \\caption{ Left: Histogram of FUSE AGN versus luminosity. The shaded area shows the number of objects exhibiting intrinsic absorption. Right: The points show outflow velocity as a function of luminosity. \\label{fig1}} \\end{figure} ",
        "conclusions": "The multiple kinematic components frequently seen in the UV absorption spectra of AGN clearly show that the absorbing medium is complex, with separate UV and X-ray dominant zones. In some cases, the UV absorption component corresponding to the X-ray warm absorber can be clearly identified (e.g., Mrk~509, \\citealt{Kriss00b}). In others, however, {\\it no} UV absorption component shows physical conditions characteristic of those seen in the X-ray absorber (NGC~3516, \\citealt{Kriss96a, Kriss96b}; NGC~5548, \\citealt{Brotherton02}). One potential geometry for this complex absorbing structure is high-density, low-column UV-absorbing clouds embedded in a low-density, high-ionization medium that dominates the X-ray absorption.  Disk-driven winds are a possible explanation for some cases of AGN outflows. By analogy to stellar winds, one would expect the terminal velocity of an AGN outflow to reflect the gravity of its origin. Disk-driven winds should therefore have velocities in the range of several thousand \\kms. Objects with broad, smooth profiles might fall in this category. The geometry proposed by \\citet{Elvis00} suggests that these objects should have only modest inclinations. However, two prime examples of Seyferts with broad smooth absorption troughs, NGC 3516 \\citep{Hutchings01} and NGC 4151 \\citep{Kriss01}, are likely the highest inclination sources in our sample given their extended, bi-conical narrow emission-line region morphologies \\citep{Miyaji92, Evans93} and their opaque Lyman limits \\citep{Kriss97}.  The lower velocities we observe in objects like NGC 3783 and NGC 5548 are more compatible with thermally driven winds from the obscuring torus \\citep{KK95, KK01}. In these thermally driven winds, photoionized evaporation in the presence of a copious mass source (the torus) locks the ratio of ionizing intensity to total gas pressure (the ionization parameter $\\Xi$) at a critical value. For AGN spectral energy distributions lacking a strong extreme ultraviolet bump, such as the composite spectra of quasars assembled by \\citet{Zheng97}, \\citet{Laor97}, and \\citet{Telfer02}, the ionization equilibrium curve exhibits an extensive vertical branch. Thus, at the critical ionization parameter for evaporation, there is a broad range of temperatures that can coexist in equilibrium at nearly constant pressure. For this reason, the flow is expected to be strongly inhomogeneous. Outflow velocities are typical of the sound speed in the heated gas, or several hundred \\kms, comparable to the velocities seen in many AGN.  In summary, we find that \\ovi\\ absorption is common in low-redshift ($z < 0.15$) AGN. 30 of 53 Type 1 AGN with $z < 0.15$ observed using \\FUSE\\ show multiple, blended \\ovi\\ absorption lines with typical widths of $\\sim 100~\\kms$ that are blueshifted over a velocity range of $\\sim$ 1000 \\kms. Those galaxies in our sample with existing X-ray or longer wavelength UV observations also show {\\sc C~iv} absorption and evidence of a soft X-ray warm absorber. In some cases, a UV absorption component has physical properties similar to the X-ray absorbing gas, but in others there is no clear physical correspondence between the UV and X-ray absorbing components."
    },
    "0403/astro-ph0403016_arXiv.txt": {
        "abstract": "X-ray and Sunyaev-Zeldovich Effect data can be combined to determine the distance to galaxy clusters. High-resolution X-ray data are now available from the {\\it Chandra} Observatory, which provides both spatial and spectral information, and Sunyaev-Zeldovich Effect data were obtained from the BIMA and OVRO arrays. We introduce a Markov chain Monte Carlo  procedure for the joint analysis of X-ray and Sunyaev-Zeldovich Effect data. The advantages of this method  are the high computational efficiency and the ability to measure simultaneously the probability distribution of all parameters of interest, such as the spatial and spectral properties of the cluster gas and also for derivative quantities such as the distance to the cluster. We demonstrate this technique by applying it to the {\\it Chandra} X-ray data and the OVRO radio data for the galaxy cluster Abell~611. Comparisons with traditional likelihood-ratio methods reveal the robustness of the method. This method will be used in  follow-up papers to determine the distances to a large sample of galaxy clusters. ",
        "introduction": "Analysis of  Sunyaev-Zeldovich Effect (SZE) and X-ray data provides a unique method of directly determining distances of galaxy clusters. Clusters of galaxies contain hot plasma ($k_B T_e \\sim$ 2-20 keV) that scatters the cosmic microwave background radiation (CMB). On average,  this inverse Compton scattering boosts the energy of the CMB photons, causing a small distortion in the CMB spectrum (Sunyaev and Zeldovich 1970,1972). For recent reviews of the SZE and its application for cosmology see Birkinshaw (1999) and Carlstrom et al. (2002).  The SZE is proportional to the integrated pressure along the line of sight, $\\Delta T \\propto  \\int n_e T_e dl$, where $n_e$ and $T_e$ are the electron density and electron temperature of the cluster plasma. The thermal X-ray emission from the same plasma has a different dependence on the density, $S_x \\propto \\int n_e^2 \\Lambda_{ee} dl$, where $\\Lambda_{ee}$ is the X-ray cooling function. Making assumptions on the distribution of the plasma (e.g., a $\\beta$ profile) and taking advantage of the different dependence on $n_e$, SZE and X-ray observations can be combined to determine the distance to galaxy clusters. These cluster distances  can be combined with redshift measurements to determine the value of the Hubble constant (Myers et al. 1997, Grainge et al. 2002, Reese et al. 2002).     In this paper we introduce a Markov chain Monte Carlo (MCMC) procedure for the joint analysis of SZE and X-ray data. The method is tested on the galaxy cluster Abell 611 which has SZE data obtained with the Caltech millimeter interferometric array at the Owens Valley Radio Observatory (OVRO) outfitted with centimeter-wave receivers (Carlstrom, Joy and Grego 1996) and X-ray data from the {\\it Chandra} X-ray Observatory. In a subsequent paper we will report the application of this technique to a sample of $\\sim$40 clusters. ",
        "conclusions": "We present a Markov chain Monte Carlo technique to derive cluster distances from SZE and X-ray data. The method was succesfully tested on the OVRO and {\\it Chandra} data of Abell~611, a galaxy cluster at z=0.288. We measure an angular diameter distance of $D_A=1.00\\pm^{0.24}_{0.21}$ Gpc (68\\% confidence level). In a previous work based on the same SZE data, but using lower-resolution X-ray data, Reese et al. (2002) derived a distance of $D_A=0.99\\pm^{0.32}_{0.29}$ Mpc.  The {\\it Chandra} X-ray data provides simultaneous spatial and spectral information, featuring the finest angular resolution to date ($\\sim 0.5$ arcsec half-power radius). The MCMC method has two major advantages: it is computationally more efficient than the traditional likelihood ratio-based methods and it provides simultaneously the probability distribution function of all model parameters.  This technique will be used in future papers to determine the distances of a large sample of galaxy clusters for which there are available high-resolution {\\it Chandra} X-ray data and BIMA/OVRO SZE data.   This work is supported by NASA LTSA grant NAG 5-7986. E.D.R. acknowledges support from NASA Chandra Postdoctoral Fellowship PF 1-20020. Partial support was also provided by NSF grants AST-0096913 and PHY-0114422. S.J.L.\\ acknowledges support from NASA GSRP Fellowship  NGT5-50173. We thank the referee for  helpful comments and suggestions.   \\newpage"
    },
    "0403/astro-ph0403220_arXiv.txt": {
        "abstract": "We use the lag-luminosity relation to calculate self-consistently the redshifts, apparent peak bolometric luminosities $L_B$, and isotropic energies $E_{\\rm iso}$ for a large sample of BATSE gamma-ray bursts.  We consider two different forms of the lag-luminosity relation; for both forms the median redshift for our burst database is 1.6. We model the resulting $E_{\\rm iso}$ sample with power law and Gaussian probability distributions without redshift evolution, both of which are reasonable models.  The power law model has an index of $\\alpha_E=1.76\\pm0.05$ (95\\% confidence), where $p(E_{\\rm iso}) \\propto E_{\\rm iso}^{-\\alpha_E}$.  The simple universal jet profile model suggested but did not require $\\alpha_E=2$, and subsequent physically reasonable refinements to this model permit greater diversity in $\\alpha_E$, as well as deviations from a power law; therefore our observed $E_{\\rm iso}$ probability distribution does not disprove the universal jet model. ",
        "introduction": "The major breakthroughs in the study of gamma-ray bursts of the past six years---most if not all bursts are cosmological, bursts do {\\it not} have constant peak luminosity, the fireballs are beamed, many bursts are associated with supernovae---resulted from the intensive study of a relatively small number of bursts without regard for whether these bursts formed a well defined statistical sample.  However, realizing that for any one burst we are only sampling an anisotropic radiation pattern from one direction, we now have to study distributions of burst properties to reconstruct the bursts' appearance from all directions. Well-defined burst samples are required to derive these distributions.  Unfortunately, we do not yet have a large sample of bursts with spectroscopic redshifts, which in most cases are required for calculating the intrinsic burst properties. The two dozen or so bursts with redshifts were detected by various detectors with different detection thresholds and energy sensitivities, and the follow-ups that determined the redshifts depended on the vagaries of weather and telescope availability. BATSE provided a large burst sample with well-understood thresholds, but without direct redshift determinations for most of these bursts. However, redshifts can be determined indirectly from the lag-luminosity (Norris, Marani \\& Bonnell 2000) or variability-luminosity (Fenimore \\& Ramirez-Ruiz 2000; Reichart et al. 2001) relations. Here we calculate self-consistently the redshifts for a large subset of the BATSE bursts using the lag-luminosity relation.  To accommodate the problematic burst GRB980425, which is significantly underluminous if the lag-luminosity relation is a single power law, Salmonson (2001) and Norris (2002) proposed that the relation should be a broken power law. This reduces the luminosity of many long-lag bursts, moving this population closer to the observer.  We use the resulting redshifts to calculate both the peak bolometric luminosity $L_B$ and the isotropic energy $E_{\\rm iso}$.  Bursts are thought to radiate anisotropically, but we sample their radiation field in only one direction.  Therefore $L_B$ and $E_{\\rm iso}$ are calculated from the observed flux as if the emissions are isotropic; corrections for the anisotropy are based on models of the relativistic jets that emit the observed gamma rays (see Bloom, Frail \\& Kulkarni 2003).  $L_B$ and $E_{\\rm iso}$ are both bolometric quantities in the burst frame.  $L_B$ is the maximum value of the luminosity while $E_{\\rm iso}$ is the total energy emitted over the duration of the burst (without correcting for the anisotropic emission).  We consider the isotropic energy $E_{\\rm iso}$ to be more fundamental, and therefore model its probability distribution.  Two models have been proposed for the structure of the jets. The uniform jet model (Frail et al. 2001; Bloom et al. 2003) assumes that all jets have a constant surface energy density $\\epsilon$ (energy per solid angle across the jet) but differing opening angles $\\theta_0$ (the angle between the jet axis and the edge of the jet). The model makes no predictions for the energy probability distribution.  On the other hand, in the universal jet profile model (Rossi, Lazzati, \\& Rees 2002; Zhang \\& Meszaros 2002) all jets have the same surface energy density $\\epsilon$ as a function of off-axis angle $\\theta$ (the angle from the jet axis), and thus the observed differences in $L_B$ or $E_{\\rm iso}$ result from the angle $\\theta_v$ between the jet axis and the line-of-sight. This model predicts that the energy probability distribution is a power law with index $\\alpha_E$, where $p(E_{\\rm iso}) \\propto E_{\\rm iso}^{-\\alpha_E}$. If $\\epsilon \\propto \\theta^k$, then $\\alpha_E=1-2/k$. Rossi et al. (2002) and Zhang \\& Meszaros (2002) suggested $k=-2$, resulting in $\\alpha_E=2$, to reproduce the distributions observed by Frail et al. (2001), although this was not a firm prediction. A Gaussian surface energy density profile results in $\\alpha_E=1$ (Lloyd-Ronning, Dai \\& Zhang 2004). Thus $\\alpha_E$ is of order 1--2. Lloyd-Ronning et al. (2004) showed that allowing the parameters of the jet profile to vary somewhat from burst to burst---the jet profile is now only quasi-universal---results both in variations in the value of $\\alpha_E$ and deviations from a pure power law. While numerical modelling of hypernovae (Zhang, Woosley, \\& MacFadyen 2003) shows that the surface energy density of the outflows should vary with off-axis angle, it does not predict the profile; thus the profiles are fitted empirically to the data.  Consequently, the observed $E_{\\rm iso}$ probability distribution is not a powerful discriminant between the two jet structure models.  This paper has two goals.  First, we calculate the redshift of a large sample of bursts self-consistently from the lag-luminosity relation.  With the redshift and the observed fluxes and fluences we calculate $L_B$ and $E_{\\rm iso}$. Second, we model the $E_{\\rm iso}$ distribution with two functional forms:  power law and lognormal.  We find the best fit values for the parameters of each functional form, and evaluate how well each functional form describes the data. We also compare the fit of the power law functional form to the predictions of the universal jet profile model.  In these calculations we use a cosmology with $H_0 = 70$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_m = 0.3$ and $\\Omega_\\Lambda = 0.7$.  The notation $p(a\\,|\\,b)$ means the probability of $a$ given $b$.  Lower case $p$ denotes a probability density while upper case $P$ (without a subscript) represents a cumulative probability.  Note that $a$ and $b$ are propositions that are true or false. A proposition can be a statement such as ``the energy distribution can be described with a lognormal functional form'' or a parameter value (equivalent to the statement ``the parameter value is 5'').  In \\S 2 we present the methodology used in this study: calculating the burst redshifts (\\S 2.1); fitting the energy probability distribution (\\S 2.2); and using the cumulative probability to test the quality of the fit (\\S 2.3).  The implementation of this methodology is discussed in \\S 3 and the results are provided by \\S 4. Finally, our conclusions are in \\S 5. ",
        "conclusions": "In this paper we have two objectives.  First, we use the lag-luminosity relation to calculate self-consistently the redshifts for 1218 BATSE bursts. For the bursts without the spectral parameters required by the calculation we use average low and high energy spectral indices, and a peak energy $E_p$ derived from the hardness ratio.  We find that the redshift is quite sensitive to the spectral parameters, particularly if the spectrum is sharply peaked. We use both single power law and broken power law lag-luminosity relations, and find that the broken power law relation does indeed predict a population of low luminosity, nearby bursts.  For both forms of the relation the median redshift is 1.58.  We use the redshifts to calculate the apparent peak bolometric flux and the isotropic energy, both assuming that the bursts radiate isotropically.  Second, we fit two functional forms to the distribution of the isotropic energy.  We find that our burst data can be fit by a power law energy distribution with $\\alpha_E=1.76\\pm0.05$ (95\\% confidence); considering the likely systematic uncertainties in addition to the statistical variance, the power law distribution is probably a good description of the data.  This value of $\\alpha_E$ is in the acceptable range for the universal jet profile model, and therefore, our work does not distinguish between the current jet structure models.  A lognormal energy distribution also describes the data; the data permit a smaller average energy if the distribution is wider.  More faint bursts will bound the lower end of this distribution."
    },
    "0403/astro-ph0403546_arXiv.txt": {
        "abstract": " ",
        "introduction": "One of the key puzzles in X-ray astronomy is to understand accretion flows in a strong gravitational field. This applies to both Active Galactic Nuclei (AGN) and Quasars, where the accretion is onto a supermassive black hole, the puzzling ultraluminous compact sources which may be intermediate mass black holes, the stellar mass Galactic black holes (GBH) and even the neutron star systems.  Neutron star radii are of order three Schwarzchild radii, i.e. the same as that for the last stable orbit of material around a black hole. Thus they have very similar gravitational potentials so should have very similar accretion flows, though of course with the major difference that neutron stars have a solid surface, so can have a boundary layer and a stellar magnetic field.  The main premise is that progress in understanding accretion in {\\em any} of these objects should give us some pointers to understanding accretion in {\\em all} of them. Galactic sources are intrinsically less luminous, but a great deal closer than the AGN, so generally are much brighter. Galactic black holes are also often generically transient, showing large variability on timescales from milliseconds to years.  These give us a sequence of spectra at differing mass accretion rates onto the central object, allowing us to test accretion models. ",
        "conclusions": "We can form a unified picture of the spectral evolution of all types of low mass X-ray binaries, where the major hard-soft spectral transition is driven by a changing inner disc radius linked to the collapse of an optically thin, hot inner flow. We see clear differences between the black holes and disc accreting neutron stars, both in the form of a unique black hole spectral signature (the high/soft state) and in terms of the evolution of their spectral shape with $L/L_{Edd}$. These differences can be modeled qualitatively and quantitatively as the {\\em same} accretion flow onto a {\\em different} object: neutron stars have a surface so have a boundary layer, while black holes have an event horizon!"
    },
    "0403/astro-ph0403400_arXiv.txt": {
        "abstract": "Calculations of weak-interaction transition rates and of nuclear formation enthalpies show that in isolated neutron stars, the solid phase, above the neutron-drip threshold, is amorphous and heterogeneous in nuclear charge.  The neutrino emissivities obtained are very dependent on the effects of proton shell structure but may be several orders of magnitude larger than the electron bremsstrahlung neutrino-pair emissivity at temperatures $\\sim 10^{9}$ K. In this phase, electrical and thermal conductivities are much smaller than for a homogeneous {\\it bcc} lattice. In particular, the reduced electrical conductivity, which is also temperature-independent, must have significant consequences for the evolution of high-multipole magnetic fields in neutron stars. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403636_arXiv.txt": {
        "abstract": "{We have surveyed the metal-rich spiral galaxy M~83 (NGC~5236) for its Wolf-Rayet population using VLT-FORS2 narrow-band imaging and follow-up spectroscopy. From a total of 280 candidates identified using He\\,{\\sc ii} $\\lambda$4686 imaging, Multi Object Spectroscopy of 198 sources was carried out, revealing 132 objects containing bona-fide Wolf-Rayet features. From this sample, an exceptional W-R content of $\\sim$1030 is inferred, with  N(WC)/N(WN)$\\sim$1.3, continuing the trend to larger values at higher metallicity amongst Local Group galaxies. More dramatic is the dominance of late-type  WC stars in M~83 with N(WC8--9)/N(WC4--7)=9 which we attribute to the sensitivity of the classification line  C\\,{\\sc iii} $\\lambda$5696 to mass-loss, providing the strength of WC winds scale with metallicity. One young massive compact cluster, \\#74 in our catalogue, hosts 20\\% of the entire galactic  population, namely $\\sim$180 late WC stars  and $\\sim$50  late WN stars. ",
        "introduction": "Wolf-Rayet stars represent excellent diagnostics of recent star formation in starbursts. Though few in number, they possess very powerful winds and so contribute significantly to the energy released into the ISM in young starbursts (Crowther \\& Dessart 1998), and are believed to represent the immediate precursors of Type Ib/c Supernovae and long duration Gamma-Ray Bursts. Individual WR stars are readily identified in nearby galaxies via their characteristic strong, broad emission lines, whilst their collective presence can be seen in the integrated light of distant starbursts.  Mass-loss rates of hot, massive stars prior to the WR phase are known to scale with metallicity (Vink et al. 2001), such that the relative number of WR stars increases with higher metallicity (Maeder \\& Meynet 1994).  In addition, the distribution amongst nitrogen-rich WN versus carbon-rich WC subclasses is also known to vary with metallicity  for Local Group galaxies (Massey \\& Johnson 1998). Studies of resolved WR populations at metallicities higher than the Solar value have proved difficult to date because the inner Milky Way is visibly obscured, whilst the other metal-rich Local Group member, M31, lies at an unfavourable inclination.   Towards this goal, M~83 (NGC~5236) is a massive spiral galaxy that is believed to be particularly metal-rich: log(O/H)+12=9.0--9.3 as derived from strong-line nebular methods (Bresolin \\& Kennicutt 2002)\\footnote{However, see Kennicutt et al. (2003) regarding the reliability of strong nebular line techniques at high metallicity.} and is the focus of the present Letter. A more detailed discussion of the WR population of M~83 and presentation of the full catalogue is given by Hadfield et al. (in preparation). ",
        "conclusions": "\\label{sect5}  We have surveyed the nearby metal-rich galaxy M~83 for the presence of WR stars using VLT-FORS2. From follow up spectroscopy of 198 out of 280 candidates, 132 regions containing WR stars are identified. Assuming intrinsic line fluxes comparable to Galactic counterparts (Schaerer \\& Vacca 1998), we identify 1030 WR stars in M~83, with N(WC)/N(WN)$\\sim$1.3, continuing the trend observed amongst Local Group galaxies to higher metallicity (Massey \\& Johnson 1998). Accounting for the remaining candidates the total WR population may be as high as 1500. More than 50\\% of the known WR are identified as late WC stars, versus $\\leq$5\\% early WC stars, which is unprecedented relative to more metal-poor Local Group galaxies. One young massive compact cluster, \\#74, hosts $\\sim$230 late WN and WC stars.  The relatively large WC population with increasing metallicity can readily be understood from comparison with evolutionary models (Maeder \\& Meynet 1994). At higher metallicity, mass-loss rates during and subsequent to the main-sequence evolution of massive stars strips away higher layers earlier on, such that a star with a particular initial mass advance to later (WC) phases at higher metallicity.  But why are there exclusively {\\it early } WC stars in the LMC, a mixed population in the Milky Way and almost exclusively {\\it late} WC stars in M~83? A decade ago it was believed that (C+O)/He increases from WC9 to WC4 (Smith \\& Maeder 1991). However, subsequent spectral analysis failed to confirm any systematic trend in C/He with subtype (Koesterke \\& Hamann 1995; Crowther et al. 2002). If carbon abundances do not play a dominant role, what does? We suggest differences in wind densities are primarily responsible. The wind densities of WC stars in the LMC  are  $\\sim$50\\% lower than those of their  Milky Way counterparts according to Crowther et al. (2002), which they attributed to a metallicity dependence of their winds. Since C\\,{\\sc iii} $\\lambda$5696 increases dramatically in strength with increasing mass-loss rate (see their Fig.12), one would expect yet higher mass-loss rates and even stronger C\\,{\\sc iii} $\\lambda$5696 in M~83 WC stars if WC winds are metallicity dependent. Our present results are fully consistent with such a population. A genuine metallicity dependence of WR winds has implications for the hard ionizing fluxes of young starbursts (Smith et al. 2002). Regardless, the presence of late WC stars is definitely a indicator of a metal-rich environment.  Finally,  should  a supernova be observed in M~83 in the near future, we now possess  a reasonable  statistical sample with which we should be able to verify whether a WR star was a potential immediate precursor."
    },
    "0403/gr-qc0403068.txt": {
        "abstract": "Post-Newtonian relativistic theory of astronomical reference frames based on Einstein's general theory of relativity was adopted by General Assembly of the International Astronomical Union in 2000. This theory is extended in the present paper by taking into account all relativistic effects caused by the presumable existence of a scalar field and parametrized by two parameters, $\\beta$ and $\\gamma$, of the parametrized post-Newtonian (PPN) formalism. We use a general class of the scalar-tensor (Brans-Dicke type) theories of gravitation to work out PPN concepts of global and local reference frames for an astronomical N-body system. The global reference frame is a standard PPN coordinate system.  A local reference frame is constructed in the vicinity of a weakly self-gravitating body (a sub-system of the bodies) that is a member of the astronomical N-body system. Such local inertial frame is required for unambiguous derivation of the equations of motion of the body in the field of other members of the N-body system and for construction of adequate algorithms for data analysis of various gravitational experiments conducted in ground-based laboratories and/or on board of spacecrafts in the solar system.  We assume that the bodies comprising the N-body system have weak gravitational field and move slowly. At the same time we do not impose any specific limitations on the distribution of density, velocity and the equation of state of the body's matter. Scalar-tensor equations of the gravitational field are solved by making use of the post-Newtonian approximations so that the metric tensor and the scalar field are obtained as functions of the global and local coordinates. A correspondence between the local and global coordinate frames is found by making use of asymptotic expansion matching technique.  This technique allows us to find a class of the post-Newtonian coordinate transformations between the frames as well as equations of translational motion of the origin of the local frame along with the law of relativistic precession of its spatial axes. These transformations depend on the PPN parameters $\\beta$ and $\\gamma$, generalize general relativistic transformations of the IAU 2000 resolutions, and should be used in the data processing of the solar system gravitational experiments aimed to detect the presence of the scalar field. These PPN transformations are also applicable in the precise time-keeping metrology, celestial mechanics, astrometry, geodesy and navigation.  We consider a multipolar post-Newtonian expansion of the gravitational and scalar fields and construct a set of internal and external gravitational multipoles depending on the parameters $\\beta$ and $\\gamma$. These PPN multipoles generalize the Thorne-Blanchet-Damour multipoles defined in harmonic coordinates of general theory of relativity. The PPN multipoles of the scalar-tensor theory of gravity are split in three classes -- {\\it active}, {\\it conformal}, and {\\it scalar} multipoles. Only two of them are algebraically independent and we chose to work with the {\\it conformal} and {\\it active} multipoles. We derive the laws of conservations of the multipole moments and show that they must be formulated in terms of the {\\it conformal} multipoles.  We focus then on the law of conservation of body's linear momentum which is defined as a time derivative of the {\\it conformal} dipole moment of the body in the local coordinates. We prove that the local force violating the law of conservation of the body's linear momentum depends exclusively on the {\\it active} multipole moments of the body along with a few other terms which depend on the internal structure of the body and are responsible for the violation of the strong principle of equivalence (the Nordtvedt effect).  The PPN translational equations of motion of extended bodies in the global coordinate frame and with all gravitational multipoles taken into account are derived from the law of conservation of the body's linear momentum supplemented by the law of motion of the origin of the local frame derived from the matching procedure. We use these equations to analyze translational motion of spherically-symmetric and rigidly rotating bodies having finite size. Spherical symmetry is defined in the local frame of each body through a set of conditions imposed on the shape of the body and the distribution of its internal density, pressure and velocity field. We prove that our formalism brings about the parametrized post-Newtonian EIH equations of motion of the bodies if the finite-size effects are neglected. Analysis of the finite-size effects reveal that they are proportional to the parameter $\\beta$ coupled with the second and higher-order rotational moments of inertia of the bodies. The finite-size effects in the translational equations of motion can be appreciably large at the latest stage of coalescence of binary neutron stars and can be important in calculations of gravitational waveform templates for the gravitational-wave interferometers.  The PPN rotational equations of motion for each extended body possessing an arbitrary multipolar structure of its gravitational field, have been derived in body's local coordinates. Spin of the body is defined phenomenologically in accordance with the post-Newtonian law of conservation of angular momentum of an isolated system. Torque consists of a general relativistic part and the PPN contribution due to the presence of the scalar field. The PPN scalar-field-dependent part is proportional to the difference between {\\it active} and {\\it conformal} dipole moments of the body which disappears in general relativity. Finite-size effects in rotational equations of motion can be a matter of interest for calculating gravitational wave radiation from coalescing binaries. ",
        "introduction": "\\subsection{General Conventions} Greek indices $\\alpha, \\beta, \\gamma,...$ run from 0 to 3 and mark space-time components of four-dimensional objects. Roman indices $i,j,k,...$ run from 1 to 3 and denote components of three-dimensional objects (zero component belongs to time). Repeated indices mean the Einstein summation rule, for instance, $A^\\alpha B_\\alpha=A^0B_0+A^1B_1+A^2B_2+A^3B_3$ and $T^k_k=T^1_{\\;1}+T^2_{\\;2}+T^3_{\\;3}$, etc.  Minkowski metric is denoted $\\eta_{\\alpha\\beta}={\\rm diag}(-1,+1,+1,+1)$. Kronecker symbol (the unit matrix) is denoted $\\delta_{ij}= {\\rm diag}(1,1,1)$. Levi-Civita fully antisymmetric symbol is $\\varepsilon_{ijk}$ such that $\\varepsilon_{123}=+1$. Kronecker symbol is used to rise and lower Roman indices. Complete metric tensor $g_{\\alpha\\beta}$ is used to rise and lower the Greek indices in exact tensor equations whereas the Minkowski metric $\\eta_{\\alpha\\beta}$ is employed for rising and lowering indices in the post-Newtonian approximations  Parentheses surrounding a group of Roman indices mean symmetrization, for example, $A_{(ij)}=(1/2)(A_{ij}+A_{ji})$. Brackets around two Roman indices denote antisymmetrization, that is $A_{[ij]}=(1/2)(A_{ij}-A_{ji})$. Angle brackets surrounding a group of Roman indices denote the symmetric trace-free (STF) part of the corresponding three-dimensional object, for instance, $$A_{<ij>}=A_{(ij)}-\\frac{1}{3}\\delta_{ij}A_{kk}\\;,\\quad\\quad A_{<ijk>}=A_{(ijk)} -\\frac{1}{5}\\delta_{ij}A_{kpp}-\\frac{1}{5}\\delta_{jk}A_{ipp}-\\frac{1}{5}\\delta_{ik}A_{jpp} \\;.$$ We also use multi-index notations, for example, $$A_L=A_{i_1i_2...i_l}\\;, \\qquad\\quad B_{P-1}=B_{i_1i_2...i_{p-1}}\\;, \\qquad\\qquad D_{<L>}=D_{<i_1i_2...i_l>}\\;.$$ Sum over multi-indices is understood as $$A_L Q^L=A_{i_1i_2...i_l}Q^{i_1i_2...i_l}\\;,\\qquad\\quad P_{aL-1} T^{bL-1}=P_{ai_1i_2...i_{l-1}}T^{bi_1i_2...i_{l-1}}\\;.$$ Comma denotes a partial derivative, for example, $\\phi_{,\\alpha}=\\partial\\phi/\\partial x^\\alpha$, where $\\phi_{,0}=c^{-1}\\partial\\phi/\\partial t,\\, \\phi_{,i}=\\partial\\phi/\\partial x^i,$ and semicolon $T^\\alpha_{\\,\\;;\\beta}$ denotes a covariant derivative. $L$-order partial derivative with respect to spatial coordinates is denoted $\\partial_L=\\partial_{i_1}...\\partial_{i_l}$. Other conventions are introduced as they appear in the text. We summarize these particular conventions and notations in the next section for the convenience of the readers. \\subsection{Particular Conventions and Symbols Used in the Paper} \\renewcommand{\\arraystretch}{1.5} \\setlength{\\arrayrulewidth}{0.1mm} %\\setlength{\\arraycolsep}{3mm} \\begin{longtable}{|c|@{\\,}p{7cm}@{\\,}|c|} \\hline \\textbf{\\color{blue}Symbol}&\\multicolumn{1}{c|}{\\textbf{\\color{blue}Description}}& \\multicolumn{1}{c|}{\\textbf{\\color{blue}Equation(s)}} \\endhead\\hline $g_{\\mu\\nu} $&physical (Jordan-Fierz frame) metric tensor&(\\ref{10.3})\\\\\\hline $\\tilde g_{\\mu\\nu} $&conformal (Einstein frame) metric tensor&(\\ref{13.18})\\\\\\hline $ g $&the determinant of $g_{\\mu\\nu}$ &(\\ref{10.1})\\\\\\hline $\\tilde g $&the determinant of $\\tilde g_{\\mu\\nu}$&(\\ref{13.20})\\\\\\hline $\\eta_{\\mu\\nu} $&the Minkowski (flat) metric tensor&(\\ref{exp})\\\\\\hline $\\Gamma^\\alpha_{\\mu\\nu} $& the Christoffel symbol  &(\\ref{covd})\\\\\\hline $R_{\\mu\\nu} $& the Ricci tensor &(\\ref{10.2})\\\\\\hline $R $& the Ricci scalar & (\\ref{10.1})\\\\\\hline $\\tilde R_{\\mu\\nu}$& the conformal Ricci tensor  &(\\ref{13.19})\\\\\\hline $T_{\\mu\\nu}$& the energy-momentum tensor of matter &(\\ref{10.2})\\\\\\hline $T $& the trace of the energy-momentum tensor &(\\ref{10.2})\\\\\\hline $\\phi$& the scalar field &(\\ref{10.1})\\\\\\hline $\\phi_0 $& the background value of the scalar field $\\phi$& (\\ref{aa})\\\\\\hline $\\zeta $& the dimensionless perturbation of the scalar field &(\\ref{aa})\\\\\\hline $\\theta(\\phi) $& the coupling function of the scalar field  &(\\ref{10.1})\\\\\\hline ${\\dAl}_g $& the Laplace-Beltrami operator &(\\ref{covd})\\\\\\hline $\\dAl $&the D'Alembert operator in the Minkowski  space-time  &(\\ref{11.29})\\\\\\hline $\\rho $& the density of matter in the co-moving frame &(\\ref{11.1})\\\\\\hline $\\rho^* $& the invariant (Fock) density of matter  &(\\ref{11.19})\\\\\\hline $\\Pi $& the internal energy of matter in the co-moving frame&(\\ref{11.1})\\\\\\hline $\\pi^{\\mu\\nu} $& the tensor of (anisotropic) stresses of matter   &(\\ref{11.1})\\\\\\hline $u^\\alpha $& the 4-velocity of matter &(\\ref{11.1})\\\\\\hline $v^i $&the 3-dimensional velocity of matter in the global frame &(\\ref{11.23})\\\\\\hline $\\omega$& the asymptotic value of the coupling function  $\\theta(\\phi) $  &(\\ref{10.5})\\\\\\hline $ \\omega' $& the asymptotic value of the derivative of the coupling function $\\theta(\\phi) $ &(\\ref{10.5})\\\\\\hline $c$& the ultimate speed of general and special theories of relativity&(\\ref{10.1})\\\\\\hline $\\epsilon $& a small dimensional parameter, $\\epsilon=1/c$ &(\\ref{exp})\\\\\\hline $h_{\\mu\\nu} $& the metric tensor perturbation, $g_{\\mu\\nu}-\\eta_{\\mu\\nu}$  &(\\ref{mtp})\\\\\\hline $\\nne{n}{\\mu}{\\nu} $& the metric tensor perturbation of order $\\epsilon^n$ in the post-Newtonian expansion of the metric tensor   &(\\ref{exp})\\\\\\hline $N $&a shorthand notation for $\\nne{2}{0}{0}$  &(\\ref{not})\\\\\\hline $L $&a shorthand notation for $\\nne{4}{0}{0}$  &(\\ref{not})\\\\\\hline $N_i $&a shorthand notation for $\\nne{1}{0}{i}$  &(\\ref{not})\\\\\\hline $L_i $&a shorthand notation for $\\nne{3}{0}{i}$  &(\\ref{not})\\\\\\hline $H_{ij} $&a shorthand notation for $\\nne{2}{i}{j}$  &(\\ref{not})\\\\\\hline $H $&a shorthand notation for $\\nne{2}{k}{k}$  &(\\ref{not})\\\\\\hline $\\tilde{N},\\; \\tilde{L}$& shorthand notations for perturbations of conformal metric $\\tilde g_{\\mu\\nu}$ &(\\ref{mp4})\\\\\\hline $\\gamma$&the `space-curvature' PPN parameter  &(\\ref{11.27})\\\\\\hline $\\beta $&the `non-linearity' PPN parameter  &(\\ref{11.28})\\\\\\hline $\\eta $& the Nordtvedt parameter, $\\eta =4\\beta-\\gamma-3$ &(\\ref{mp3})\\\\\\hline $G $&the observed value of the universal gravitational constant  &(\\ref{10.6})\\\\\\hline ${\\cal G} $&the bare value of the universal gravitational constant  &(\\ref{10.7}), (\\ref{10.6})\\\\\\hline $x^\\alpha=(x^0,x^i) $&the global coordinates with $x^0=ct$ and $x^i\\equiv{\\bm x}$  & \\\\\\hline $w^\\alpha=(w^0,w^i) $&the local coordinates with $w^0=cu$ and $w^i\\equiv{\\bm w}$  & \\\\\\hline $U $&the Newtonian gravitational potential in the global frame  &(\\ref{12.5})\\\\\\hline $U^{\\sss(A)} $&the Newtonian gravitational potential of body A in the global frame  &(\\ref{12.9a})\\\\\\hline $U_i $&a vector potential in the global frame  &(\\ref{12.8})\\\\\\hline $U_i^{\\sss(A)} $&a vector potential of body A in the global frame  &(\\ref{12.9a})\\\\\\hline $\\chi,\\;\\Phi_1,\\ldots,\\Phi_4 $&various special gravitational potentials in the global frame  &(\\ref{12.6}), (\\ref{12.9ex}) \\\\\\hline $V,\\; V^i $&potentials of the physical metric in the global frame &(\\ref{mp1}), \\ref{mp2})\\\\\\hline $\\sigma,\\;\\sigma^i $&the {\\it active} mass and current-mass densities in the global frame  &(\\ref{13.0}), (\\ref{13.1})\\\\ \\hline $I_{\\sss{<L>}} $&the {\\it active} Thorne-Blanchet-Damour mass multipole moments in the global frame  &(\\ref{13.9})\\\\ \\hline $S_{\\sss{<L>}} $&the {\\it active} spin multipole moments in the global frame  &(\\ref{13.10})\\\\\\hline $\\bar V $&potential of the scalar field in the global frame  &(\\ref{mp3})\\\\\\hline $\\bar\\sigma $&scalar mass density in the global frame  &(\\ref{13.13})\\\\\\hline $\\bar I_{\\sss{<L>}} $&scalar mass multipole moments in the global frame  &(\\ref{13.17})\\\\\\hline $\\tilde V $&gravitational potential of the conformal metric in the global frame  &(\\ref{mp4})\\\\\\hline $\\tilde\\sigma $&the {\\it conformal} mass density in the global frame  &(\\ref{13.21})\\\\\\hline $\\tilde I_{\\sss{<L>}} $&the {\\it conformal} mass multipole moments in the global frame  &(\\ref{13.30})\\\\\\hline $\\mathbb{M} $&conserved mass of an isolated system   &(\\ref{13.35})\\\\\\hline $\\mathbb{P}^i $&conserved linear momentum of an isolated system  &(\\ref{13.36})\\\\\\hline $\\mathbb{S}^i $&conserved angular momentum of an isolated system  &(\\ref{13.361})\\\\\\hline $\\mathbb{D}^i $&integral of the center of mass of an isolated system  &(\\ref{13.35a})\\\\\\hline $\\hat{A} $&symbols with the hat stand for quantities in the local frame  & \\\\\\hline $(B) $&sub-index referring to the body and standing for the internal solution in the local frame  &(\\ref{1.1}), (\\ref{1.2})\\\\\\hline $(E) $&sub-index referring to the external with respect to (B) bodies and standing for the external solution in the local frame&(\\ref{1.1}), (\\ref{1.2})\\\\\\hline $(C) $&sub-index standing for the coupling part of the solution in the local frame  &(\\ref{1.2})\\\\\\hline $P_{\\sss L} $&external STF multipole moments of the scalar field  &(\\ref{1.7a})\\\\\\hline $Q_{\\sss L} $&external STF gravitoelectric multipole moments of the metric tensor  &(\\ref{1.8a})\\\\\\hline $C_{\\sss L} $&external STF gravitomagnetic multipole moments of the metric tensor  &(\\ref{1.9a})\\\\\\hline $Z_{\\sss L},\\; S_{\\sss L} $&other sets of STF multipole moments entering the general solution for the space-time part of the external local metric&(\\ref{1.9a})\\\\\\hline $Y_{\\sss L}, B_{\\sss L}, D_{\\sss L}, E_{\\sss L}, F_{\\sss L}, G_{\\sss L}$& STF multipole moments entering the general solution for the space-space part of the external local metric  &(\\ref{1.10a})   \\\\\\hline ${\\mathcal V}_i,\\;\\Omega_i $&linear and angular velocities of kinematic motion of the local frame; we put them to zero throughout the rest of the paper &(\\ref{1.8a}), (\\ref{1.8aaa})\\\\\\hline $\\nu^i $&3-dimensional velocity of matter in the local frame  &(\\ref{1.12})\\\\\\hline ${\\cal I}_{\\sss{L}} $&active Thorne-Blanchet-Damour  STF mass multipole moments of the body in the local frame  &(\\ref{1.31})\\\\\\hline $\\sigma_{\\sss B}$&active mass density of body B in the local frame  &(\\ref{pz3})\\\\\\hline $\\bar{\\cal I}_{\\sss{L}} $&scalar STF mass multipole moments of the body in the local frame  &(\\ref{1.33})\\\\\\hline $\\bar\\sigma_{\\sss B} $&scalar mass density of body B in the local frame  &(\\ref{pz4})\\\\\\hline $\\tilde{\\cal I}_{\\sss{L}} $&conformal STF mass multipole moments of the body in the local frame  &(\\ref{1.34})\\\\\\hline $\\tilde\\sigma_{\\sss B} $&conformal mass density of body B in the local frame  &(\\ref{pz5})\\\\\\hline $\\sigma^i_{\\sss B} $&current mass density of body B in the local frame  &(\\ref{pz6})\\\\\\hline $S_{\\sss L} $&spin STF multipole moments of the body in the local frame  &(\\ref{1.32})\\\\\\hline $\\xi^0,\\;\\xi^i $&Relativistic corrections in the post-Newtonian transformation of time and space coordinates &(\\ref{2.2}), (\\ref{2.3})\\\\\\hline $x^i_{{\\sss B}},\\;\\nnv{i},\\;\\nnaa{i} $&position, velocity and acceleration of the body's center of mass with respect to the global frame &(\\ref{2.3}), (\\ref{2.8}), (\\ref{2.10})\\\\\\hline $\\nnr{i} $&$x^i-x^i_{{\\sss B}}(t),$ i.e. the spacial coordinates taken with respect to the center of mass of body B in the global frame  &(\\ref{2.3})\\\\\\hline ${\\cal A},\\;{\\cal B}_{\\sss{<L>}} $&functions appearing in the relativistic transformation of time  &(\\ref{2.8}), (\\ref{eq1}) \\\\\\hline ${\\cal D}_{\\sss{<L>}},\\;{\\cal F}_{\\sss{<L>}},\\;{\\cal E}_{\\sss{<L>}} $& functions appearing in the relativistic transformation of spacial coordinates  &(\\ref{eq2})\\\\\\hline $\\Lambda^{\\beta}_{\\;\\,\\alpha} $&matrix of transformation between local and global coordinate bases  &(\\ref{cb})\\\\\\hline $\\mathfrak{B},\\;\\mathfrak{D},\\;\\mathfrak{B}^i,\\;\\mathfrak{P}^i,\\;\\mathfrak{R}^{i}_{\\;j} $& PN corrections in the matrix of transformation $\\Lambda^{\\beta}_{\\;\\,\\alpha}$&(\\ref{ma00})--(\\ref{ma03})\\\\\\hline $\\bar U,\\;\\bar U^i,\\;$ etc.& external gravitational potentials  &(\\ref{3.1})--(\\ref{3.4})\\\\\\hline $\\bar U_{,L}\\nnxe,\\;\\bar U^i_{,L}\\nnxe $& $l$-th spatial derivative of an external potential taken at the center of mass of body B  &(\\ref{3.13})\\\\\\hline $\\mathcal{U}^{{\\sss(B)}} $& PN correction in the formula of matching of the local Newtonian potential  &(\\ref{3.5})\\\\\\hline ${F}^{ik} $&the matrix of relativistic precession of local coordinates with respect to global coordinates&(\\ref{5.8})\\\\ \\hline ${\\cal M}_*,\\;{\\cal J}^i_*,\\;{\\cal P}^i_* $& Newtonian-type mass, center of mass, and linear momentum of the body in the local frame  &(\\ref{a})--(\\ref{c})\\\\\\hline ${\\rm M} $& general relativistic PN mass of the body in the local frame &(\\ref{ij3})\\\\\\hline ${\\cal M} $& active mass of the body in the local frame &(\\ref{ij2})\\\\\\hline $\\tilde {\\cal M} $& conformal mass of the body in the local frame  &(\\ref{ij1})\\\\\\hline ${\\cal I}^{\\sss{(2)}} $&rotational moment of inertia of the body in the local frame  &(\\ref{fop})\\\\\\hline ${\\cal N}^{\\sss L} $& a set of STF multipole moments  &(\\ref{aer})\\\\\\hline ${\\cal P}^i $& PN linear momentum of the body in the local frame &(\\ref{6.2b})\\\\\\hline $\\Delta\\dot{\\cal P}^i $& scalar-tensor PN correction to $\\dot{\\cal P}^i$  &(\\ref{bvo})\\\\\\hline $\\tilde{\\cal M}_{ij} $&conformal anisotropic mass of the body in the local frame  &(\\ref{brt})\\\\\\hline ${\\mathbb {F}}^i_N,\\;\\Delta{\\mathbb {F}}^i_{N},\\;{\\mathbb {F}}^i_{pN},\\;\\Delta{\\mathbb {F}}^i_{pN} $& gravitational forces in the expression for $Q_i$  &(\\ref{6.9})--(\\ref{fopm})\\\\\\hline ${\\cal S}^i$&the bare post-Newtonian definition of the angular momentum (spin) of a body&(\\ref{spin-3}) \\\\\\hline ${\\cal T}^i$&the post-Newtonian torque in equations of rotational motion&(\\ref{spin-5}) \\\\\\hline $\\Delta{\\cal T}^i$&the post-Newtonian correction to the torque ${\\cal T}^i$ &(\\ref{spin-6}) \\\\\\hline $\\Delta{\\cal S}^i$&the post-Newtonian correction to the bare spin ${\\cal S}^i$ &(\\ref{spin-7}) \\\\\\hline ${\\cal R}^i$&velocity-dependent multipole moments &(\\ref{spin-8}) \\\\\\hline ${\\cal S}^i_+$&the (measured) post-Newtonian spin of the body&(\\ref{spin-9})\\\\\\hline $r $& radial space coordinate in the body's local frame, $r=|{\\bm w}|$ &(\\ref{as7})\\\\\\hline ${\\Omega}^j_{\\sss B} $& angular velocity of rigid rotation of the body B referred to its local frame &(\\ref{qma})\\\\\\hline $I^{(2l)}_{\\sss B} $& $l$-th rotational moment of inertia of the body B  &(\\ref{9.6.3})\\\\\\hline $\\mathbb{I}_{\\sss B}^L $& multipole moments of the multipolar expansion of the Newtonian potential in the global coordinates    &(\\ref{aq9})\\\\\\hline $R_{\\sss B} $&$|{\\bm R}_{\\sss B}|$, where ${\\bm R}_{\\sss B}={\\bm x}-{\\bm x}_{\\sss B}$ &(\\ref{re4})\\\\\\hline $R_{\\sss{BC}}^i $&$x^i_{\\sss C} - x^i_{\\sss B}$ &(\\ref{pn7})\\\\\\hline $F^i_{\\sss N},\\;F^i_{\\sss{EIH}},\\;F^i_{\\sss{\\cal S}},\\;F^i_{\\sss{\\cal I}GR},\\; \\delta F^i_{\\sss{\\cal I}GR} $& forces from the equation of motion of spherically-symmetric bodies  &(\\ref{gko1})--(\\ref{gko5})\\\\ \\hline $\\mathfrak{M}_{\\sss B} $&Nordtvedt's gravitational mass of the body B  &(\\ref{dm})\\\\\\hline \\end{longtable} ",
        "conclusions": ""
    },
    "0403/astro-ph0403510_arXiv.txt": {
        "abstract": "Optical and infrared monitoring of the afterglow site of gamma-ray burst (GRB) 031203 has revealed a brightening source embedded in the host galaxy, which we attribute to the presence of a supernova (SN) related to the GRB (``SN 031203''). We present details of the discovery and evolution of SN 031203 from 0.2 to 92 days after the GRB, derived from SMARTS consortium photometry in I and J bands. A template type Ic lightcurve, constructed from SN 1998bw photometry, is consistent with the peak brightness of SN 031203 although the lightcurves are not identical. Differential astrometry reveals that the SN, and hence the GRB, occurred less than $300 h_{71}^{-1}$ pc (3 $\\sigma$) from the apparent galaxy center. The peak of the supernova is brighter than the optical afterglow suggesting that this source is intermediate between a strong GRB and a supernova. ",
        "introduction": "Since the discovery of GRB afterglows, the evidence for a physical connection between gamma-ray bursts (GRBs) and core-collapse supernovae (SNe) has mounted \\citep[see reviews by][]{vanp99,mez01}. Particularly compelling were observations of lightcurves and broadband photometry of SN-like features embedded in GRB afterglow light \\citep[see][]{blo03}. Recently, spectroscopic evidence \\citep{Stanek03,Hjorth03,kdw+03} confirmed that GRBs are produced in the death of massive stars \\citep{woo93}. To date, SN signatures have been reliably found in only a few GRBs (see \\citealt{blo03}) necessitating the search for and the study of new GRB-related SNe.  GRB 031203 triggered the IBIS instrument onboard the {\\it Integral} satellite on 3 December 2003 at 22:01:28 UT \\citep{Gotz03}, leading to quick discoveries of X-ray \\citep{Camp03} and radio afterglows \\citep{Frail03, SKF03}. Spectroscopy of the host galaxy coincident with the radio transient yielded a redshift of $z=0.1055$ \\citep{Proch04}, likely the redshift of the burst itself. The low redshift (second only to the unusual GRB 980425) of GRB 031203 presents a rare opportunity to create a well-sampled SN lightcurve using modest aperture telescopes. We began our observations of the field 5 hours after trigger and continued monitoring periodically for several months. We reported our discovery of an increase in brightness of the aperture magnitude of the host, and suggested the emergence of a supernova was responsible \\citep{Bailyn03}. Hereafter, since the explosion date of the SN is likely that of the GRB, we designate the SN associated with GRB 031203 as ``SN 031203''. Monitoring of the SN by other groups has now confirmed the presence of SN 031203 both photometrically \\citep{Bersier04} and spectroscopically \\citep{Tag04}.  In this paper we present optical and infrared data obtained with the SMARTS 1.3m telescope and ANDICAM instrument between 0.2 and 92 days after the detection of GRB 031203. Observations and data reduction are reported in section 2. Section 3 describes the aperture photometry and image subtraction carried out on this data. The resultant evidence of a SN associated with GRB 031203 is presented. A comparison between this SN and SN 1998bw is made in section 4. ",
        "conclusions": ""
    },
    "0403/astro-ph0403456_arXiv.txt": {
        "abstract": "We describe the design and construction of a formatted fiber field-unit, SparsePak, and characterize its optical and astrometric performance. This array is optimized for spectroscopy of low-surface brightness, extended sources in the visible and near-infrared. SparsePak contains 82, 4.7$^{\\prime\\prime}$ fibers subtending an area of 72$^{\\prime\\prime}\\times$71$^{\\prime\\prime}$ in the telescope focal plane, and feeds the WIYN Bench spectrograph. Together, these instruments are capable of achieving spectral resolutions of $\\lambda/\\Delta\\lambda\\sim 20,000$ and an area--solid-angle product of $\\sim 140$ arcsec$^2$ m$^2$ per fiber. Laboratory measurements of SparsePak lead to several important conclusions on the design of fiber termination and cable curvature to minimize focal ratio degradation. SparsePak itself has throughput $>80$\\% redwards of 5200\\AA\\, and 90-92\\% in the red. Fed at f/6.3, the cable delivers an output 90\\% encircled energy at nearly f/5.2. This has implications for performance gains if the WIYN Bench Spectrograph had a faster collimator. Our approach to integral-field spectroscopy yields an instrument which is simple and inexpensive to build, yet yields the highest area--solid-angle product per spectrum of any system in existence. An Appendix details the fabrication process in sufficient detail for others to repeat. SparsePak was funded by the National Science Foundation and the University of Wisconsin-Madison Graduate School, and is now publicly available on the WIYN Telescope through the National Optical Astronomical Observatories. ",
        "introduction": "Observational astronomy consists of obtaining subsets of a fundamental data hyper-cube of the apparent distribution of photons in angle$^2$ on the sky $\\times$ wavelength $\\times$ time $\\times$ polarization. Information-gathering systems (``instruments'') are designed to make science-driven trades on the range and sampling of each of these dimensions. Here we describe an instrument optimized for the study of the stellar and ionized gas kinematics in disks of nearby and distant galaxies. Such studies require bi-dimensional spectroscopy at medium spectral resolution ($5000<R<20,000$, where $R\\equiv\\lambda/\\Delta\\lambda$) of extended sources over a relatively narrow range of wavelength, e.g., $\\sim$600 spectral channels, with no consideration of time-sampling or polarization.  What is paramount for our application is the ability to gather sufficient signal at low light levels and medium spectral resolution. Etendue at constant spectral resolution and sampling is the relevant figure of merit (etendue is the product of area, solid angle, and system throughput).  To characterize the light-gathering aperture alone in what follows we refer to ``grasp'' (the area $\\times$ solid angle product).  Field-of-view and spatial resolution are also of merit, but of secondary importance. For our particular application, we require the number of spatial resolution elements be sufficient to resolve galaxy disks out to 3-4 scale lengths at several points per scale-length -- about 15 points across a diameter. The matching of overall scale is dictated, then, by the telescope and fiber size needed to reach the required spectral resolution and etendue with an affordable spectrograph.  That is, the galaxies are to be chosen to fit the instrument.   Bacon \\etal\\ (1995) and Ren \\& Allington-Smith (2002) discuss the trades between spatial and spectral dynamic range and sampling for a variety of spectroscopic instruments -- all fundamentally limited by the two-dimensional sampling geometry of their detector focal planes. An evaluation of this discussion yields that integral-field spectroscopy (IFS) and Fabry-Perot (FP) imaging are the preferred methods for bi-dimensional spectroscopy at low light levels. Each has its merits. Relative to imaging FP systems, IFS systems trade spatial resolution and coverage for spectral resolution and spectral coverage. For emission-line studies in need of high spatial resolution FP are optimal. For lower spatial resolution, particularly at high spectral resolution where the FP ``bull's eye'' is small, IFS is particularly competitive. (The ``bull's eye'' is the angular diameter in which the central wavelength shifts by less than the spectral resolution.) For absorption line studies where many spectral channels are needed, IFS is superior. This condition becomes even more pronounced if spectral resolution and etendue are valued more highly than spatial resolution, as we do here.    Within the realm of ``integral-field'' spectroscopy there is still a wide range of spatial and spectral sampling (see Ren \\& Allington-Smith 2002). Our interest is in fiber-fed or image-slicing systems which reformat the telescope focal plane (i.e., into a slit), optimizing for spectral sampling while still retaining bi-dimensional spatial coverage (cf. SAURON, Bacon \\etal\\ 2001a, whic is optimized for spatial sampling at the sacrifice of spectral sampling). While image slicers minimize entropy increases, fibers relax the restriction on the spatial sampling. In some situation, the entropy increase (i.e., information loss) from fibers via focal-ratio degradation (FRD; see Angel \\etal\\ 1977; Barden, Ramsey \\& Truax, 1981; Clayton 1989; or Carrasco \\& Parry 1994) may be more than compensated by the increased flexibility in how the fibers sample the telescope focal plane and are mapped into the dispersive optical system.  For example, fibers offer the advantage of formatted sampling, whereby fibers intentionally do not sample an integral area, but instead some structured, coherent, yet dispersed pattern.  While this ``patterning'' has not been used in fixed-bundle arrays, there are clear advantages to this approach. For an extended object, all spatial elements do not carry equal scientific weight, and there is a trade between coverage and sampling. Fibers allow these trades to be fine-tuned. In the specific example of galaxy kinematic studies, one may choose to allocate fibers to sample both major and minor axes, or perhaps uniformly cover a larger physical area with sparse sampling such that lower surface-brightness outer regions are still sampled more frequently than the inner, higher surface-brightness cores. The latter is the approach we have adopted here.  These considerations led us to our design of a {\\it formatted} fiber array for the WIYN Bench Spectrograph -- the first of its kind. It is a simple, inexpensive instrument with dramatic gains in information gathering power for a broad range of scientific programs which require bi-dimensional, medium-resolution spectroscopy at low surface-brightness levels. The SparsePak design is in an orthogonal direction to most 8m-class IFU instruments which strive to maximize spatial resolution at high cost and complexity and at the loss of medium spectral resolution for background limited observations.  Our strategy was to take advantage of an existing spectrograph capable of echelle resolutions with a very long slit (equivalent to about 12 arcmin) to optimize the study of galaxy kinematics at low surface brightness. The result of this mating of fiber cable to spectrograph yields a bi-dimensional spectroscopic system capable of achieving spectral resolutions and grasp comparable to the best systems on {\\it any} telescope. The specific grasp of SparsePak -- the grasp per spatial resolution element -- is the highest of any spectroscopic instrument (see Figure 1). It is this singular attribute that allows SparsePak to be used effectively at low surface-brightness and medium spectral resolutions. This fiber array was completed on March 10, 2001, installed in May 2001, and successfully commissioned over the following month.  In a series of two papers we describe the design, construction, and performance of this array and its associated spectrograph.  We start in \\S 2 with a basic description of SparsePak and the spectrograph it feeds, as well as a comparison of SparsePak to existing fiber-arrays on WIYN and other telescopes. In \\S 3 of this paper we present the design goals and constraints which led to the SparsePak formatted-array design. A synopsis of SparsePak's technical attributes are presented in \\S 4. The construction of single-fiber reference cables are also documented. This cable provides a benchmark for evaluating the laboratory-measured performance of SparsePak.  Sections 5 and 6 contain the astrometric and optical properties of the array, as measured in our lab. In \\S 7 is a summary of our results.  An appendix contains a technical description of the fabrication process in sufficient detail for others to repeat the process.  In Paper II (Bershady \\etal\\ 2003) we establish the on-sky performance of the array and spectrograph in terms of throughput, spectral resolution and scattered-light; we demonstrate the ability to perform precision sky-subtraction and spectrophotometry; and we present examples of commissioning science data which highlight the capabilities for which SparsePak was designed. ",
        "conclusions": "In this paper we have presented the design, construction, and laboratory measurements of the SparsePak FFU, a formatted fiber-optic array designed to mate the Nasmyth imaging port to the Bench Spectrograph on the WIYN 3.5-m telescope. A physical description of the SparsePak array is found in \\S 2 and 4, with a complete description of the design and assembly contained in the Appendix. The latter is included specifically to allow others to repeat this relatively simple instrument-building process.  The primary scientific motivation for this array is to measure the kinematics of stars and gas in nearby galaxy disks. Examples of science capabilities are found in Paper II, and in several recent studies using SparsePak to measure stellar velocity dispersions to estimate dynamical mass and the radial dependence of the mass-to-light ratio in spiral disks radius (Bershady \\etal\\ 2002, Verheijen \\etal\\ 2003, 2004), the Tully-Fisher relation in barred spirals (Courteau \\etal\\ 2003), and rotation curves and ionized gas kinematics in low-surface-brightness galaxies (Swaters \\etal\\ 2003).  Here we have focused on the confluence of science objectives with technical, observational, and physical constraints that shaped the specific SparsePak design. The salient design drivers are the balancing of trade-offs in spatial and spectral coverage and resolution, plus an optimization for superior sky-subtraction. For this latter purpose we have developed an analytic expression for the optimum number of sky fibers in the background- and detector-limited regimes. Placement of sky and source fibers along the spectrograph slit are also important considerations in our design -- both for sky subtraction and to offset the effects of vignetting within the spectrograph. In Paper II we explore the sky-subtraction performance of our design. The results presented here and in Paper II are applicable to other fiber-fed spectrographs.  The SparsePak performance can be summarized in terms of its instrumental grasp of 137 m$^2$ arcsec$^2$ per fiber, and 11,200 m$^2$ arcsec$^2$ overall. In more details, the cable has a throughput of 89-92\\% redwards of 500 nm (92\\% is the expected best-case for two fused-silica--air interfaces). The cable throughput drops rapidly below 500 nm. The overall system throughput (telescope plus cable plus spectrograph) are measured and described in Paper II. The fibers produce FRD which takes an input f/6.3 beam and degrades to f/5.7 to f/4.1 at EE95. There is no wavelength dependence to the observed FRD between 500 and 800 nm. The non-telecentricity error budget is below 0.1 deg (differential) at the telescope focal-plane, and does not contribute significantly to the observed FRD. The range of FRD depends, to first order, on a fiber's slit position (due to mechanical curvature of spectrograph feed) and, to second order, on its position in the head (due to edge stresses).  The implications are that future fiber-arrays manufactured with the process described here should have thicker buffers of inactive fibers (at least 2 fibers thick) on the outside walls of the array, and straighter feeds.  There are several conclusion which may be drawn from our study. First, we have shown that economical, high-performance optical cables are possible to build which are relatively long in length.  The losses within the fibers themselves are relatively small; in a fiber cable of 25m length the end-losses dominate. This means that significantly longer cables can be constructed with high-throughput performance in the red. This may be useful for connecting spectrographs to multiple telescopes at a given site.  Second, our laboratory measurements indicate the promise of significantly lower levels of FRD with modest modifications to our fiber-head design construction and closer attention to the bending geometry near the spectrograph termination -- even for very large (500 $\\mu$m) fibers. This optimism should be tempered by the fact that all other existing cables on the Bench Spectrograph exhibit comparable FRD to SparsePak. However, it is unclear whether the other Bench fiber feeds have been optimized to minimize fiber stress.   Finally, within the context of the current fiber cables feeding the WIYN Bench Spectrograph, it is clear that much of the light is being lost due to the injection of a beam faster than the spectrograph was designed to handle. In Paper II we document how much of the light is lost from the SparsePak cable due to geometric vignetting factors within the Bench Spectrograph. Roughly speaking, a modest decrease in the collimator focal-length would gain back factors of two in throughput while at worst decreasing the spectral resolution by only 30\\%. For many programs this trade of throughput for resolution is a winning proposition. Moreover, the resolution losses are not even this severe since the current system is under-sampled; higher resolution can be regained for stellar surveys by using smaller fibers with little light-loss in the telescope focal plane -- thanks to WIYN's excellent image quality. Given the cost of running the WIYN telescope and the number of nights the Bench Spectrograph is in use, it would seem unimaginable were such an upgrade not implemented in short order."
    },
    "0403/astro-ph0403383_arXiv.txt": {
        "abstract": "s{We present new results on the clustering and three-dimensional distribution of radio galaxies from the Texas-Oxford NVSS Structure (TONS) survey. The TONS survey was constructed to look at the distribution of radio galaxies in a region of moderate (0 $\\ltsimeq$ $z$ $\\ltsimeq$ 0.5) redshifts by matching NVSS sources with objects in APM catalogues to obtain a sample of optically bright (R $\\le$ 19.5), radio faint (1.4-GHz flux density S$_{1.4} \\ge$ 3 mJy) radio galaxies over large areas on the sky. We find that redshift spikes, which represent large concentrations of radio galaxies which trace ($\\approx$ 100 $\\rm {Mpc}^3$) super-structures are a common phenomena in these surveys. Under the assumption of quasi-linear structure formation theory and a canonical radio galaxy bias, the structures represent $\\approx$ 4-5$\\sigma$ peaks in the primordial density field and their expected number is low. The most plausible explanation for these low probabilities is an increase in the radio galaxy bias with redshift. To investigate potential mechanisms which have triggered the radio activity in these galaxies - and hence may account for an increase in the bias of this population, we performed imaging studies of the cluster environment of the radio galaxies in super-structure regions. Preliminary results show that these radio galaxies may reside preferentially at the edges of rich clusters. If radio galaxies are preferentially triggered as they fall towards rich clusters then they would effectively adopt the cluster bias. } ",
        "introduction": "Radio galaxies are ideal probes of large-scale structure as they are biased tracers of the underlying mass and can be easily detected out to high redshifts. By using biased galaxies populations, one can efficiently trace huge super-structures (i.e clusters of clusters of galaxies) which are still in the linear regime and can therefore be directly traced back to rare fluctuations in the initial density field at recombination. However, in order to be useful probes, it is vital to understand how different populations of radio galaxies trace the underlying dark matter (i.e. their bias) and how this has evolved with time. This is directly related to how the radio activity is triggered in different populations and in different environments. ",
        "conclusions": "The TONS08 radio galaxies within super-structure regions are generally in moderately rich (Abell class 0) environments. However, 63 per cent of the radio galaxies are within a projected distance of 3 Mpc from the centre of a cluster candidate. The fact that we see so many radio galaxies near rich clusters suggests that the radio galaxies are associated with rich clusters but often only on the edges of high overdensity regions. This explains why we find that the radio galaxies are only in moderately rich environments: many of the radio galaxies are further than 0.5 Mpc from the cluster centre. One possible scenario is that of radio galaxies at the centres of poor groups of galaxies being preferentially triggered as the group falls down large-scale structure filaments towards rich clusters. The radio galaxies would then effectively adopt the cluster bias, and the number of redshift spikes we see in the data would become consistent with the number that we expect.  \\hspace{1mm}  \\noindent This material is based in part upon work supported by the Texas Advanced Research Program under Grant No. 009658-0710-1999."
    },
    "0403/astro-ph0403106_arXiv.txt": {
        "abstract": "The results of analysis of HST $I-$band imaging of a sample of 41 radio galaxies spanning three orders of magnitude in radio luminosity at redshift $z\\simeq0.5$ are presented. These results represent the first stage in a coordinated programme to study the connection between radio luminosity and host-galaxy properties, black-hole masses and cluster environments in radio galaxies spanning a wide range in radio luminosity over a restricted range in redshift. The full sample is comprised of objects drawn from four complete, low-frequency selected radio samples with progressively fainter flux-density limits (3CRR, 6CE, 7CRS and the new TexOx-1000 sample). Modelling of the HST imaging data shows that the host galaxies have surface-brightness distributions consistent with those expected for classic ellipticals (S\\'{e}rsic parameter, $\\beta\\simeq0.25$), with $\\beta$ in the range $0.17<\\beta<0.30$, and a mean of $<\\beta>=0.23\\pm0.01$. The luminosities of the host galaxies are found to be comparable with those of galaxies drawn from the bright end of the local cluster galaxy luminosity function, spanning the range $0.7L^{\\star}<L<10 L^{\\star}$, with a mean of $3.2\\pm0.3 L^{\\star}$, after correcting for the effects of passive evolution. In addition, the radio galaxies are shown to follow a Kormendy ($\\mu_{e}-r_{e}$) relation indistinguishable from that of both powerful low-redshift radio galaxies and local Abell brightest cluster galaxies. Combining our new results with those in the literature it is found that the scalelengths and Kormendy relations of 3C-class radio galaxies do not vary significantly over the redshift range $0.0<z<0.8$, providing no evidence for dynamical evolution of this class of host galaxy within this redshift interval. Converting the host-galaxy luminosities into black-hole mass estimates, using the local $M_{bh}-M_{bulge}$ correlation, predicts that the radio galaxies harbour central black holes with masses in the range $10^{8.1}\\Msun <M_{bh}< 10^{9.5}\\Msun$, with a geometric mean of $<M_{bh}>=10^{8.87\\pm 0.04}\\Msun$.  Finally, a significant ($\\simeq 3\\sigma$) correlation is found between black-hole mass and 151-MHz radio luminosity for those objects in the sample with either high-excitation nuclear spectra (HEG) or classical double (CD) radio structures. ",
        "introduction": "The underlying mechanisms which govern the large range in low-frequency radio luminosity observed in radio-loud active galactic nuclei (AGN) are currently unknown. Over the previous two decades a wealth of observational effort has been invested in studying the redshift r\\'{e}gime $0<z<1$, investigating whether the dominant factors are linked to the Mpc-scale cluster environments (eg. Prestage \\& Peacock 1988; Hill \\& Lilly 1991; Ellingson, Yee \\& Green 1991; Wold et al. 2000; Best 2000) or the kpc-scale properties of the host galaxies (e.g. Smith \\& Heckman 1989; Best, Longair \\& R\\\"{o}ttgering 1998; McLure \\& Dunlop 2000; Dunlop et al. 2003; Zirm, Dickinson \\& Dey 2003).  One of the incontrovertible observational facts which has emerged from these studies, at least at $z<1$, is that the host galaxies of all powerful radio-loud AGN are massive $L>L^{\\star}$ ellipticals (eg. Taylor et al. 1996; McLure et al. 1999; Dunlop et al. 2003). The apparent uniformity of radio-loud AGN host galaxies has taken on added importance over the last few years, following the discovery in nearby (distance $\\ltsim$ 150 Mpc) inactive galaxies that a reasonably accurate estimate ($\\Delta M_{bh}\\simeq 0.3$ dex) of the central black-hole mass can be obtained via its correlation with the mass of the host spheroidal component (Kormendy \\& Richstone 1995; Magorrian et al. 1998; Gebhardt et al. 2000; Ferrarese \\& Merritt 2000; McLure \\& Dunlop 2002; Marconi \\& Hunt 2003; Tremaine et al. 2002). Moreover, recent progress has also indicated that a similarly accurate black-hole mass estimate ($\\Delta M_{bh}\\simeq 0.4$ dex) can be obtained for broad-line AGN using emission-line widths to derive the virial mass estimate (eg. Kaspi et al. 2000; McLure \\& Dunlop 2002; McLure \\& Jarvis 2002; Vestergaard 2002). Consequently, a large body of work has appeared in the recent literature investigating the possible link between radio luminosity and black-hole mass in radio-loud AGN (eg.  Laor 2000; Lacy et al. 2001; McLure \\& Dunlop 2001a; McLure \\& Dunlop 2002; Bettoni et al. 2003; Dunlop et al. 2003).  Unfortunately, observational studies have traditionally been subject to a degeneracy between radio luminosity and redshift produced as a by-product of flux-limited radio samples. In order to study the properties of radio-loud AGN separated by a large dynamic range in radio luminosity, it has previously been necessary to select samples consisting of objects covering a wide range of redshifts. This has led to difficulties in interpreting the data due to the complication of potentially significant evolutionary effects. However, by selecting our sample of objects from four complete, low-frequency selected radio samples with successively fainter flux-density limits, it has been possible to construct a sample of radio galaxies which spans three decades in radio luminosity at a virtually constant cosmic epoch ($0.4<z<0.6$). In this respect the motivation for this study is identical to that pursued by Hill \\& Lilly (1991), who investigated the cluster environments of a sample of 45 radio galaxies spanning a similar range in radio luminosities at $z\\simeq 0.5$. By successfully disentangling the effects of luminosity and redshift for the first time, Hill \\& Lilly (1991) were able to demonstrate that $z\\simeq 0.5$ marks an apparent epoch-dependent change in cluster environment, with radio galaxies tending to inhabit significantly richer cluster environments than their low-redshift counterparts. Furthermore, the wide range of radio luminosities in the Hill \\& Lilly sample enabled them to determine that extended radio luminosity did not appear to be a strong function of cluster richness.  Developments over the intervening decade have provided our new study with two crucial advantages over the previous work of Hill \\& Lilly (1991). Firstly, the availability of the 7CRS (Willott et al. 2003; Lacy et al. 1999) and TexOx-1000 (Hill \\&  Rawlings 2003) radio samples has now allowed us to compile a sample of 41 radio galaxies at $z\\simeq 0.5$, which is complete over the full three-decades range in radio power. Secondly, with the high-resolution imaging provided by the Hubble Space Telescope (HST), and the improvement in follow-up observations available with ground-based telescopes, we are now in a position to obtain a wide range of high quality data-sets for our new sample which were not obtainable only ten years ago. Consequently, our new sample of $z\\simeq 0.5$ radio galaxies is the basis for an ambitious project, which will for the first time systematically investigate the connection between the radio luminosity of radio galaxies, the properties of their host galaxies, their location on the fundamental plane, the richness and mass of their cluster environments and the masses of their central black holes.  The results presented in this paper represent the first stage of this wide-ranging project, and concentrate solely on the analysis of our deep $I-$band HST WFPC2 imaging data. The results of our extensive follow-up observations of the sample are deferred to a series of future papers.  The structure of the paper is as follows, in Section 2 the details of the radio-galaxy sample are described, while in Section 3 the HST observations and data reduction are discussed. In Section 4 the modelling of the HST imaging is described, with the main results presented and compared to literature results on low-redshift radio galaxies in Section 5. In Section 6 the host-galaxy properties of our 3CRR sub-sample are compared to literature results to explore the evidence for dynamical evolution of the most powerful radio galaxies within the redshift range $0.0<z<0.8$. In Section 7 the properties of the radio-galaxy hosts are compared to those of local brightest cluster galaxies. In Section 8 the masses of the radio galaxies' central black holes are estimated, and in Section 9 the relationship between black-hole mass and radio luminosity is investigated. The main conclusions are summarized in Section 10. Unless otherwise stated, throughout this paper the following cosmology is assumed: $H_{0}=70$ km\\,\\,s$^{-1}$Mpc$^{-1}$, $\\Omega_{m}=0.3$, $\\Omega_{\\Lambda}=0.7$.    \\begin{table*} \\begin{center} \\caption{The full ZP5 sample. Column one lists the radio-galaxy names, while columns two and three list the J2000 source coordinates. Column 4 lists the radio-galaxy redshifts and column five lists the logarithm of the 151-MHz luminosities in units of$~$WHz$^{-1}$sr$^{-1}$. Column 6 classifies the radio structures as either Classical Double (CD), Fat Double (FD), Jet (J) or Compact (CM), the same classifications as used by Owen \\& Laing (1989). The radio structure classifications are based on consistent 1.4 GHz VLA A-array data (1.5 arcsec resolution) for the 6CE, 7CRS and TOOT sub-samples. For the 3CRR sub-sample, for which structural classifications are well established, the classifications were based on the high-quality maps available in the literature (at a variety of frequencies/resolutions). Column 7 describes the optical nuclear spectra in terms of the presence of high-excitation narrow emission lines (HEG) or not (LEG), using the classification scheme of Jackson \\& Rawlings (1997). Those objects for which the available optical spectra are of insufficient quality to allow a reliable HEG/LEG determination are listed with a `?'. To avoid confusion we include here in parenthesis the old names of the 7CRS objects which have been adopted in previous papers: 7C0213+3418 (5C6.63), 7C0219+3419 (5C6.201), 7C0219+3423 (5C6.214), 7C0220+2952 (5C6.233), 7C0223+3415 (5C6.279), 7C0810+2650 (5C7.7).} \\begin{tabular}{lcccccl} \\hline Source       &   RA & DEC &z   &  $L_{151}$ &Radio Structure&Class\\\\ \\hline 3C 16        & 00 37 45.39 & +13 20 09.6 &0.405&  26.82 &CD&HEG\\\\ 3C 19        & 00 40 55.01 & +33 10 07.3 & 0.482&  26.96&CD&LEG\\\\ 3C 46        & 01 35 28.47 & +37 54 05.7 & 0.437&  26.84&CD&HEG\\\\ 3C 172       & 07 02 08.32 & +25 13 53.9 & 0.519&  27.17&CD&HEG\\\\ 3C 200       & 08 27 25.38 & +29 18 45.5 & 0.458&  26.92&CD&LEG\\\\ 3C 225B      & 09 42 15.41 & +13 45 51.0 & 0.582&  27.50&CD&HEG?\\\\ 3C 228       & 09 50 10.79 & +14 20 00.9 & 0.552&  27.37&CD&HEG\\\\ 4C 74.16     & 10 14 14.84 & +74 37 37.4 & 0.568&  27.16&CD&HEG\\\\ 3C 244.1     & 10 33 33.97 & +58 14 35.8 & 0.428&  27.10&CD&HEG\\\\ 3C 274.1     & 12 35 26.64 & +21 20 34.7 & 0.422&  27.02&CD&HEG\\\\ 3C 295       & 14 11 20.65 & +52 12 09.0 & 0.464&  27.68&CD&HEG\\\\ 3C 330       & 16 09 35.01 & +65 56 37.7 & 0.550&  27.43&CD&HEG\\\\ 3C 341       & 16 28 04.04 & +27 41 39.3 & 0.448&  26.88&CD&HEG\\\\ 3C 427.1     & 21 04 07.07 & +76 33 10.8 & 0.572&  27.53&CD&LEG\\\\ 3C 457       & 23 12 07.57 & +18 45 41.4 & 0.428&  27.00&CD&HEG\\\\ \\hline 6C 0825+3407 &  08 25 14.59 & +34 07 16.8 &0.406&  26.09&FD&LEG\\\\ 6C 0850+3747 &  08 50 24.77 & +37 47 09.1 &0.407&  26.15&FD&HEG\\\\ 6C 0857+3945 &  08 57 43.56 & +39 45 29.0 &0.528&  26.34&CD&HEG\\\\ 6C 1111+3940 &  11 11 19.39 & +39 40 14.5 &0.590&  26.33&CD&LEG\\\\ 6C 1132+3439 &  11 32 45.74 & +34 39 36.2 &0.512&  26.33&CD&HEG\\\\ 6C 1200+3416 &  12 00 53.34 & +34 16 47.3 &0.530&  26.17&CD&LEG\\\\ 6C 1303+3756 &  13 03 44.26 & +37 56 15.2 &0.470&  26.29&CD&HEG\\\\ \\hline 7C 0213+3418 &  02 13 28.39 & +34 18 30.6 &0.465&  25.66&CD&LEG\\\\ 7C 0219+3419 &  02 19 15.89 & +34 19 43.2 &0.595&  26.13&FD&HEG?\\\\ 7C 0219+3423 &  02 19 37.83 & +34 23 11.2 &0.595&  25.98&J&HEG?\\\\ 7C 0220+2952 &  02 20 34.26 & +29 52 19.5 &0.560&  26.07&FD&LEG?\\\\ 7C 0223+3415 &  02 23 47.24 & +34 15 11.9 &0.473&  25.55&CD&HEG\\\\ 7C 0810+2650 &  08 10 26.10 & +26 50 49.1 &0.435&  25.58&CM&HEG\\\\ 7C 1731+6638 &  17 31 43.84 & +66 38 56.7 &0.562&  25.62&CD&HEG\\\\ 7C 1806+6831 &  18 06 50.16 & +68 31 41.9 &0.580&  26.36&CD&HEG\\\\ \\hline TOOT 0009+3523 &  00 09 46.90 & +35 23 45.1 &0.439&  24.79&FD&LEG\\\\ TOOT 0013+3459 &  00 13 13.29 & +34 59 40.6 &0.577&  25.75&FD&LEG\\\\ TOOT 0018+3510 &  00 18 53.93 & +35 10 12.1 &0.416&  25.16&J&LEG\\\\ TOOT 1255+3556 &  12 55 55.83 & +35 56 35.8 &0.471&  25.01&J&LEG\\\\ TOOT 1301+3658 &  13 01 25.03 & +36 58 09.4 &0.424&  24.76&J&LEG\\\\ TOOT 1303+3334 &  13 03 10.29 & +33 34 07.0 &0.565&  25.66&J&HEG\\\\ TOOT 1307+3639 &  13 07 27.07 & +36 39 16.4 &0.583&  25.30&J&LEG\\\\ TOOT 1309+3359 &  13 09 53.95 & +33 59 28.2 &0.464&  24.91&CM&HEG\\\\ TOOT 1626+4523 &  16 26 48.50 & +45 23 42.6 &0.458&  25.03&FD&LEG\\\\ TOOT 1630+4534 &  16 30 32.80 & +45 34 26.0 &0.493&  25.17&J&LEG\\\\ TOOT 1648+5040 &  16 48 26.19 & +50 40 58.0 &0.420&  25.12&CM&LEG\\\\ \\hline \\end{tabular} \\label{tab1} \\end{center} \\end{table*}  \\begin{figure*} \\centerline{\\epsfig{file=md1346rv_fig1.ps,width=12cm,angle=0}} \\caption{151-MHz radio luminosity versus redshift for the narrow-line radio galaxy component of the 3CRR, 6CE and 7CRS samples. The members of the ZP5 sample are highlighted as enlarged symbols in the redshift interval $0.4<z<0.6$, including those drawn from the new TexOx-1000 (TOOT) survey.} \\label{fig1} \\end{figure*}  \\begin{table*} \\begin{center} \\caption{List of sources which met the redshift and radio-luminosity criteria for inclusion in the ZP5 sample, but were excluded due to the presence of broad emission lines. In column 2 the sources are classified as either quasar (Q) or weak quasar (WQ) according to the classification scheme of Jackson \\& Rawlings (1997). We note here that the excluded quasars are confined to the two most radio luminous sub-samples, in agreement with the findings of Willott et al. (2000) that the quasar fraction in complete, low-frequency selected radio samples is an increasing function of radio luminosity. The implications of the possible quasar cores in the ZP5 sample (see Table 3) will be fully discussed in a forthcoming paper (Mitchell et al., in prep).} \\begin{tabular}{lcl} \\hline Source         &Class&Comment\\\\ \\hline 3C 47          &Q &\\\\ 3C 147         &Q &\\\\ 3C 215         &Q &\\\\ 3C 275.1       &Q &\\\\ 3C 334         &Q &\\\\ 3C 455         &WQ&\\\\ 6C 1220+3723   &Q & B1950.0 name - as in Rawlings et al. (2001)\\\\ 5C 7.118       &WQ& B1950.0 name - as in Willott et al. (2003)\\\\ TOOT 1252+3310   &WQ& Broad H$\\beta$ detected in follow-up spectroscopy (Willott et al., in prep)\\\\ \\hline \\end{tabular} \\label{tab2} \\end{center} \\end{table*} ",
        "conclusions": "The results of a detailed analysis of deep, high-resolution HST images of a sample of 41 radio galaxies which spans three decades in radio luminosity at $z\\simeq 0.5$ have been presented. The host-galaxy properties derived from two-dimensional image modelling have been discussed and compared to existing results in the literature. The principle conclusions of this study can be summarized as follows: \\begin{enumerate}  \\item{The morphologies of the host galaxies of all four radio-galaxy sub-samples are found to be largely consistent with a canonical elliptical-galaxy $r^{1/4}$ surface-brightness distribution. No trend is detected for host-galaxy morphology to vary significantly with radio luminosity.}   \\item{The luminosities of the host galaxies are found to be consistent with those of galaxies drawn from the bright end of the local cluster-galaxy luminosity function, in the range $0.7L^{\\star}<L<10L^{\\star}$, with a mean value of $3.2\\pm0.3 L^{\\star}$.}   \\item{The radio-galaxy hosts are found to lie on the same Kormendy ($\\mu_{e}-r_{e}$) relation as powerful $z\\simeq 0$ radio galaxies, after allowing for the effects of passive evolution of their stellar populations under the assumption of a high formation redshift ($z_{for}\\,\\gtsim$\\,3).}  \\item{By combining the new results on $z\\simeq 0.5$ 3CRR galaxies with previous results in the literature, it is found that there is no evidence for dynamical evolution among the hosts of 3C-class radio galaxies in terms of scalelength, luminosity and Kormendy relation in the redshift interval $0.0<z<0.8$. However, although no evidence is found for dynamical evolution among the hosts of the most powerful radio sources as a class, dynamical evolution of the individual host galaxies themselves cannot be ruled out.}  \\item{The scalelengths of the ZP5 host galaxies are found to be consistent with those of brightest cluster galaxies (BCGs) at the same cosmic epoch ($z\\simeq 0.5$). When compared to local Abell cluster BCGs the ZP5 host galaxies are found to be a factor of $\\simeq 1.5$ smaller on average. However, when the comparison is restricted to local clusters of Abell class 0 only, the ZP5 and BCG scalelength distributions are found to be statistically indistinguishable.}  \\item{Converting the derived host-galaxy luminosities into black-hole mass estimates using the local $M_{bh}-M_{bulge}$ relation predicts that the host galaxies harbour central black-holes in the mass range $10^{8.1}\\Msun<M_{bh}<10^{9.5}\\Msun$, with a geometric mean of $10^{8.87\\pm0.04}\\Msun$}.  \\item{Significant correlations ($\\simeq 3\\sigma$) are found between black-hole mass and extended low-frequency radio luminosity for sub-samples of the ZP5 objects which have either classical double (CD) radio structures or high-excitation nuclear spectra (HEG). However, at present we are unable to determine whether the objects which do not follow the same correlation, predominantly the lower-luminosity TOOT sub-sample, do so because of their LEG spectra, fat double/jet radio structures, or alternatively, as a result of selection effects biasing the TOOT sub-sample towards the inclusion of objects with higher host luminosities/black-hole masses.}  \\end{enumerate}"
    },
    "0403/astro-ph0403330_arXiv.txt": {
        "abstract": "{ It is expected that the next generation of high-contrast imaging instruments will deliver the first unresolved image of an extrasolar planet. The emitted thermal infrared light from the planet should show no phase effect assuming the planet is in thermal equilibrium. But the reflected visible light will vary with phase angle. Here we study the photometric variation of the reflected light with orbital phase of a ringed extrasolar planet. We show that a ring around an extrasolar planet, both obviously unresolved, can be detected by its specific photometric signature. A simple quantitative model is discussed, taking into account the basic optical and geometrical properties of the ringed planet. ",
        "introduction": "The discovery of extrasolar planets by radial velocity measurements has provided the first dynamical characteristics of planets (orbital elements and mass). The next step will be to investigate physical characteristics (albedo, temperature, radius etc.) and planet surroundings. Among the latter are the satellites and rings of planets. Here we present the first of a series of papers on the detectability of exoplanet surroundings, namely rings and satellites. It would seem {\\it a priori} that their detection will require imaging at very high angular resolution, beyond the capabilities of the most ambitious imaging projects such as Darwin/TPF (\\cite{leger_et_al96, angel_97, beichman99}). On the contrary, we show in these papers that angularly unresolved rings and satellites are detectable by the transit method and by the same imaging and polarimetric detection projects in preparation for extrasolar planets. In this paper we start the study of the detectability of rings.  The detection of rings would not constitute just an astrophysical curiosity, it would have some impact on the understanding of planetary systems and on the strategy for their detection.  - {\\it Dynamics of planetary systems.} The existence of a ring would reveal indirectly the existence of satellites around the planet. Indeed the latter (so-called confining or {\\it shepherd} satellites) confine the rings and prevent them to dilute rapidly. The mass of a confining satellite can be extremely small. For instance, the $3.4\\times10^{-7}M_{\\oplus}$ satellite Janus confines the outer edge of the Saturn A ring. In the absence of resonances, and depending on the ring viscosity, the confinement is, for a satellite mass $M_S$ and distance to the planet $d_S$, proportional to $M_S^2/d_S^{4}$ (\\cite{borderies_82}). In this sense, the detection of rings can be considered as an indirect detection of extremely small extrasolar bodies.   Satellites of giant planets in the habitable zone are also of interest because they are potential abodes of life (\\cite{williams_et_al1997}). In addition, it will be interesting to search for a correlation between the presence of rings and the planet orbit in order to understand if migration has an impact on the formation of rings: For instance a correlation of the frequency of rings with closer Jupiters could reveal that migration favors their formation (by a stronger disk-protoplanet mass exchange), while an anti-correlation could reveal that migration inhibits ring or satellite formation. These are just suggestions to be investigated by detailed formation models.    - {\\it Strategy for planet detection.} The radius $r_p$ of planets detected in the thermal infrared by the future Darwin/TPF space mission is expected to be inferred from the thermal flux $F_{p,IR}$ through the relation $F_{p,IR}\\propto r_p^2$ (\\cite{schneider01}, 2002). However, the latter relation holds only for spherical bodies, so that the presence of a ring (non-spherical by definition) would give a wrong value for $r_p$. In particular, one cannot exclude that small (e.g. Earth-like) planets have giant rings (Sect. \\ref{ringed_earth}). In addition, as will be explained in the discussion (Sect. \\ref{discussion}), a ring may hide the planet in some configurations during some parts of its orbital revolution (\\cite{desmarais_et_al2002a}, 2002b). We discuss in Sect. \\ref{large and thick ring} the operational consequences of this effect: In reflected light, if there is no planet detection at a given period of observation, this non-detection does not constitute a proof that there is no planet.  It is expected that the next generation of space or ground-based high-contrast imaging instruments (with adaptive optics, coronagraphy, pupil apodization) will be able to deliver the first image of an unresolved extrasolar planet. Two options are currently under consideration for TPF: thermal infrared and visible light. The emitted thermal infrared light from the planet should show no phase effect assuming the planet is in thermal equilibrium. But the reflected visible light will vary with phase angle, as should be shown by a broad-band photometric follow-up of the planet during its orbital motion.  Therefore we argue that it is of interest to study how the presence of a ring around a planet would influence its magnitude as a function of its orbital position. The aim of this paper is to show that the reflected light curve of a ringed planet is significantly different from that of a ringless planet, thus revealing the presence of a ring. The specific signature of a ring was discussed by Schneider (2001). Here a more quantitative model is discussed, in which the basic optical properties of the planet and the ring are taken into account, together with all geometrical parameters describing the ringed planet.  In Sect. \\ref{section_ring}, we first briefly review the known rings in the Solar System. In Sect. \\ref{section_model}, we describe our numerical model and list our different assumptions. Sect. \\ref{discussion} discusses first the case of the photometry of a lambertian ringless planet, so that we can always refer to this basic case during the discussion of our ringed planet simulations. We finally show selected sets of light curves of a ringed planet for different geometrical and optical parameters that demonstrate the specific photometric signature of a ring around an extrasolar planet. Following papers will discuss the detection of satellites and rings by direct imaging, transits, astrometry and polarimetry. ",
        "conclusions": "This work demonstrates that a ring around an extrasolar planet significantly affects the reflected light curve of an extrasolar planet during its orbital motion. A ring could thus be detected, although both planet and ring would obviously remain unresolved. This may be achieved simply by photometric monitoring of the planet reflected light. We identified the following signatures, which would require only moderate photometry accuracy to be observed (around $3\\sigma\\approx0.5$ mag):   i) {\\it Light curve dichotomy, with strong slope changes at the equinoxes}, due to the ring being alternately seen in reflection and in transmission.  ii) {\\it Light curve dependence on wavelength} in dual-band photometry (methane band and continuum for instance), or {\\it spectral variations} if spectroscopy is possible.  iii) The {\\it $\\phi$-shift} of the light curve extrema, due to the longitude of the ring obliquity.  iv) {\\it Temporary extinction of the planet} during the orbital motion, due to the rising shadow of the ring on the planet.  v) {\\it High brightness in the reflected light}, leading to abnormally large planet radius, or/and abnormally low mass density.  vi) {\\it Disagreement between the planet radii} measured from reflected light and thermal infrared emission.  Although future space missions studies concentrate mainly on infrared instruments for technical and scientific reasons, this work shows the additional interest of shorter wavelengths (visible band) for extrasolar planet characterization."
    },
    "0403/astro-ph0403089_arXiv.txt": {
        "abstract": "A Sloan Digital Sky Survey (SDSS) source sample containing 44,200 quasar redshifts is examined. Although arguments have been put forth to explain some of the structure observed in the redshift distribution, it is argued here that this structure may just as easily be explained by the presence of previously predicted preferred redshifts. ",
        "introduction": "It was realized early in the SDSS survey \\citep{ric02} that quasar redshifts near z = 2.7 would be under-sampled simply because at that redshift the colors of stars resemble those of quasars and it would be difficult to keep the quasar identification efficiency high. Based on the knowledge gained during the commissioning period, and the results from the first 2555 quasars, the target selection algorithm was refined to the point where it was felt by 2000 November that it was appropriate for use in obtaining a statistical sample of quasars \\citep{ric02}. However, after additional observations were carried out, and many more sources with redshifts above z = 3 were obtained, it became apparent that there was a second dip in the redshift distribution at z = 3.5. This was again deemed to be a problem with the source selection algorithm, although it is not clear how this was determined, since a real source deficiency at this redshift would presumably give the same result. It may, however, have been decided from the evidence for a trough visible in the $i$-redshift plot \\citep[Fig 3]{sch03} which, when redshifts are assumed to be purely cosmological, can only be explained by a selection effect. On the other hand, it is exactly what would be expected if the redshifts are intrinsic. However, further changes to the target selection algorithm were made and it was \\em finalized \\em a second time on 2001 August 24. Spectra taken since that date were deemed to constitute the basis for the statistical sample of SDSS quasars \\citep{ric02}. The first release of SDSS quasar redshift data (DR1) contained the 16,713 redshifts measured prior to 2001 October \\citep{sch03}. Although the target selection algorithm was being refined throughout most of the period the DR1 observations were made, the sample considered here contains an additional $\\sim 27,500$ redshifts that were obtained since 2001 October, after the algorithm was finalized.  The primary SDSS quasar science goals were to study (1), the evolution of the quasar luminosity function and (2), the spatial clustering of quasars, as a function of redshift. However, a reasonably complete statistical sample of quasar redshifts can also be used to examine the long-standing controversy surrounding the very nature of quasar redshifts. Peaks in the distribution of quasar redshifts have long been claimed as possible evidence for the existence of preferred redshifts \\citep{kar71,kar77,bur01,bel02a,bel02c}. However, only source samples with a few hundred redshifts have been available to test this claim. Now that surveys containing tens of thousands of quasar redshifts are available, if selection effects can be properly accounted for, although this will not be easy, any peaks still visible in the redshift distribution could be a much more reliable indicator for the existence of preferred redshifts. Alternatively, if predicted redshift peaks do not appear, it can mean either that intrinsic redshifts do not exist, or that they have been smeared out by a significant Doppler component.  There are currently only two models that have been proposed that clearly define the preferred redshift values expected. It is of interest now to see if one of these models might equally well explain the observed redshift distribution of 44,200 SDSS quasars. ",
        "conclusions": "In this examination of 44,200 SDSS quasar redshifts it has been found that there are at least 5 peaks in the redshift distribution that correspond to the pure harmonics of z = 0.62 predicted by eqn 2. Although valid arguments have been put forward to explain the redshift valleys at z = 2.7 and 3.5, there is no proof yet that they have correctly explained the causes of these low source counts, since eqn 2 predicts low densities at both these redshifts, reducing the effect of any selection effect. Eqn 2 also predicts the low source count seen near z = 4.03 where no selection effect has yet been proposed.  It is concluded here that the intrinsic redshift model given by eqn 2 can better explain the structure present in the distribution of these 44,200 SDSS quasar redshifts than can the selection effects suggested by others. SDSS investigators are urged to keep this in mind if they hope to insure that real density fluctuations are not inadvertently removed. If we are to obtain a truly reliable statistical sample, adjustments that affect the SDSS redshift distribution must not involve the assumption that quasar redshifts are purely cosmological. %"
    },
    "0403/astro-ph0403262_arXiv.txt": {
        "abstract": "We analyze the Far-UV/UV spectra of four central stars of planetary nebulae with strong wind features --- NGC~2371, Abell~78, IC~4776 and NGC~1535,  and derive their photospheric and wind parameters by modeling high-resolution FUSE (Far-Ultraviolet Spectroscopic Explorer) data in the Far-UV and HST-STIS and IUE data in the UV with spherical non-LTE line-blanketed model atmospheres.  Abell~78 is a hydrogen-deficient transitional [WR]-PG~1159 object, and we find NGC~2371 to be in the same stage, both migrating from the constant-luminosity phase to the white dwarf cooling sequence with $\\Teff \\simeq 120$~kK, $\\Mdot \\simeq 5\\E{-8}$~\\Msunyr.  NGC~1535 is a ``hydrogen-rich'' O(H) CSPN, and the exact nature of IC~4776 is ambiguous, although it appears to be helium burning.  Both objects lie on the constant-luminosity branch of post-AGB evolution and have $\\Teff \\simeq 65$~kK, $\\Mdot \\simeq 1\\E{-8}$.  Thus, both the H-rich and H-deficient channels of PN evolution are represented in our sample. We also investigate the effects of including higher ionization stages of iron (up to \\FeX) in the model atmosphere calculations of these hot objects (usually neglected in previous analyses), and find iron to be a useful diagnostic of the stellar parameters in some cases.  The Far-UV spectra of all four objects show evidence of hot ($T \\sim 300$~K) molecular hydrogen in their circumstellar environments. ",
        "introduction": "\\label{sec:intro}  After leaving the asymptotic giant branch (AGB), low and intermediate-mass stars become central stars of planetary nebulae (CSPN).  The spectral types of the central stars are numerous and diverse (\\eg, O, Of, sdO, [WR], [WR]-PG~1159, as well as white dwarf-types), but all of these can be generally divided into two groups based on surface abundances: hydrogen-rich (which show obvious hydrogen lines in their spectra) and hydrogen-deficient (which don't). H-deficient stars are in a post-helium flash, helium-burning phase, while the H-rich stars may be either hydrogen or helium burning.  About 10-20~\\% of CSPN are H-deficient (\\citealp{demarco:02,koesterke:98b} and references therein).  The two groups are thought to represent different ``channels'' of CSPN evolution, terminating with either H-rich DA type or He-rich DO type white dwarfs (\\eg, \\citealp{napiwotzki:99}).  During the CSPN phase the star sheds mass in the form of a stellar wind as it evolves toward the white dwarf cooling sequence.  It is generally believed that the winds of these objects are driven by radiation pressure, as with population-I O-stars and WR stars.  If the driving of the wind is mainly dictated by the stellar luminosity, the wind parameters can, in principle, be used to determine the distances to CSPN \\citep{kudritzki:99,kudritzki:00,tinkler:02}.  The spectra of the majority of the H-deficient class have strong wind signatures, show no photospheric lines and are very similar to those of WC stars --- the carbon-rich Wolf-Rayets which are evolved massive stars.  The low-mass central stars are termed ``[WC]'' stars to denote their association with planetary nebulae.  ``PG~1159''- type stars are extremely hot, H-deficient white dwarf-type objects, thought to represent an entry point onto the white dwarf cooling sequence.  They show mainly absorption lines in the optical spectra but possess a few wind lines in the UV.  Because PG~1159 objects have similar abundances to [WC] stars and about half of the PG~1159 stars are CSPN \\citep{werner:01}, they are thought to be descendants of the [WC]s. Linking these two classes is a rare group of objects named ``[WC]-PG~1159'' stars, whose spectra show both emission lines and some absorption lines.  Only three definite cases of such stars are known so far: Abell~78, Abell 30 and Longmore 4 \\citep{koesterke:01}. \\citet{parthasarathy:98} have posited that all ``weak emission line stars'' ([WELS]) are actually [WC]-PG~1159, based on optical spectra, but \\citet{werner:01} cautioned against this statement until better spectra of the sample is available.  The H-rich CSPN sample of \\citet{mendez:88} are similar in that their spectra also feature absorption lines as well as UV wind emission features.  [WC], [WC]-PG~1159, and PG~1159 stars all reside in the same area of the H-R diagram (HRD).  In terms of evolution, this region corresponds to the bend where the star leaves the constant-luminosity post-AGB phase and is transitioning to the white dwarf cooling sequence. [WC] and PG~1159 stars lie close to each other in the H-R diagram, yet have very different spectra (in the optical, a [WC] spectrum is littered with strong emission lines, while that of a PG~1159 presents an absorption line spectrum).  This suggests a dramatic drop in mass-loss rates over a short time of evolution \\citep{koesterke:98b}.  The commonly accepted evolutionary sequence for H-deficient CSPN is: [WC]\\ra [WC]-PG~1159\\ra PG~1159\\ra DO WD \\citep{hamann:96}.  This sequence spans the time where the stellar winds are in the process of ``turning-off''.  Therefore, understanding the objects in this part of the H-R diagram is not only a prerequisite for understanding the evolution of central stars, but will also give insight into the driving of stellar winds in general.  In the classification of [WR]-type objects, wind parameters are more fundamental discriminators than stellar temperature \\citep{crowther:99a, acker:03}.  As the central star's winds die down, the wind features in the FUV/UV are the last to fade \\citep{koesterke:98a}.  Therefore, to estimate accurate parameters of the extended atmospheres this phase is best studied in the FUV/UV.  In this paper, we present a FUV/UV-based spectral analysis of four CSPN which show strong wind signatures - NGC~2371, Abell~78, IC~4776, and NGC~1535.  We model the FUV/UV spectra of these stars using stellar atmospheres codes to determine the wind parameters such as \\Teff, \\Mdot, and \\vinf\\ and discuss their evolutionary implications. Abell~78 has been classified as a [WO]-PG~1159 transition object \\citep{mendez:91,crowther:98}, and we argue that NGC~2371 also belongs to this class.  NGC~1535 is an O(H) star belonging to the H-rich sample of \\citet{mendez:88}. The FUV spectra of IC~4776 is similar to that of NGC~1535, but we find evidence that this object may be a He-burner.  Comparing with evolutionary calculations, our derived parameters place the latter two objects along the constant-luminosity path of the HRD, while the former two lie on the bend marking the transition from that phase to the white dwarf cooling sequence.   This paper is arranged as follows.  The observations and data reduction are described in \\S~\\ref{sec:obs}.  A comparison of the spectra of the objects is presented in \\S~\\ref{sec:description}.  Our models and parameter determinations are described in \\S~\\ref{sec:modeling}.  The implications of our results are discussed in \\S~\\ref{sec:discussion} and our conclusions in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have performed a FUV/UV spectral analyses on four CSPN which display combined absorption and wind line spectra, a sign that the winds of these objects are in the process of switching off.  Two objects are H-deficient (NGC~2371 and Abell~78), one is H-rich (NGC~1535), and for one our analysis is inconclusive (IC~4776).  Thus both channels of the CSPN phase are represented. We derived the stellar wind parameters of the sample PN by modeling their FUV and UV spectra, where strong wind signatures are seen.  NGC~1535 and Abell~78 have been previously analyzed in the optical using plane-parallel models(\\citealp{werner:03}, \\citealp{mendez:88}), and the latter also with wind models \\citep{koesterke:98b}.  Our parameters are in good agreement with the previous results.  The most striking feature in all the FUV spectra is the \\OVI\\ doublet. In massive O-type stars, this ion originates from X-rays due to shocks in the radiatively driven winds \\citep{bianchi:02}.  We are able to fit this feature in our sample CSPN without the inclusion of X-rays due to their higher effective temperatures.  Two objects have $\\Teff \\simeq 65$~kK, with $\\Mdot\\sim10^{-8}$~\\Msunyr.  NGC~1535 has been classified as a Hydrogen-rich O(H) star based on its optical spectrum, and we are able to match its FUV/UV spectrum using solar abundances.  From the spectrum of IC~4776, we cannot constrain the abundances, but comparison with evolutionary tracks suggest this object is a He-burner.  Our models for these objects show a jump in the ionization structure of the CNO elements (and S and P) between 55-70~kK.  It is in this temperature range that \\citet{tinkler:02} find that H-rich CSPN exhibit a large jump in \\Mdot\\ by an factor of 10-100, which they suspect is due to such a bi-stability jump.  This supports the idea of \\citet{vink:01} that a bi-stability jump at such high temperatures may be due to CNO elements.  The two H-deficient objects Abell~78 and NGC~1271 are significantly hotter ($\\Teff>110$~kK), with $\\Mdot\\sim 10^{-7}$~\\Msunyr.  Our derived parameters for the [W01]-PG~1159 star Abell~78 generally agree with those of past analyses.  We derive similar parameters for NGC~2371, and suggest that it is a [WR]-PG~1159 also, but its wind is more ionized and shows no \\OV\\ features (a [WO0]-PG~1159?).  We find evidence of iron deficiency in both of these objects, supporting the findings of \\citet{werner:03} in Abell~78 and of \\citet{miksa:02} in PG~1159 stars.  They lie on the bend in the theoretical evolutionary tracks between the constant luminosity phase and the WD cooling sequence, having post-AGB ages of 10-15~kyr.  For [WC] stars, it seems that as the star evolves away from the post-AGB phase through the [WC] sequence, the temperature and terminal velocity of the wind increase as the wind density decreases \\citep{acker:03}.  Except for Abell~78, the post-AGB ages predicted by the evolutionary models are typically 2-4 times lower than their kinematic ages. However, kinematic ages are lower limits to the post-AGB age (as the expansion is slower in the initial phase) thus the actual discrepancy may be smaller.  Our FUV/UV analysis has provided wind parameters for H-rich and H-deficient CSPN at the stage in post-AGB evolution where the winds are fading.  However, because a precise determination of the mass is not possible, the objects do not necessarily represent the same evolutionary sequence.  This work has also provided information on the interstellar and circumstellar environment from our measurements of the column densities and temperatures of \\HI\\ and \\Htwo\\ along the sight-lines.  Our determinations of the \\HI\\ column density and \\EBMV\\ imply that the relationship between these two quantities in the circumstellar (PN) environment of CSPN differs from that of the ISM, having lower dust-to-gas ratios (probably due to the destruction of dust by the radiation field).  The high resolution of the FUSE data allow us to detect hot \\Htwo\\ associated with the nebulae.  For all four objects in our sample, a single component of \\Htwo\\ gas at typical ISM temperatures (\\ie, $T \\sim 80$~kK) was not adequate to fit the absorption spectrum.  A second, hotter ($T\\sim300$~kK) component was necessary, which we assume to be associated with the circumstellar environment.  With the advent of many FUSE observations of CSPN, it is becoming apparent that hot circumstellar \\Htwo\\ is not uncommon CSPN at quite different evolutionary stages: from old PN with white dwarf nuclei \\citep{herald:02} to young, compact PN \\citep{herald:04a}.  The objects in this paper lie either along the constant luminosity section of the post-AGB evolutionary tracks, or on the transition bend to the WD cooling sequence, and thus represent intermediate stages to the previously mentioned cases.  Given the intense UV radiation fields emitted by the CSPN, it is likely that the nebular \\Htwo\\ exists in clumps, shielded by neutral and ionized hydrogen, as appears to be the case in the Helix nebula \\citep{speck:02}."
    },
    "0403/astro-ph0403054_arXiv.txt": {
        "abstract": "Radio continuum observations at 20 and 6~cm of the highly inclined Virgo spiral galaxy NGC~4522 are presented. Both, 20 and 6~cm total emission distributions are asymmetric with an extended component to the west where extraplanar atomic gas and H$\\alpha$ emission are found. The 6~cm polarized emission is located at the eastern edge of the galactic disk. Its peak is located about 1~kpc to the east of the total emission peak. We argue that this phenomena is a characteristic feature for cluster galaxies which are experiencing significant pressure from the intracluster medium. The degree of polarization decreases from the east to the west. The flattest spectral index between 20 and 6~cm coincides with the peak of the 6~cm polarized emission. These findings are consistent with a picture of a large scale shock due to ram pressure located at the east of the galaxy where cosmic rays are accelerated. We conclude that it is likely that the galaxy experiences active ram pressure. ",
        "introduction": "The Virgo cluster spiral galaxy NGC~4522 is one of the few Virgo galaxies where we can directly observe the effects of ram pressure due to the galaxy's rapid motion in the intra-cluster medium. Kenney \\& Koopmann (1999) observed this galaxy in the optical and the H$\\alpha$ line with the WIYN telescope. Whereas the highly inclined, old stellar disk appears symmetric, the distribution of ionized gas is highly asymmetric. The H$\\alpha$ disk is sharply truncated beyond 0.35\\,$R_{25}$. Ten percent of the H$\\alpha$ emission arises from extraplanar H{\\sc ii} regions that are exclusively located to the west of the galactic disk. Kenney \\& Koopmann (1999) argue that this ionized gas distribution is reminiscent of a bow shock morphology, which suggests that the gas is pushed to the west by ram pressure.  Vollmer et al. (2000) obtained an H$\\alpha$ velocity field with the Fabry-Perot Interferometer at the Observatoire de Haute Provence. The galactic disk shows a symmetrically rising rotation curve. The velocities of the extraplanar emission regions are not part of this regular rotation. Their kinematic behaviour cannot be reproduced by rotation within a gravitational potential of any known disk or halo model. Thus this gas is located out of the galactic plane and/or is accelerated/decelerated.  NGC~4522 has a large line-of-sight velocity with respect to the cluster mean ($\\sim$1300~km\\,s$^{-1}$) and is located at a distance of 3.3$^{\\rm o}$ ($\\sim$1~\\footnote{We use a distance of 17~Mpc to the Virgo cluster}Mpc) south from the Virgo cluster center (M87). The intracluster medium (ICM) density at these distances might not be high enough to strip the gas at a galactic radius of 0.35\\,$R_{25}$. This is why Vollmer et al. (2000) argued that NGC~4522 has been severely affected by ram pressure in the past and the extraplanar filaments represent gas that has been pushed to larger distances from the galaxy center and is now falling back onto the galactic disk.  Recently, Kenney et al. (2003) observed NGC~4522 in the H{\\sc i} 21~cm line with the VLA. They found an atomic gas distribution that is similar to that of the ionized gas: a truncated H{\\sc i} disk with two high column density blobs in the north-west and the south-west. The H{\\sc i} velocity field of the extraplanar emission regions shows clear deviations from the overall rotation pattern.  Even with the distributions and velocity fields of the atomic and ionized gas it is still unclear if the galaxy is experiencing ongoing ram pressure stripping or if the extraplanar gas is falling back to the galaxy after a past ram pressure stripping event. On the one hand the galaxy's location and radial velocity excludes a simple radial orbit advocated by Vollmer et al. (2000), on the other hand with the ongoing stripping scenario it is not clear if ram pressure is strong enough at such a large distance from the cluster center where the ICM density is lower than average.  Observations of the polarized radio continuum radiation can give important and complementary information on the gas dynamics as has been shown for the barred galaxy NGC~1097 (Beck et al. 1999). This radiation traces the ordered large scale ($\\sim$1~kpc) magnetic field in galactic ISM. This magnetic field is very sensitive to (i) compression and (ii) shear motion, which are both difficult to detect in radial velocity fields. Thus, enhanced polarized radio continuum radiation can be a sign of compression or strong shear motion in the gas. Both kinematic features are present during a ram pressure stripping event (see e.g. Vollmer et al. 2000): a galaxy on a highly eccentric orbit within the cluster experiences compression due to ram pressure when it passes through the cluster core (compression phase); after the core passage ram pressure ceases and the gas starts to fall back onto the galactic disk, which gives rise to strong shear motions.  Observationally, there are only two Virgo cluster galaxies observed deeply in the polarized 6~cm radio continuum, because this is very time consuming (NGC~4254, Soida et al. 1996, and NGC~4654, Chyzy et al. in prep.). Both galaxies show an asymmetric distribution of polarized radio continuum emission.  In order to interpret the polarization data, numerical modeling is important, because different dynamical features, such as compression or shear, can produce polarized emission maxima. Otmianowska \\& Vollmer (2003) solved the induction equation using the velocity fields of Vollmer et al. (2002) in order to calculate the evolution of the large scale magnetic field during a stripping event. In a second step they calculated the polarized radio continuum emission and made maps of its evolution that can be directly compared to observations. We will apply this method to a dynamical model designed for NGC~4522 in a forthcoming paper.  In this article we present radio continuum observations at 6~cm and 20~cm obtained with the VLA. The outline of the article is as follows. The observations are described in Sec.~1. The results are discussed in Sect.~2. We compare our results to radio observation of field galaxies (Sect.~4) and compare our data with existing H{\\sc i} and H$\\alpha$ maps (Sect.~5). The discussion in Sect.~6 is followed by the conclusions in Sect.~7. ",
        "conclusions": "We present VLA D array observations at 6 and 20~cm of the highly inclined Virgo spiral galaxy NGC~4522. The results are \\begin{enumerate} \\item the 20~cm and 6~cm total emission distributions are asymmetric; they are more extended to the west, where extraplanar H$\\alpha$ and H{\\sc i} emission is found, \\item the 6~cm polarized emission is located at the eastern edge of the galactic disk; its peak is displaced to the east of the total emission peak, \\item the degree of polarization increases from the galaxy center towards the eastern peak and towards the northeastern edge of the disk, \\item the spectral index between 20 and 6~cm decreases from west to east and its maximum coincides with the peak of the 6~cm polarized emission, \\item polarized radio continuum emission is a powerful tool to detect interactions of spiral galaxies with their cluster environment. \\end{enumerate} Points 3 and 4 are consistent with a picture of a large scale shock due to ram pressure located at the east of the galaxy possibly associated with particle acceleration. We argue that the asymmetry of the 6~cm polarized emission is characteristic for cluster galaxies and is a sign of interaction with its cluster environment. We conclude that it is probable that the galaxy experiences active ram pressure, but it is not clear whether we are observing it before, during, or after peak ICM pressure."
    },
    "0403/astro-ph0403324_arXiv.txt": {
        "abstract": "I briefly review our current understanding of dark matter and dark energy. The first part of this paper focusses on issues pertaining to dark matter including observational evidence for its existence, current constraints and the `abundance of substructure' and `cuspy core' issues which arise in CDM. I also briefly describe MOND. The second part of this review focusses on dark energy. In this part I discuss the significance of the cosmological constant problem which leads to a predicted value of the cosmological constant which is almost $10^{123}$ times larger than the observed value $\\la/8\\pi G \\simeq 10^{-47}$GeV$^4$. Setting $\\la$ to this small value ensures that the acceleration of the universe is a fairly recent phenomenon giving rise to the `cosmic coincidence' conundrum according to which we live during a special epoch when the density in matter and $\\la$ are almost equal. Anthropic arguments are briefly discussed but more emphasis is placed upon dynamical dark energy models in which the equation of state is time dependent. These include Quintessence, Braneworld models, Chaplygin gas and Phantom energy. Model independent methods to determine the cosmic equation of state and the Statefinder diagnostic are also discussed. The Statefinder has the attractive property $\\atridot/a H^3 = 1 $ for LCDM, which is helpful for differentiating between LCDM and rival dark energy models. The review ends with a brief discussion of the fate of the universe in dark energy models. ",
        "introduction": "\\index{dark matter}  Observations of the cosmic microwave background (CMB) and the deuterium abundance in the Universe suggest that $\\om_{\\rm baryon}h^2 \\simeq 0.02$, or $\\om_{\\rm baryon} \\simeq 0.04$ if the current Hubble expansion rate is $h = H_0/100 {\\rm km/sec/Mpc} = 0.7$. Although $\\om_{\\rm baryon}$ is much larger than the observed mass in stars, $\\om_{\\rm stars} \\simeq 0.005$ \\footnote{This suggests that most of the baryonic matter at $z=0$ is not contained in stars but might be contained in hot gas \\cite{bosma03}.}, it is nevertheless very much smaller than the total energy density in the universe inferred from the observed anisotropy in the cosmic microwave background \\cite{spergel03} \\beq \\Omega_{\\rm total} \\equiv \\frac{8\\pi G\\rho_{\\rm total}}{3H^2} = 1.02 \\pm 0.02~. \\eeq Both dark matter and dark energy are considered essential missing pieces in the cosmic jigsaw puzzle \\beq \\om_{\\rm total} - \\om_{\\rm baryons} = ~? \\eeq Although the nature of neither dark matter (DM) nor dark energy (DE) is currently known, it is felt that both DM and DE are non-baryonic in origin, and that DM is distinguished from DE by the fact that the former clusters on sub-Megaparsec scales (in order to explain galactic rotation curves) whereas the latter has a large negative pressure (and can make the universe accelerate). In addition there is strong evidence to suggest that \\beq \\Omega_m \\simeq 1/3, ~~ \\Omega_{\\rm DE} \\simeq 2/3~. \\eeq  \\n In these lectures I will briefly review some properties of both dark matter and dark energy.  Though the observational evidence favouring a flat Universe with $\\Omega_{\\rm total} \\simeq 1$ is fairly recent, the nature of the `unseen' component of the universe (which dominates its mass density), is a long-standing issue in modern cosmology. Indeed, the need for dark matter was originally pointed out by Zwicky (1933) who realized that the velocities of individual galaxies located within the Coma cluster were quite large, and that this cluster would be gravitationally bound only if its total mass substantially exceeded the sum of the masses of its component galaxies. For clusters which have relaxed to dynamical equilibrium the mean kinetic and potential energies are related by the virial theorem \\cite{coles} \\beq K + \\frac{U}{2} = 0~, \\eeq where $U \\simeq -GM^2/R$ is the potential energy of a cluster of radius $R$, $K \\simeq 3 M\\langle v_r^2\\rangle/2$ is the kinetic energy and $\\langle v_r^2\\rangle^{1/2}$ is the dispersion in the line-of-sight velocity of cluster galaxies. (Clusters in the Abell catalogue typically have $R \\simeq 1.5 h^{-1}$ Mpc.) This relation allows us to infer the mean gravitational potential energy if the kinetic energy is accurately known. The mass-to-light ratio in clusters can be as large as $M/L \\simeq 300 M_\\odot/L_\\odot$. However since most of the mass in clusters is in the form of hot, x-ray emitting intracluster gas, the extent of dark matter in these objects is estimated to be $M/M_{\\rm lum} \\simeq 20$, where $M_{\\rm lum}$ is the total mass in luminous matter including stars and gas.  \\begin{figure}[ht] \\centering \\includegraphics[width=10cm]{five.eps} \\caption{\\footnotesize The observed rotation curve of the dwarf spiral galaxy M33 extends considerably beyond its optical image (shown superimposed); from Roy ~\\protect\\cite{dproy}. } \\label{fig:flat} \\end{figure}  In individual galaxies the presence of dark matter has been convincingly established through the use of Kepler's third law \\beq v(r) = \\sqrt{\\frac{G M(r)}{r}} \\eeq to determine the `rotation curve' $v(r)$ at a given radial distance from the galactic center.  Observations of galaxies taken at distances large enough for there to be no luminous galactic component indicate that, instead of declining at the expected rate $v \\propto r^{-1/2}$ true if $M \\simeq {\\rm constant}$, the velocity curves flattened out to $v \\simeq {\\rm constant}$ implying $M(r) \\propto r$ (see fig \\ref{fig:flat}). This observation suggests that the mass of galaxies continues to grow even when there is no luminous component to account for this increase. Velocity curves have been compiled for over 1000 spiral galaxies usually by measuring the 21 cm emission line from neutral hydrogen (HI) \\cite{dark1,dark2}. The results indicate that $M/L = (10 - 20) M_\\odot/L_\\odot$ in spiral galaxies and in ellipticals, while this ratio can increase to $M/L \\simeq (200 - 600) M_\\odot/L_\\odot$ in low surface brightness galaxies (LSB's) and in dwarfs. For instance, a recent measurement of the Draco dwarf spheroidal galaxy located at a distance of only 79 kpc from the Milky Way shows the presence of a considerable amount of dark matter $M/L\\vert_{\\rm Draco} = 440 \\pm 240 M_\\odot/L_\\odot$ \\cite{kleina01} ! It is interesting that the total mass of an individual galaxy is still somewhat of an unknown quantity since a turn around to the $v \\propto r^{-1/2}$ law at large radii has not been convincingly observed.  An important difference between the distribution of dark matter in galaxies and clusters needs to be emphasised: whereas dark matter appears to {\\em increase} with distance in galaxies, in clusters exactly the reverse is true, the dark matter distribution actually {\\em decreases} with distance. Indeed, for certain dwarfs (such as DD0154) the rotation curve has been measured to almost 15 optical length scales indicating that the dark matter surrounding this object is extremely spread out (see also figure \\ref{fig:flat}). A foreground cluster, on the other hand, acts as a gravitational lens which focuses the light from background objects such as galaxies and QSO's thereby allowing us to determine the depth of the cluster potential well. Observations of strong lensing by clusters indicate that dark matter is strongly concentrated in central regions with a projected mass of $10^{13} - 10^{14} M_\\odot$ being contained within 0.2 - 0.3 Mpc of the central region. As we shall see later, this observation may prove to be problematic for alternatives to the dark matter hypothesis such as the Modified Newtonian Dynamics (MOND) approach of Milgrom \\cite{milgrom}.  As discussed earlier, the fact that only $4\\%$ of the cosmic density is baryonic suggests that the dark matter which we are observing could well be non-baryonic in origin. The need for non-baryonic forms of dark matter gets indirect  support from the fact that baryonic models find it difficult to grow structure from small initial conditions and hence to reconcile the existence of a well developed cosmic web of filaments, sheets and clusters at the present epoch with the exceedingly small amplitude of density perturbations ($\\delta\\rho/\\rho \\sim 10^{-5}$ at $z \\simeq 1,100$) inferred from COBE measurements and more recent CMB experiments. Indeed, it is well known that, if the effects of pressure are ignored, linearized density perturbations in a spatially flat matter dominated universe grow at the rate $\\delta \\propto t^{2/3} \\propto (1+z)^{-1}$, where $1+z = a_0/a(t)$ is the cosmological redshift. (Contrast this relatively slow growth rate with the exponential `Jeans instability' of a static matter distribution $\\delta \\propto \\exp{\\sqrt{4\\pi G\\rho} t}$.) In a baryonic universe, the large radiation pressure (caused by thompson scattering of CMB photons off electrons) ensures that density perturbations in the baryonic component can begin growing only after hydrogen recombines at $z \\simeq 1,100$ at which point of time baryons and radiation decouple. Requiring $\\delta > 1$ today implies $\\delta > 10^{-3}$ at recombination, which contradicts CMB observations by over an order of magnitude ! In non-baryonic models on the other hand, the absence of any significant coupling between dark matter and radiation allows structure to grow much earlier, significantly before hydrogen in the universe has recombined. After recombination baryons simply fall into the potential wells created for them by the dominant non-baryonic component. As a result a universe with a substantial non-baryonic component can give rise to the structure which we see today from smaller initial fluctuations.  The growth of structure via gravitational instability depends both upon the nature of primordial perturbations (adiabatic/isocurvature) and upon whether the dark matter species is {\\rm hot} or {\\rm cold}. The issue of density perturbations has been discussed in considerable detail by Ruth Durrer at this school and I shall not touch upon this important topic any further. Let me instead say a few words about hot and cold dark matter. Non-baryonic {\\em Hot Dark Matter} (HDM) particles are assumed to have decoupled from the rest of matter/radiation when they were relativistic and so have a very large velocity dispersion (hence called `hot'). {\\em Cold Dark Matter} (CDM) particles, on the other hand, have a very small velocity dispersion and decouple from the rest of matter/radiation when they are non-relativistic. The free-streaming (collisionless phase mixing) of non-baryonic particles as they travel from high density to low density regions (and vice versa) introduces an important length scale called the `free-streaming distance' $\\lambda_{\\rm fs}$ -- which is the mean distance travelled by a relativistic particle species until its momentum becomes non-relativistic. In both HDM and CDM the processed {\\em final} spectrum of density perturbations differs from its initial form. In the case of HDM this difference arises because fluctuations on scales smaller than $\\lambda_{\\rm fs}$ are wiped out due to free streaming with the result that the processed final spectrum has a well defined cutoff on scales smaller than $\\lambda \\sim \\lambda_{\\rm fs}$. Perhaps the best example of HDM is provided by a light neutrino of mass about $30$ eV. In this case $\\lambda_{\\rm fs} \\simeq 41(30 {\\rm eV}/m_\\nu)$ Mpc with the result that large proto-pancakes having masses comparable to those of rich clusters of galaxies $M \\sim 10^{15} M_\\odot$ are the first objects to form in HDM. Smaller objects (galaxies) are formed by the fragmentation of the proto-pancake. This {\\em top-down} scenario for structure formation was originally suggested by Zeldovich and coworkers in connection with adiabatic baryonic models and subsequently applied to HDM. It has since fallen out of favour mainly due to the strong observational constraints on the mass of the neutrino $\\sum_{\\nu_i}m_{\\nu_i} < 0.7$ eV and on the relic neutrino density $10^{-3} \\lleq \\Omega_\\nu h^2 \\lleq 10^{-1}$ \\cite{elgaroy02,spergel03,ellis03,minaka02}. It also faces considerable difficulty in forming structure sufficiently early to explain the existence of galaxies and QSO's at high redshifts.  In contrast to HDM, constituents of CDM have a much smaller free-streaming distance. Because of this small scales are the first to go non-linear and gravitational clustering proceeds in a {\\em bottom up} fashion in this scenario. A key quantity defining gravitational clustering is the power spectrum of density perturbations $P(k) \\equiv |\\delta_k|^2$, which is related to the mean square density fluctuation via \\beq \\bigg\\langle \\left (\\frac{\\delta\\rho}{\\rho}\\right )^2 \\bigg\\rangle = 4\\pi \\int_0^\\infty P(k) k^2 dk~. \\label{eq:pow1} \\eeq Inflationary models predict $P_i(k) \\propto k^n$, $n \\simeq 1$, at an early epoch. As the universe expands the power spectrum gets modified. The `processed' final spectrum depends upon the nature of dark matter, the epoch of matter-radiation equality and other cosmological quantities. The final and initial spectra are related through a transfer function \\beq P_f(k) = P_i(k)\\times T^2(k)~. \\label{eq:pow_spec} \\eeq CDM-type spectra have the following approximate form of the transfer function \\cite{sahni84,ss84,sc95} \\beq T(k) = \\left (1 + \\frac{A k^2}{\\log{(1 + Bk)}}\\right )^{-1}~. \\label{eq:transfer} \\eeq Equations (\\ref{eq:pow_spec}) \\& (\\ref{eq:transfer}) illustrate the `turn around' of the power spectrum from its primordial scale invariant form $P(k) \\propto k$ on the largest scales to $P(k) \\propto k^{-3}\\log^2k$ on small scales. (The precise location of the turn-around and the amplitude of $P(k)$ depend upon specific details of the cosmological model, see for instance \\cite{bbks86}.)  The `standard' cold dark matter (SCDM) paradigm, which assumed that $\\Omega_{\\rm CDM} = 1$, was introduced during the early 1980's at roughly the same time when HDM was perceived to be in trouble (see \\cite{kolbt90,khlopov,coles,sc95} for references to earlier work on this subject). Although SCDM was very successful in explaining a host of observational details, it was clear already a decade ago, that the processed power spectrum of SCDM lacked sufficient power on large scales and so fell short of explaining the angular two point correlation function for galaxies on scales $\\sim 50$ Mpc \\cite{efstath}. The relevant cosmological quantity in this case is the shape of the power spectrum of density perturbations, which for CDM-like models, can be characterised by the `shape parameter' $\\Gamma = \\Omega_m h$. SCDM models with $\\Omega_m = 1$ and the HST-determined value $h \\simeq 0.7$ predict $\\Gamma \\simeq 0.5$ which is much larger than the observed value $\\Gamma = 0.207 \\pm 0.030$ inferred from observations of galaxy clustering in the sloan digital sky survey (SDSS) \\cite{pope}. A modification of SCDM called LCDM assumes that, in addition to CDM the universe consists of a smoothly distributed component called a cosmological constant or a Lambda-term. LCDM models with $h \\simeq 0.7$ and $\\Omega_m = 0.3$ predict a smaller value for the shape parameter, $\\Gamma \\simeq 0.2$, and the resulting amplitude and shape of the power spectrum is  in excellent agreement with several different sets of observations as demonstrated in figure \\ref{fig:power}.  \\begin{figure}[ht] \\centering \\includegraphics[width=10cm]{kspace.ps} \\caption{\\footnotesize The power spectrum inferred from observations of large scale structure, the Lyman$\\alpha$ forest, gravitational lensing and the CMB. The solid line shows the power spectrum prediction for a flat scale-invariant LCDM model with $\\Omega_m = 0.28$, $\\Omega_b/\\Omega_m = 0.16$, $h = 0.72$; from Tegmark et al. ~\\protect\\cite{tegmark03a}. } \\label{fig:power} \\end{figure}   From (\\ref{eq:pow1}), (\\ref{eq:pow_spec}) \\& (\\ref{eq:transfer}) we find that on small scales, the contribution to the {\\em rms} density fluctuation from a given logarithmic interval in $k$ is \\beq \\left (\\frac{\\delta\\rho}{\\rho}\\right )^2_k \\sim k^3 P_f(k) \\propto \\log^2{k}~, \\eeq which illustrates the fact that, although the smallest scales are the first to go non-linear, there is significant power to drive gravitational instability rapidly to larger scales in this model. Indeed, detailed N-body simulations of large scale structure show that filaments defining the cosmic web first form on the smallest scales. The scale-length characterizing the cosmic web grows as the universe expands, until at the present epoch the cosmic web consists of a fully developed supercluster-void network with a scale-length of several tens of Megaparsec \\cite{sss96,sss03,masaar,rien}.  Promising candidates for cold dark matter include a $100 - 1000$ GeV particle called a neutralino. The neutralino is a weakly interacting massive particle (WIMP). As its name suggests it is neutral and is a fermionic partner to the gauge and Higgs bosons (usually called the `bino, wino and higgsino'). It is believed that the lightest supersymmetric particle will be stable due to R-parity which makes the neutralino an excellent candidate for cold dark matter (see \\cite{roszko99,jungman96} for reviews of particle dark matter). A radically different particle candidate for cold dark matter is an ultra-light pseudo-Goldstone boson called an axion with a mass of only $m_a \\sim 10^{-5 \\pm 1}$ eV. Although ultralight, the axion is `cold' because it was created as a zero-momentum condensate. Its existence is a by-product of an attempt to resolve QCD of what is commonly called the `strong CP problem' which arises because non-perturbative effects in QCD give rise to an electric dipole moment for the neutron -- in marked contrast with observations \\cite{kolbt90}. Other candidates for non-baryonic cold dark matter include string theory motivated modulii fields \\cite{brustein}; non-thermally produced super-heavy particles having a mass $\\sim 10^{14}$ GeV and dubbed Wimpzillas \\cite{kolb}; as well as axino's and gravitino's -- superpartners of the axion and graviton respectively \\cite{roszko99}.  Since WIMP's cluster gravitationally, one should expect to find a flux of these particles in our own solar system and attempts are being made to determine dark matter particles by measuring the scattering of WIMP's on target nucleii through nuclear recoils. Now the earth orbits the sun with a velocity $\\simeq 30$ km/sec, even as the sun orbits the galaxy with $v_{M_\\odot} \\simeq 220$ km/sec. Furthermore the plane of the Earth's orbit is inclined at an angle of $60^\\circ$ to the glactic plane, because of which the dark matter flux on Earth is expected to be larger in June (when the Earth's velocity and the Sun's velocity add together) than in December (when these two velocities subtract). The resulting rate variation is about $7\\%$ between the flux measured during summer and winter. Precisely such a signal was reported by the DAMA experiment whose data (collected since 1996) appears to show a yearly modulation with greater events reported in June than in December \\cite{dama}. However results obtained by the DAMA group remain controversial since they have not been substantiated by other groups which report negative results for similar searches (see \\cite{munoz,khalil} for recent reviews on this subject).  \\begin{figure}[ht] \\centering \\includegraphics[width=5cm]{sol.eps} \\caption{\\footnotesize The Earth's motion around the Sun; from Khalil and Munoz (2001). } \\bigskip \\medskip \\label{fig:sol} \\end{figure}   Despite the excellent agreement of LCDM with large scale observations, some concerns have recently been expressed about the ability of this model to account for a number of smaller scale observations which can be summarized as follows:  \\begin{itemize}  \\item {\\em The substructure problem} is used to describe the fact that the cold dark matter model (with or without a cosmological constant) predicts an excessive number of dark matter subhaloes (or substructure) within a larger halo. If one (perhaps naively) associates each halo with a gravitationally bound baryonic object then the predicted number of dwarf-galaxy satellites within the local group exceeds the observed number by over an order of magnitude. Indeed, detailed N-body simulations as well as theoretical estimates predict around 1000 dark matter satellites in our local group which is much larger than the 40 or so observed at present \\cite{klypin99,moore99,kauffmann93,bullock00, somerville02,benson02,mateo98,tasi02,evans01}.  \\item {\\em The cuspy core problem}  CDM predicts a {\\em universal density profile} for dark matter halos in the wide range $10^7 M_\\odot - 10^{15} M_\\odot$ which applies both to galaxy clusters as well as individual galaxies including dwarfs and LSB's. \\footnote{Low Surface Brightness Galaxies (LSB's) are dominated by their dark matter content and therefore provide particularly good astrophysical objects with which to test dark matter models.} The density profile originally suggested by Navarro, Frenk and White \\cite{nfw} is \\beq \\rho(r) = \\rho_0\\left (r_s/r\\right )\\bigg\\lbrack 1 + \\left (\\frac{r}{r_s}\\right )\\bigg \\rbrack^{-2}~, \\eeq which gives $\\rho \\propto r^{-1}$ for $r \\ll r_s$ and $\\rho \\propto r^{-3}$ for $r \\gg r_s$, where $r_s$ is the scale radius and $\\rho_0$ is the characteristic halo density. (Other groups using higher resolution computations found somewhat steeper density profiles at small radii, such as $\\rho \\propto r^{-1.5}$ \\cite{moore99b,jing00}.)  {\\em The cuspy core problem} refers to the apparent contradiction between N-body experiments -- which show that the density profile in CDM halos has a $1/r$ (or steeper) density cusp at the center, and observations -- which appear to favour significantly shallower density cores in galaxy clusters as well as in individual dwarf and LSB galaxies (see \\cite{flores94,burkert95,bolatto02,blok97,primack01,tasi02,klypin04,rhee03,sand03,ma04} for detailed discussions of this issue).  \\end{itemize}  Although disconcerting, given the very considerable success of LCDM in explaining gravitational clustering on large scales, it may at this point be premature to condemn this model on the basis of small scale observations alone. It could be that the difficulties alluded to above are a result of an oversimplification of the complex physical processes involved and that a more careful analysis of the baryonic physics on small scales including the hydrodynamical effects of star formation and supernova feedback needs to be undertaken. For instance both dwarfs and LSB's have very shallow potential wells, a strong burst of star formation and supernova activity may therefore empty dark matter halos of their baryonic content resulting in a large number of `failed galaxies' and providing a possible resolution to the `satellite catastrophe'. (The failed galaxies will act as gravitational lenses and should therefore be detectable through careful observations.) Other explanations include the effects of tidal stripping recently discussed in \\cite{klypin04}. Likewise issues involving beam smearing, the influence of bars and the interaction of baryons and dark matter in the central regions of galaxies and clusters could be intricately linked with the central cusp issue and must be better understood if one wishes to seriously test the CDM hypothesis on small scales.  In concluding this discussion on dark matter I would like to briefly mention Modified Newtonian Dynamics (MOND) which, in some circles, is regarded as an alternative to the dark matter hypothesis. As the name suggests, MOND is a modification of Newtonian physics which proposes to explain the flat rotation curves of galaxies without invoking any assumptions about dark matter. Briefly, MOND assumes that Newtons law of inertia ($F = ma$) is modified at sufficiently low accelerations ($a < a_0$) to \\beq {\\bf F} = m{\\bf a} \\mu(a/a_0)~, \\eeq where $\\mu(x) = x$ when $x \\ll 1$ and $\\mu(x) = 1$ when $x \\gg 1$ \\cite{milgrom,sanders02}. It is easy to see that this results in the modification of the conventional formula for gravitational acceleration ${\\bf F} = m{\\bf g_N}$, resulting in the following relation between the true acceleration and the Newtonian value: $a = \\sqrt{g_N a_0}$. For a body orbiting a point mass $M$, $g_N = GM/r^2$. Since the centripetal acceleration $a = v^2/r$ now equals the {\\em true} acceleration $a$, one gets \\beq v^4 = GMa_0~, \\eeq \\ie for sufficiently low values of the acceleration the rotation curve of an isolated body of mass $M$ does not depend upon the radial distance $r$ at which the velocity is measured, in other words not only does this theory predict flat rotation curves it also suggests that the individual halo associated with a galaxy is infinite in extent ! (This latter prediction may be a problem for MOND since recent galaxy-galaxy lensing results \\cite{hoekstra02} suggest that galaxy halo's may have a maximum extent of about 0.5 Mpc.) The value of $a_0$ needed to explain observations is $a_0 \\sim 10^{-8}$cm/s$^2$ which is of the same order as $cH_0$ ! This has led supporters of this hypothesis to conjecture that MOND may reflect ``the effect of cosmology on local particle dynamics'' \\cite{sanders02}. Although MOND gives results which are in good agreement with observations of individual galaxies, it is not clear whether it is as successful for explaining clusters for which strong gravitational lensing indicates a larger mass concentration at cluster centers than accounted for by MOND \\cite{sanders02,combes02}. Another difficulty with MOND is that it is problematic to embed this theory within a more comprehensive relativistic theory of gravity and hence, at present, it is not clear what predictions a MOND-type theory may make for gravitational lensing and other curved space-time effects. For some recent developments in this direction see \\cite{beken04}.   To summarise, current observations make a strong case for clustered, non-baryonic dark matter to account for as much as a third of the total matter density in the Universe $\\Omega_m \\simeq 1/3$. The remaining two-thirds is thought to reside in a relative smooth component having large negative pressure and called Dark Energy. ",
        "conclusions": "From the theoretical standpoint the single most important question to be asked of dark energy is  \\centerline{Is $w = -1$ ?}   \\smallskip \\n Rephrased in terms of the Statefinder diagnostic the question is:  \\smallskip  \\centerline{Is $\\atridot/a H^3 = 1$ ?}  \\smallskip \\n If future observations do answer this question in the affirmative\\footnote{\\ie if $w = -1$ is measured to satisfyingly high accuracy} then, in all likelyhood the cosmological constant is the vacuum energy, and one will need to review the cosmological constant problem again, in order to fathom why the formally infinite quantity $\\tik$ is in fact so very small.  \\begin{figure}[ht] \\centering \\includegraphics[width=9cm]{confcmbclust.eps} \\caption{\\footnotesize Target statistical uncertainty of the SNAP experiment is shown overlayed with current results from CMB and LSS observations. From Aldering \\cite{aldering}. } \\bigskip \\medskip \\label{fig:snap} \\end{figure}  If on the other hand, either $w \\neq -1$ or if the DE density is shown to be time dependent, then the cosmological constant problem may need to be decoupled from the DE conundrum and searches for evolving DE models which produce $\\rho_{\\rm DE} \\simeq 10^{-47}$GeV$^4$ without exacerbating `cosmic coincidence' will need to be examined deeply in the light of developments both in high energy physics and in gravitation theory (superstring/M-theory, extra dimensions etc.). In either case the key to determining the properties of DE to great precision clearly lies with ongoing and future astrophysical experiments and observations.  Since the original discovery of an accelerating universe \\cite{perl98a,perl98b,riess98} the Sn data base has grown considerably and data pertaining to $\\sim 200$ type Ia supernovae are avaliable in the literature \\cite{tonry03,knop03,barris03,riess04}. Although systematic effects such as luminosity evolution, dimming by intervening extragalactic material (alternatively brightening due to gravitational lensing) continue to be a cause of some concern -- recall that a luminosity evolution of $\\sim 25\\%$ over a lookback time of $\\sim 5$ Gyr is sufficient to nullify the DE hypothesis \\cite{riess99} -- it is reassuring that recent observations of CMB anisotropies and estimates of galaxy clustering in the 2dF and SDSS surveys, make a strong and independent case for dark energy \\cite{spergel03,tegmark03a,tegmark03b}. Indeed, a joint analysis of CMB data from WMAP + HST Key Project determination of $H_0$ imply $w_{\\rm DE} < -0.5$ at the $95\\%$ confidence level \\cite{spergel03}.  It is of paramount importance that Sn observations continue to be supplemented by other investigations which are sensitive to the geometry of space and can be used as independent tests of the DE hypothesis. The volume-redshift test, Sunyaev-Zeldovich surveys, the Alcock-Paczynski test, the angular size-redshift test and gravitational lensing have all been suggested as possible probes of dark energy, and will doubtless enrich the theory vs observations debate in the near future. In addition, the proposed SNAP satellite which aims to measure light curves of $\\sim 2000$ supernovae within a single year \\cite{snap}, should provide a big step forward in our understanding of type Ia supernovae and help determine the cosmological parameters to great precision, as shown in figure \\ref{fig:snap}."
    },
    "0403/astro-ph0403438_arXiv.txt": {
        "abstract": "Of all pulsars known Vela has been one of the most productive in terms in understanding pulsars and their characteristics. We present the latest results derived from Australian telescopes. These include a more accurate pulsar distance, a more precise pulsar local space velocity, a new model of the spin up and the association of a radio nebula with the X-ray pulsar wind nebula. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403112_arXiv.txt": {
        "abstract": "We present \\hst\\ observations that show a bifurcation of colors in the middle main  sequence of the globular cluster   $\\omega$ Centauri.  We see  this in three different  fields,  observed with different cameras and filters.  We also  present high precision  photometry of a central ACS field which shows a  number of main-sequence turnoffs and subgiant branches.  The double  main  sequence,   the  multiple turnoffs    and subgiant  branches, and other population   sequences discovered in the past along    the red  giant branch   of   this cluster  add   up to a fascinating but frustrating puzzle.   We suggest various explanations, none of them very conclusive. ",
        "introduction": "A number of  properties (total mass, chemical composition, kinematics, and spatial   distribution of  the  stars)  make $\\omega$  Centauri  a peculiar object among Galactic   globular clusters.  The  most evident anomaly is the large spread in metallicity  seen both in spectroscopic (Norris \\&  Da Costa 1995)  and photometric (Hilker \\&  Richtler 2000, Lee et al.\\ 1999, Pancino et al.\\ 2000) investigations.  Most of  the  fascinating results   on \\omcen~  come  from the evolved stellar  population, which can be studied   in detail from the ground. In this paper we use \\hst~ data to explore  the cluster's turnoff (TO) and main-sequence (MS) populations.  While studies  of the evolved red giant  branch    (RGB)  can   explore   metallicity,   kinematic,  and spatial-distribution issues, we  need the fainter stars  if we hope to learn anything about ages  and mass functions, and  to give us  better statistics (there are $\\sim$10 MS stars for every RGB star).  The present  paper was stimulated by preliminary  results by one of us (Anderson 1997, 2002, 2003), on the presence  of multiple turnoffs and of a bifurcated main sequence (MS).  Here we  confirm that the unusual features    found  in   the  color-magnitude  diagrams  are  not  some data-reduction artifact, or a local phenomenon.  The features are real and  are present throughout the  cluster.  Unfortunately, the striking results  we present  here lead  to more  questions   than they answer. Though it will take  a lot of time  to fully exploit the  \\omcen~ data stored in the \\hst~ archive  (and we  are  working on this), we  think that these new results are  worthy of immediate publication because of their importance to the ongoing debate on the nature of this object. ",
        "conclusions": "One way  to interpret the observations  is  to assume  that either the photometric  calibration or the isochrones  are  in error.  If the net error is 0.06  mag in  $V-I$  color, then perhaps  the metal-poor (MP) population ([Fe/H]\\ $=\\!-1.6$) follows  along the bMS instead of along the  rMS.  If this is  the case then the  rMS would  correspond to the metal-rich  population ([Fe/H]\\ $=\\!-0.5$), as  the 0.06 mag $V\\!-\\!I$ separation  cannot be explained  by the metallicity difference between the  MP and Mint  populations.  (While it may  be conceivable that the isochrones  could have  errors in an   absolute sense, they  should be reliable in a differential sense.)  There are additional problems with this  interpretation.  First is that  only 5$\\%$ of  the RGB stars are metal rich, but in this scenario over 70$\\%$  of the MS stars would be metal  rich.  This would  imply  drastically different mass functions, such as have never  been seen before  anywhere (see Piotto  \\& Zoccali 1999).    Furthermore,  there is  no   actual   gap in  the   observed metallicity distribution and in the color  distribution of the RGBs of the MP and  Mint populations.  Most importantly, the  fact that the MS extension of the LTO runs parallel to the rMS (on the  red side of it, panel e) makes this scenario impossible.  The second interpretation is that the rMS corresponds to the MP stars, but the bMS corresponds to  a super-metal-poor population, with [Fe/H] $\\ll\\!-2$.  However, such a  large population of metal-poor stars  has never been observed in \\omcen\\ or in any other globular.  The third  possibility is  that  the populations of   the two MSs have sensibly different helium  content ($Y$).  Norris, Freeman, \\& Mighell (1996)  have shown that  the metallicity distribution of \\omcen\\ stars can be well  fitted by two separate  components, and argued  that this can be explained by two successive epochs of star formation.  Assuming for the more metal-rich  ([Fe/H]\\ $=\\!-1.0$) Mint population  a helium content of  $Y\\sim0.30$, we  find that the  corresponding  MS would be $\\sim0.07$ magnitude bluer in ($V-I$) than the  MP MS (assumed to have a canonical $Y=0.23$, and [Fe/H]\\ $=\\!-1.6$). Note that Norris et al.\\ (1996) found that the  ratio of the Mint  to  MP population should  be 0.2, compatible, within the uncertainties, with the  value we find for the rMS/bMS ratio.  Panel e  of Fig.\\ 1 shows that  the bMS could well be  connected with the intermediate TO-SGB.  Panel  a of Fig.\\ 2 shows that this  intermediate SGB is slightly  brighter  than the luminosity expected for a metallicity similar to the Mint, and the expected TO is redder than the observed one. These observational facts are consistent with this  population beeing helium enhanced  and slightly younger, as expected  if  the  helium  enhancement  is  due to self-pollution from intermediate AGB MP   stars.    The dramatic increase    of  s-process heavy-element   abundances  with metallicity  found   by Smith et al.\\ (2000) in \\omcen\\ RGB stars furtherly support the hypothesis that Mint stars  could have formed from  material  polluted by ejecta from 1.5-3 $m_\\odot$ AGB stars.  The  presence of a  population with  high helium content could also  account  for the  anomalously  hot HB  of  \\omcen, following  the calculations   of  D'Antona et al.\\  (2002).   All this notwithstanding, a $Y\\geq0.30$   is higher   than  any value   so  far measured   in Galactic GCs  (Salaris  et al.\\ 2004),   and not easy to understand.  As a fourth possibility, if we assume that  the rMS corresponds to the majority    of the cluster   stars,   the bMS  could  correspond  to a population of  stars located behind \\omcen.   As  shown in panel  d of Fig.\\ 2, if the   bMS is populated by   stars located 1.6  kpc  beyond \\omcen, we can easily fit it with an [Fe/H]  $=\\!-1$ isochrone.  Panel e of Fig.\\ 1 appears to strengthen this hypothesis: we see the bMS get closer and  closer to the  rMS, crossing it at  $H\\alpha\\sim18.5$, and apparently continuing into a broadened TO and SGB.  This broadening of the  intermediate   TO could  be   the  result of  a   spread  in both metallicity and  distance.  The overall  appearance of the CMD is that there are two sequences, shifted   by up to $\\sim0.3$--0.5  magnitude. The  hypothesis of a background  agglomerate of stars with metallicity around  [Fe/H] $\\sim\\!-1.0$ would also naturally  explain  why the bMS appears to intersect the rMS at  $V_{606}\\sim23.5$ (cf.\\ Figs.\\ 1c and 1d).  Such a background object would naturally explain the observation that   the  giants of  different   metallicity appear to have somewhat different spatial  distributions  (Jurcsik 1998,  Hilker  \\&  Richtler 2000), though this  spatial variation could  be explained by merger or self-enrichment scenarios as well.  Leon, Meylan, \\&  Combes  (2000) have  identified a tidal  tail around \\omcen.  Tidal tails often have a  clumpy nature.  However, the number of stars in the bMS  seems to be too  large and the sequence too sharp to be interpreted as a  part  of a clump   in a tidal tail behind  the cluster.  Another possibility is that the object  in the background is a distinct cluster  or a dwarf  galaxy.  As  it should cover  at least 20--30 arcmin in  the sky (this  is the extent of  the region where we identified a DMS) and   be located at about 7   kpc from the  Sun, the object should be extended  by at least  40--60 pc.  The probability of observing such an object in the direction of \\omcen~ is extremely low. However,  if   this object  happens  to be  gravitationally  linked to \\omcen~(either  because it  was part of  the  same original  system or because it  is the remnant of some  merging event), that would enhance the probability of seeing it in the same direction as \\omcen.  We note that the idea of a population of  stars behind the cluster has been suggested before.  Ferraro et al.\\  (2002) measured a bulk motion for  the RGB-a stars   with respect to  the  other cluster  stars, and interpreted this as evidence that it  could be a background object, or a  merger product that has  not  yet phase-mixed.  However, Platais et al.\\    (2003) find  this   motion  spurious,  attributing   it   to a color/magnitude term  in    the proper  motions.  Moreover,   Anderson (2003), using very accurate WFPC2 proper motions, contradicts the bulk motions seen by Ferraro et  al.\\ (see his  Fig.\\ 1).  In any case, the background population  we consider  here could  not correspond to  the very  metal-rich population; our Fig.\\ 1e  makes it clear that the LTO and the bMS are not related to each other."
    },
    "0403/astro-ph0403397_arXiv.txt": {
        "abstract": "Using a sample of 2408 time--resolved spectra for 91 BATSE GRBs presented by Preece et al., we show that the relation between the isotropic--equivalent luminosity ($L_{\\rm{iso}}$) and the peak energy ($E^{'}_{\\rm{p}}$) of the $\\nu F_{\\nu}$ spectrum in the cosmological rest frame, $L_{\\rm{iso}}\\propto E_{\\rm{p}}^{'2}$, holds within these bursts, and also holds among these GRBs, assuming that the burst rate as a function of redshift is proportional to the star formation rate. The possible implications of this relation for the fireball models are discussed by defining a parameter $\\omega\\equiv (L_{\\rm{iso}}/10^{52} {\\rm{erg\\, s}^{-1}})^{0.5}/(E^{'}_{\\rm{p}}/200\\, {\\rm{ keV}})$. It is found that $\\omega$ is narrowly clustered in $0.1-1$. We constrain some parameters for both the internal shock and external shock models from the requirement of $\\omega\\sim 0.1-1$, assuming that these model parameters are uncorrelated. The distributions of the parameters suggest that if the prompt gamma--rays are produced from kinetic--energy--dominated internal shocks, they may be radiated from a region around $R\\sim 10^{12}-10^{13}$ cm (or Lorentz factor $\\sim 130-410$) with a combined internal shock parameter $\\zeta_{\\rm{i}}\\sim 0.1-1$ during the prompt gamma--ray phase, which are consistent with the standard internal shock model; if the prompt gamma--rays of these GRBs are radiated from magnetic--dissipation--dominated external shocks, the narrow cluster of $\\omega$ requires $\\sigma\\sim 1-470$, $\\Gamma\\sim 216-511$, $E\\sim10^{51}-10^{54}$ ergs, $n\\sim 0.5-470$ cm$^{-3}$, and $\\zeta_{\\rm{e}}\\sim 0.36-3.6$, where $\\sigma$ is the ratio of the cold-to-hot luminosity components, $\\Gamma$ the bulk Lorentz factor of the fireball, $E$ the total energy release in gamma--ray band, $n$ the medium number density, and $\\zeta_{\\rm{e}}$ a combined external shock parameter, which are also in a good agreement with the fittings to the afterglow data. These results indicate that both the kinetic--energy--dominated internal shock model and the magnetic--dissipation--dominated external shock model can well interpret the $L_{\\rm{iso}}\\propto E_{\\rm{p}}^{'2}$ relation and the value of $\\omega$. ",
        "introduction": "% Gamma--ray bursts (GRBs) are now believed to be produced by jets powered by central engines with a standard energy reservoir at cosmological distances (see series reviews by Fishman \\& Meegan 1995; Piran 1999; van Paradijs et al. 2000; Cheng \\& Lu 2001; M\\'{e}sz\\'{a}ros 2002; Zhang \\& M\\'{e}sz\\'{a}ros 2003).  The most impressive features of GRBs are the great diversities of their light curves and spectral behaviors, and extremely large luminosities. These spectra are well fitted by the Band function (Band et al. 1993). However, the radiation mechanism at work during the prompt phase remains poorly understood. Although the spectral behavior and the luminosity are dramatically different from burst to burst, the isotropic--equivalent luminosity, $L_{\\rm{iso}}$, (or isotropic--equivalent energy radiated by the source, $E_{\\rm{iso}}$), and $E^{'}_{\\rm{p}}$, the peak energies of $\\nu F_{\\nu}$ spectrum in the rest frame among GRBs, obey an empirical relation of $L_{\\rm{iso}}\\propto E_{\\rm{p}}^{'2}$ (Amati et al. 2002; Yonetoku et al. 2003; Sakamoto et al. 2004; Lamb et al. 2003a, b, c). This relation was revisited in standard synchrotron/inverse--Compton/synchro--Compton models (Zhang \\& M\\'{e}sz\\'{a}ros 2002). Recently, Sakamoto et al. (2004) and Lamb et al. (2003a, b, c) pointed out that HETE--2 observations not only confirm this correlation, but also extend it to the population of X--ray flashes, which are thought to be a low energy extension of typical GRBs (Heise et al. 2001, Kippen et al. 2003). Based on this relation, Atteia (2003) also constructed a simple redshift indicator for GRBs.  One may ask: whether or not this relation holds in any segment within a GRB? The answer remains unknown. If the answer is positive, combining the results mentioned above, one might suggest that this relation is a universal law during the prompt gamma--ray phase, and presents some constraints on fireball models. In this Letter, we investigate this issue. Using a sample of 2408 time--resolved spectra for 91 BATSE GRBs presented by Preece et al. (2000), we show that this relation holds within these bursts, and also holds among these GRBs, assuming that the burst rate as a function of redshift is proportional to the star formation rate. We suggest that both the kinetic--energy--dominated internal shock model and the magnetic--dissipation--dominated external shock model may well interpret this relation.  Throughout this work we adopt $H_0=65$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm{m}}=0.3$, and $\\Omega_\\Lambda=0.7$. ",
        "conclusions": "Using a sample of 2408 time--resolved spectra for 91 long, bright GRBs presented by Preece et al. (2000), we show that the $L_{\\rm{iso}}\\sim E_{\\rm{p}}^{'2}$ relation holds within these BATSE bursts, and this relation also holds among these GRBs by assuming that the burst rate as a function of redshift is proportional to the star formation rate.  We discuss possible implications of this relationship for the fireball models by defining a parameter $\\omega\\equiv (L_{\\rm{iso}}/10^{52} {\\rm{erg\\ s}^{-1}})^{0.5}/(E^{'}_{\\rm{p}}/200\\, {\\rm{keV}})$. It is found that $\\omega$ is not influenced by the Doppler--boosting effect, and it is determined by the gamma--ray emission region and shock parameters in the kinetic--energy--dominated internal shock model or determined by the parameters of both the shock and the environment in the magnetic--dissipation--dominated external shock model. We derive the distributions of some parameters for both the internal shock model and the external shock model from the requirement of $\\omega\\sim 0.1-1$. We suggest that if the prompt gamma--rays are produced from a kinetic--energy--dominated internal shock, they may be radiated from a region around $R\\sim 10^{12}-10^{13}$ cm (or Lorentz factor $\\sim 130-410$) with an internal shock parameter $\\zeta_{\\rm{i}}\\sim 0.1-1$, which is consistent with the standard internal shock model; if the prompt gamma--rays of these GRBs are radiated from magnetic--dissipation--dominated external shocks, the $\\omega\\sim 0.1-1$ requires $\\sigma \\sim 1-470$, $\\Gamma \\sim 216-511$, $\\zeta_{\\rm{e}}\\sim 0.36-3.6$, $E\\sim 10^{51}-10^{54}$ ergs, and $n\\sim 0.5-470$ cm$^{-3}$. Please note that the distributions for these model parameters for both the internal and external shock models are based on the assumption that they are uncorrelated. Although these parameters seem to be uncorrelated, we do not know if it is really the case. If these parameters are correlated during prompt gamma--ray phase, these distributions are not valid.  We would like to thank the referee, Dr. Don Lamb, for his valuable comments, which have enabled us to improve greatly the manuscript. This work is supported by the National Natural Science Foundation of China (grants 10233010 and 10221001), the National 973 Project (NKBRSF G19990754), the Natural Science Foundation of Yunnan (2001A0025Q), and the Research Foundation of Guangxi University."
    },
    "0403/astro-ph0403442_arXiv.txt": {
        "abstract": "The variability of the Fe \\Ka line near 6.5 keV seems to be reduced compared to the variability of the hard X-rays which presumably drive the line emission. This is observed both in active galactic nuclei and galactic black hole binaries. We point out that such reduced variability, as well as lack of coherence between the variations of the line and the continuum, are a natural prediction of a propagation model of variability in the geometry of inner hot accretion flow. We compute detail model predictions of the variability characteristics which could be compared with current and future data. We also point out that the model requires a gradual disappearance of the cold disc, rather than a sharp transition from the cold disc to a hot flow. ",
        "introduction": "The Fe \\Ka fluorescent/recombination line near 6.5 keV is an important diagnostic of accretion flows around compact objects (see Reynolds \\& Nowak 2003 for a recent review). It is the strongest hard X-ray ($E>1$ keV) line which can originate in the innermost regions of accretion flows ($\\le 100\\,\\Rg$). It is indeed observed in energy spectra from all kinds of accreting sources: black hole and neutron star X-ray binaries, cataclysmic variables, active galactic nuclei (AGN).  The line is produced when plasma is irradiated by hard X-ray ($E > 7$ keV) radiation. If the plasma is  Thomson thick the continuum spectral component formed as a result of the irradiation (``Compton reflection'') has a characteristic shape, peaking at 20--30 keV (Lightman \\& White 1988). The line and the Fe K-shell absorption edge are superposed on this continuum (George \\& Fabian 1991; Matt, Perola \\& Piro 1991). The properties of the line and edge depend on the ionization of the reflecting medium: with increasing ionization the line and edge shift towards higher energies, while their strength increase  (e.g.\\ \\.{Z}ycki \\& Czerny 1994 and references therein). The profile of the line may be modified by relativistic and kinematic effects, if the line originates in e.g.\\ a rotating accretion disc (Fabian et al.\\ 1989). The line and absorption edge are then broadened and smeared. Such broad features are commonly seen in Seyfert galaxies (e.g.\\ MCG-6-30-15, Tanaka et al.\\ 1995, Fabian et al.\\ 2002) and black hole binaries (BHB; e.g.\\ Cyg X-1, Done \\& \\.{Z}ycki 1999; GRS 1915+105, Martocchia et al.\\ 2002), providing a clear evidence for a relativistic accretion disk extending deep into the gravitational potential of the central black hole.  X-ray emission from accreting sources is highly variable in a broad range of time-scales. Most of the variability power is located in the range of Fourier frequency $f \\approx$(0.1--1)$M/(10\\,\\MSun)$ Hz, corresponding to a time-scale ($T=1/f$) of a few seconds for a $10\\,\\MSun$ stellar black hole and a few weeks for a $10^7\\,\\MSun$ AGN (BHB review in McClintock \\& Remillard 2003; Markowitz et al.\\ 2003b for AGN). Typical power density spectrum is roughly a power law with slope $\\alpha\\approx 0$ ($P(f)\\propto f^{\\alpha}$) at low $f$, steepening to $\\alpha\\approx -1$ at $f \\approx 0.1$ Hz and to $\\alpha\\approx -2$ at $f=1$--3 Hz, for the BHB in low/hard state. Typical root-mean-square (r.m.s.) variability is 20--30\\%. The variability is stochastic rather than caused by deterministic chaos type of process (Czerny \\& Lehto 1997).  The observed variability of the Fe \\Ka line and the entire reprocessed component is somewhat surprising: they generally show rather less variability than the high energy continuum which is presumably driving the line emission and Compton reflection. This is seen both in AGN and BHB. In AGN the \\Ka line seems either not to respond to continuum variations on time-scales of minutes to days (e.g.\\ Reynolds 2000; Done, Madejski \\& \\.{Z}ycki 2000; Chiang et al.\\ 2000), or the line variability appears to be uncorrelated with that of the continuum (Vaughan \\& Edelson 2001). In particular, studies of r.m.s.\\ variability amplitude as a function of energy demonstrate the reduced variability in a relatively model independent way (Inoue \\& Matsumoto 2001; Markowitz, Edelson \\& Vaughan 2003a). The short term variability of the reprocessed component is more difficult to measure because of insufficient statistics, but where such studies were possible, the results also suggested reduced variability (e.g.\\ Done et al.\\ 2000).  A similar effect is seen in BHB systems. The time resolved spectral analysis is more difficult for BHB than for AGN,  but Fourier spectroscopy (study of Fourier-frequency resolved spectra; Revnivtsev, Gilfanov \\& Churazov 1999, 2001) can be used to investigate the variability characteristics on short time scales. Indeed, the amplitude of the reprocessed component in high Fourier frequency spectra is smaller compared to that in low Fourier frequency spectra (Revnivtsev et al.\\ 1999, 2001). In other words, the reprocessed component responds to the variability of the primary continuum on long time scales ($T=1/f \\ge 1$ sec) but it does not do so on short time scales.  In AGN one reason for the lack of variability might be the contribution to the reprocessing by a distant matter, e.g.\\ the obscuring dusty torus. Indeed, a narrow Fe \\Ka line is observed in many Seyfert galaxies. However, it is the variability properties of the broad component (hence produced very close to the central black hole) that are so puzzling.  One suggested explanation invoked complex ionization effects in the illuminated surface of the disc (e.g.\\ Nayakshin, Kazanas \\& Kallman 2000). Formation of a hot ionized skin (where the Fe \\Ka line is not produced) with thickness proportional to the X--ray flux may lead to anti-correlation between the line flux and the continuum flux. This model was quantitatively tested by \\.{Z}ycki \\& R\\'{o}\\.{z}a\\'{n}ska (2001), who concluded that it does indeed predict certain decrease of amplitude of variability of the line, although it is not possible to obtain an absolutely constant line flux.  In this paper we demonstrate that the reduced variability of the Fe \\Ka line is a natural result from a propagation model of X-ray emission in the geometry of a hot inner accretion flow. This geometry is one of the possibilities for accretion flow in low/hard states of accreting black holes (e.g.\\ Di Salvo et al.\\ 2001; see Done 2002 for review). It is supported by X--ray spectral studies  (review in Poutanen 1999), in particular the correlation between spectral slope and amplitude of reflection (Zdziarski, Lubi\\'{n}ski \\& Smith 1999) and long time scale spectral evolution of black hole binaries (Zdziarski et al.\\ 2003). The physical mechanism of formation the two-phase plasma flow may be related to plasma evaporation/condensation (R\\'{o}\\.{z}a\\'{n}ska \\& Czerny 2000 and references therein). The idea of propagating X-ray emitting structures was put forward by Miyamoto et al.\\ (1988), based on early variability studies of Cyg X-1. This was further developed and tested by e.g.\\ Nowak et al.\\ (1999), Misra (2000), and formulated in a more general form by Kotov, Churazov \\& Gilfanov (2001). Recently, \\.{Z}ycki (2003) showed that the Fourier-frequency resolved spectra can be reproduced in the propagation model, while Uttley (2004) argued for this model based on the r.m.s.--flux relation in accreting pulsar SAX J1808.4-3658 (see also Uttley \\& M$^{\\rm c}$Hardy 2001). ",
        "conclusions": "\\label{sec:discuss}   We have computed the variability properties of the Fe \\Ka line from X--ray reprocessing in a propagation model of X--ray emission in accreting compact objects. The model combines results from various spectral and timing studies of accreting black holes. The former suggest the geometry of the standard optically thick accretion disc truncated at a radius larger than the radius of the last stable orbit (e.g.\\ Esin, McClintock \\& Narayan 1997; Gierli\\'{n}ski et al.\\ 1997; \\.{Z}ycki, Done \\& Smith 1998; Done \\& \\.{Z}ycki 1999; Zdziarski et al.\\ 1999, 2003). The latter postulate correlated flares (avalanches) and spectral evolution during flares, in order to explain power spectra and time lags (Poutanen \\& Fabian 1999; Kotov et al.\\ 2001). Combined spectral-timing studies suggest a connection between the geometry and timing properties (Revnivtsev et al.\\ 1999, 2001; \\.{Z}ycki 2002, 2003). Most of the observational results were obtained for black hole binaries, but, were possible to conduct, analysis of data for Seyfert galaxies confirm the general similarity between the two classes of objects (e.g.\\ spectral studies by Done et al.\\ 2000; Chiang et al.\\ 2000; Lubi\\'{n}ski \\& Zdziarski 2001; timing studies of Czerny et al.\\ 2001; Uttley \\& M$^{\\rm c}$Hardy 2001; Vaughan et al.\\ 2003; Markowitz et al.\\ 2003b; Papadakis 2004).  \\begin{figure} \\epsfxsize = 8 cm \\epsfbox[18 250 620 710]{rmse.ps} \\caption{ Dependence of r.m.s.\\ variability on energy. Clear dip at Fe \\Ka line energy is seen, from the weakly variable line. \\label{fig:rmse}} \\end{figure}  The variability properties of the Fe \\Ka line seem to be similar in both classes of objects in that the line appear to be less variable than the continuum which drives it, and, where the variability is detected, it does not seem to be clearly correlated with the continuum. This is contrary to simple(st) ideas, whereby the continuum and the line are produced in the same region and thus should be closely related. We note that the observational situation is far from clear, though. Time resolved spectral analysis would be the most direct method to determine the variability of the line, but the results may be model dependent (see detailed discussion in Zdziarski et al.\\ 2003). Decomposition of counts from a medium energy resolution instrument like RXTE/PCA into the line and continuum depends on the model assumed for both the line and the reflected continuum. In particular, the effects of relativistic smearing of the reflected continuum was not taken into account. We note though that according to Markowitz et al.\\ (2003a) the results on line variability in a sample of Seyfert galaxies were insensitive to assumptions about the amplitude of the reflected component: whether its relative amplitude was fixed, or allowed to vary in accord with the \\Ka line.  Our model does reproduce the reduced variability of the Fe \\Ka line, compared to the variability of its driving continuum. Line variations may also appear not exactly correlated with continuum variations, because of the time delay between the peaks of the line and continuum fluxes. Both effects are necessary consequences of the adopted geometry, of a truncated disc with inner hot flow. The same geometry can also explain the hard X--ray time lags (Kotov et al.\\ 2001; \\.{Z}ycki 2003) through spectral evolution during flares (Poutanen \\& Fabian 1999). Quantitatively, our assumed ratio of heating to cooling rates, $C(r) \\propto \\lh(r)/\\lsoft(r)$, is a much weaker function of radius than could be expected for a compact (size $\\ll r$) active region and a sharply truncated disc. As already mentioned in \\.{Z}ycki (2003), in that latter geometry the supply of soft photons from the disc would diminish so rapidly that the predicted energy spectrum would be much too hard to be consistent with the data. This implies a {\\em gradual\\/} disappearance of the cold disc, which may be an interesting clue as to how the physical process of disc evaporation proceeds (R\\'{o}\\.{z}a\\'{n}ska \\& Czerny 2000).  We emphasize that the reduction of variability of the \\Ka line discussed in the present paper is the same phenomenon as the decreasing reflection amplitude with Fourier frequency, found in X-ray data of Cyg X-1 and GX 339-4 by Revnivtsev et al.\\ (1999, 2001) and modelled by \\.{Z}ycki (2003). In the present paper this effect was analysed with tools usually applied to AGN data, which are usually analysed in time domain rather than Fourier domain. We note, that it does not seem possible to design the parameters of the model, so that the line flux remains exactly constant. This is because, even though the line may be constant during each flare (assuming $C(r) \\propto 1/\\lh(r)$), the number of flares active at any time varies, and this causes variation of the total flux of the line.  A robust feature of the presented model is a connection of the time scale of the line {\\em response\\/} to the time scale of continuum variability. This is because the line flux responds on the time scale related to the duration of a flare. This in turn has to be chosen such that the observed power spectra are reproduced. The peaks in $f\\times P(f)$ are at rather long time-scales, much longer than just the light travel time through the region of most efficient energy generation. For example, the longest flares in our computations last $\\approx 10^6$ sec, which in light travel time corresponds to large distance of $2\\approx \\times 10^4\\,\\Rg$ (for $10^7\\,\\MSun$). This is simply a manifestation of the well known observational fact that time-scales of maximum X-ray variability power are much longer than the naively expected short dynamical time-scale. The physical mechanisms of producing the relatively long time scales of variability are rather unclear, but some interesting possibilities were recently considered in literature. 3-D magnetohydrodynamical simulations of Narayan, Igumenshchev \\& Abramowicz (2003) reveal slow ($v\\ll v_{\\rm ff}$) drift of plasma clumps across lines of magnetic field, which lines are compressed by the accretion flow. King et al.\\ (2004) considered a model involving magnetic dynamos operating locally in the disc. The longer and larger flares are results of correlations between dynamos acting in neighboring locations (radii). While individual dynamos operate of short (dynamical) time scales producing short flares, the correlated behaviour produces longer lasting, large events. Whatever the exact physical processes are, it is clear that the accretion flow is highly inhomogeneous and structured, factors that any realistic modelling should allow for.  In the presented model, the line flux is leading the continuum flux. We ignored the light travel time delay of the line photons after the continuum, but this is unlikely to affect our result, since the line is supposed to originate close to the location of emission of primary radiation. The light travel time delay may be estimated as $\\delta t \\sim r\\times \\Rg/c$ (assuming the height of the emission region $h \\sim r$), which gives $\\delta t\\sim 10^3$ sec, i.e.\\ about the time bin in our simulations. Obviously, there may be additional effects due to, for example, adjustment of properties of the reprocessing medium to the increasing irradiation flux, which might affect the result to some extent.  The dependence of r.m.s.\\ variability amplitude on energy is a relatively model independent demonstration of the reduced variability of the line. Our computations qualitatively reproduce the minimum of r.m.s.$(E)$ at the energy of the line. Generally, the r.m.s.\\ spectra are energy dependent variability amplitude, or, equivalently, can be thought of as representing energy spectra of the variable component of the spectrum. This can in principle be computed also in narrow ranges in Fourier frequency (Fourier frequency resolved spectroscopy, Revnivtsev et al.\\ 1999, 2001; \\.{Z}ycki 2002, 2003), but the quality of AGN data is not sufficient to make use of this technique as yet.  \\begin{figure*} \\parbox{\\textwidth}{ \\parbox{0.45\\textwidth}{ \\epsfxsize = 0.4\\textwidth \\epsfbox[18 160 620 710]{flli4.ps} }\\hfil { \\parbox{0.45\\textwidth}{ \\epsfxsize = 0.4\\textwidth \\epsfbox[18 160 620 710]{fllifl4.ps} } }} \\caption{ Line flux vs.\\ the 7--30 keV flux. {\\it Left panel:\\/} each point shows one time bin of 5000 sec.\\ from a $10^7$ sec.\\ observation. The line flux follows linearly the continuum flux for low values of the latter, corresponding to the emission coming from  outer regions (active region above the cold disc). The higher the continuum flux, the stronger the deviation from a linear relation, and the bigger the spread of points. {\\it Right panel:\\/} results from spectra binned in flux. The slope of the best fit line is $\\approx 0.56$. \\label{fig:flli}} \\end{figure*}  Concluding, the propagation model of high energy emission in the geometry of a truncated accretion disc provides a framework for understanding many of the observed spectral and temporal characteristics of X-ray radiation from accreting black holes."
    },
    "0403/astro-ph0403168_arXiv.txt": {
        "abstract": "We discuss the prompt emission of GRBs, allowing for $\\gamma\\gamma$ pair production and synchrotron self-absorption. The observed hard spectra suggest heavy pair-loading in GRBs. The re-emission of the generated pairs results in the energy transmission from high-energy gamma-rays to long-wavelength radiation. Due to strong self-absorption, the synchrotron radiation by pairs is in optically thick regime, showing a thermal-like spectral bump in the extreme-ultraviolet/soft X-ray band, other than the peak from the main burst. Recently, the prompt soft X-ray emission of GRB 031203 was detected thanks to the discovery of a delayed dust echo, and it seems to be consistent with the model prediction of a double-peak structure. The confirmation of the thermal-like feature and the double-peak structure by observation would indicate that the dominant radiation mechanism in GRBs is synchrotron rather than inverse-Compton radiation. ",
        "introduction": "In the past few years, a standard model was well established in which the gamma-ray burst (GRB) afterglows result from the relativistic blast-waves sweeping up the ambient medium of GRBs (\\Mesz 2002). However, the prompt emission of GRBs is believed to be irrelevant to ambient medium, and its radiation mechanism is still poorly known so far.  The recent definite proof of GRB 030329 associated with a type Ib/c supernova confirmed, as long suspected, that GRBs, at less the long class, originate from explosions of massive stars in distant galaxies (Stanek \\etal 2003; Hjorth \\etal 2003). Since GRBs are events occurring on stars, the emission region may be compact, and the huge energy release will lead to the formation of $e^\\pm,\\gamma$ fireballs, exhibiting thermal-like spectra. But the GRB spectra are non-thermal and hard, with a significant fraction of the energy above the \\pair pair formation energy threshold. For a photon with tens of MeV to escape freely, avoiding $\\gamma\\gamma$ interactions, the fireball must be ultra-relativistic expanding, with $\\Gamma\\ga100$ (Lithwick \\& Sari 2001, and references therein). The afterglow studies has also confirmed the presence of ultra-relativistic motion. However, if the intrinsic emission, before leaking out from fireball, includes radiation with even higher energy, say, beyond GeV, these radiation still suffers $\\gamma\\gamma$ absorption, leading to pair loading in GRBs. In context of relativistic fireball model, Li \\etal (2003; Li03 hereafter) found that, in a wide range of model parameters, the resulting pairs may dominate those electrons associated with fireball baryons. The presence of abundant pairs would affect the behaviors of the early afterglow from reverse shocks (Li03), and may also emit particular signals in the bursting phase.  We discuss in this Letter the prompt GRB emission, with emphasis on the re-emission by the secondary \\pair pairs. If the energy density in the emission region is dominated by magnetic field, the pairs would re-emit mainly by synchrotron radiation, rather than IC process (e.g., Pilla \\& Loeb 1998). Due to strong self-absorption, the pair emission appears as a thermal-like bump in the GRB spectrum, similar to the feature discussed by Kobayashi, \\Mesz \\& Zhang (2004) in the context of reverse shock emission. (Fan \\& Wei 2004 have also studied the pair emission, but with less stress on the self-absorption effect.) We further show that the intense soft X-ray emission in GRB 031203, inferred by the delayed dust halo (Vaughan \\etal 2004 [V04]; Watson \\etal 2004 [W04]), can be accounted for by the spectral bump due to pair-loading. This is of significant interest, since this feature could give a diagnostics for the magnetic field in the fireball (Kobayashi, \\Mesz \\& Zhang 2004) and the dominant radiation mechanism in GRBs. ",
        "conclusions": "We have studied the prompt GRB emission, allowing for $\\gamma\\gamma$ pair production and synchrotron self-absorption. Inferred by the observed characteristics of GRB emission, the resulting pairs usually dominate the baryonic electrons. The pairs will give rise to further emission by synchrotron radiation if in the strong magnetic field, which is also responsible to the prompt sub-MeV emission. However, due to strong self-absorption the pair emission exhibits a thermal-like bump in the extreme UV/soft X-ray band, other than the peak in the hard X-ray band. Since the feature emerges for $Y<1$, its observation gives a diagnostics for the magnetic energy density in the fireball (Kobayashi, \\Mesz, \\& Zhang 2004). The recent observation of a dust halo around GRB 031203 infers a spectral peak of the prompt burst emission in the soft X-ray band, which seems to be consistent with the predicted double-peak structure.  Some primary hypotheses have been taken in our calculation. First, we assume that the emission region is transparent for Compton scattering, even though the secondary pairs increase significantly the total optical depth. For typical parameter values this assumption is protected. However, in some extreme cases with quite small $\\Gamma$ and $R$, the secondary pairs may form an optically thick screen again (Guetta, Spada \\& Waxman 2001; Kobayashi, Ryde \\& MacFadyen 2002), which degrades the gamma-rays and results in an X-ray flash (XRF; \\Mesz \\etal 2002). If so, our calculation using eq. (\\ref{Npm}) may underestimate the pair-loading in XRFs, which may need detailed works of numerical simulation (e.g., Pe'er \\& Waxman 2003). Secondly, we assume strong magnetic field, $Y<1$, in the emission region. If $Y>1$, the pairs lose most energy by IC scattering the GRB photons, and the IC photons are not self-absorbed again since beyond the optically thick regime, hence no effective energy exchange between pairs and photons is established and the bump disappears. Therefore, once UV/soft X-ray bumps are detected this will infer $Y<1$ and that it is synchrotron rather than IC radiation that gives rise to the sub-MeV emission of GRBs."
    },
    "0403/astro-ph0403218_arXiv.txt": {
        "abstract": "We report the discovery of a locus of stars in the SDSS $g-r$ vs. $u-g$ color-color diagram that connects the colors of white dwarfs and M dwarfs. While its contrast with respect to the main stellar locus is only $\\sim$1:2300, this previously unrecognized feature includes 863 stars from the SDSS Data Release 1. The position and shape of the feature are in good agreement with predictions of a simple binary star model that consists of a white dwarf and an M dwarf, with the components' luminosity ratio controlling the position along this binary system locus. SDSS DR1 spectra for 47 of these objects strongly support this model. The absolute magnitude--color distribution inferred for the white dwarf component is in good agreement with the models of Bergeron et al. (1995). ",
        "introduction": "Modern large-scale accurate photometric surveys offer an unprecedented view of stellar populations. Here we discuss a population of unresolved binary stars which account for fewer than 10$^{-3}$ of stars detected by the Sloan Digital Sky Survey (York et al. 2000). Despite this low occurance frequency, the sample presented here is sufficiently large ($\\sim$1000 stars) to characterize their broad-band optical properties.  \\subsection{ Sloan Digital Sky Survey}  The Sloan Digital Sky Survey (SDSS; Abazajian et al.~2003, and references therein) is revolutionizing stellar astronomy by providing homogeneous and deep ($r < 22.5$) photometry in five passbands ($u$, $g$, $r$, $i$, and $z$; Fukugita et al. 1996, Gunn et al. 1998, Hogg et al. 2001, Smith et al. 2002), accurate to 0.02 mag (Ivezi\\'{c} et al.~2003). Ultimately, up to 10,000 deg$^2$ of sky in the Northern Galactic Cap will be surveyed. The survey sky coverage will result in photometric measurements for over 100 million stars and a similar number of galaxies. Astrometric positions are accurate to better than 0.1 arcsec per coordinate (rms) for point sources with $r<20.5^m$ (Pier et al.~2003), and the morphological information from the images allows robust star-galaxy separation to $r \\sim$ 21.5$^m$ (Lupton et al.~2003).  Here we report the results of a color-based search for binary stars in the recent SDSS Data Release 1 (see www.sdss.org), which includes 53 million unique objects detected in 2099 deg$^2$ of sky.   \\subsection{     The Stellar Locus in the SDSS Photometric System     }  The effective temperature is the dominant parameter that determines the position of the majority of stars in optical color-color diagrams constructed with broad-band filters (Lenz et al. 1998, and references therein). The effective temperature range results in a well-defined stellar locus in color-color diagrams (for more details see  Finlator et al. 2000, and references therein). ",
        "conclusions": "The accurate multi-band SDSS photometry for a large number of stars allowed detection of a new feature in the broad-band optical color-color diagrams: a ``bridge'' of stars, well-separated from the main stellar locus, connects the positions of M dwarfs and white dwarfs. The bridge characteristics are consistent with a binary system than includes an M dwarf and a white dwarf, with the system's position on the bridge determined by the components' luminosity ratio. This conclusion is strongly supported by SDSS spectra for 47 such systems.  The distance to these systems can be estimated in a straightforward way because a  photometric parallax relation for M dwarfs can be applied to $i$ and $z$ band measurements, where the contribution from the white dwarf is negligible. With a known system distance, the white dwarf luminosity-color distributions can be determined and compared to models. We find that models by Bergeron et al. (1995) are in good agreement with the data.  This work analyzed only about a quarter of the data that will be obtained by the SDSS. Thus, the color selection method presented here will eventually yield $\\sim$4,000 unresolved M dwarf-- white dwarf binary systems.   \\vskip 0.3in {\\it Acknowledgements}  We thank Princeton University for a generous support of this research.  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/."
    },
    "0403/gr-qc0403069_arXiv.txt": {
        "abstract": " ",
        "introduction": "Thanks to the recent technological advance, there are many on-going projects of gravitational wave detection in the world; the Laser Interferometric Gravitational Wave Observatory (LIGO)~\\cite{LIGO}, VIRGO~\\cite{VIRGO}, GEO-600~\\cite{GEO}, ACIGA~\\cite{ACIGA}, TAMA300~\\cite{TAMA} and the Large-scale Cryogenic Gravitational wave Telescope (LCGT)~\\cite{LCGT} which are ground-based laser interferometers, and EXPLORER~\\cite{EXPLORER}, ALLEGRO~\\cite{ALLEGRO}, NIOBE~\\cite{NIOBE}, NAUTILUS~\\cite{NAUTILUS} and AURIGA~\\cite{AURIGA} which are bar detectors. Furthermore, there are some future space interferometer projects such as the Laser Interferometer Space Antenna (LISA)~\\cite{LISA} and the DECi hertz Interferometer Gravitational wave Observatory (DECIGO)~\\cite{DECIGO}. The detection of gravitational waves provides us with not only a direct experimental test of general relativity but also a new window to observe our Universe. To use them as a new tool of observation, it is necessary to derive theoretical waveforms. Once we know them, we may appeal to the matched filtering technique to extract source's information from gravitational wave signals. However, because the signals are expected to be very weak and the amount of data will be enormous for long-term continuous observations, it is essentially important to develop efficient data analysis methods.  For these ground-based as well as future space-based interferometers, coalescence of compact object binaries is the most important source of gravitational waves. The process of a binary coalescence can be divided into three distinct phases. During the initial {\\it inspiral} phase, the gravitational radiation reaction timescale is much longer than the orbital period. The gravitational waves from the inspiral phase carry the information of the masses, spins and so on, of the system. After the inspiral phase, the binary becomes dynamically unstable and starts to merge. This phase is called the {\\it merger} phase. The gravitational waves from the merger phase give us the information about fully general relativistic dynamics of the system. Finally, if a black hole is formed after the merger phase, the system enters the {\\it ringdown} phase where it gradually settles down to a stationary Kerr black hole. During this process, the black hole emits gravitational waves with frequencies and damping rates specific to its mass and spin. Thus, the gravitational waves in the ringdown phase carry the information of the mass and spin of the final black hole.  In this paper, we consider an effective search method for gravitational ringing of distorted spinning (Kerr) black holes. The ringdown waves are described by quasi-normal modes of a black hole.  The quasi-normal modes are complex frequency wave solutions of the perturbed Einstein equations with purely outgoing-wave boundary condition at infinity and ingoing-wave at horizon, with vanishing incoming-wave amplitude. A quasi-normal mode is characterized by the central frequency $f_c$, usually called the (quasi-)normal-mode frequency, and the quality factor $Q$ which is inversely proportional to the imaginary part of the complex frequency. Their properties were analyzed extensively by Leaver~\\cite{Leaver}, and it is known that the least-damped (fundamental) mode belongs to the $\\ell=m=2$ spin-$2$ spheroidal harmonic modes. we note that a ringdown signal decays exponentially. Therefore, unless the signal-to-noise ratio (SNR) is very large, the signal will be soon buried in the noise after a few oscillation cycles. So, it is essentially important to keep the loss of SNR as small as possible when we construct a search method.  Here we consider a black hole characterized by its mass $M$ and the dimensionless spin parameter $a=J/M^2$ where $J$ is the spin angular momentum. The parameter $a$ takes a value in the range $[0,1)$, with $a=0$ corresponding to a Schwarzschild black hole and $a=1$ to an extreme Kerr black hole. It is known that the imaginary part of the $\\ell=m=2$ least-damped mode is the smallest of all the quasi-normal modes, and results of black hole perturbation calculations as well as numerical relativity simulations strongly suggest that a ringdown wave is dominated by this $\\ell=m=2$ least-damped mode unless $a$ is extremely close to unity. Hence we focus on this single mode. Then a ringdown waveform is expressed as \\begin{eqnarray} h(f_c,\\,Q,\\,t_0,\\,\\phi_0;\\,t) = \\cases{ e^{ - \\frac {\\pi \\,f_c\\,(t-t_0)}{Q}}\\,\\cos(2\\,\\pi \\,f_c\\,(t-t_0)-\\phi_0) & for $t \\geq t_0$ \\,, \\cr 0 & for $t < t_0$ \\,, \\cr} \\label{eq:RDwave} \\end{eqnarray} where we set the amplitude to unity for simplicity, and $t_0$ and $\\phi_0$ are the initial time and phase of the ringdown wave, respectively. For the $\\ell=m=2$ least-damped mode, analytical fitting formulas for the central frequency $f_c$ and quality factor $Q$ were given by Echeverria~\\cite{ech} as \\begin{eqnarray} f_c &\\simeq& 32{\\rm kHz}\\,[1-0.63\\,(1-a)^{0.3}] \\left({M \\over M_{\\odot}}\\right)^{-1} \\,, \\label{eq:fcMa} \\\\ Q &\\simeq& 2.0\\,(1-a)^{-0.45} \\,. \\label{eq:QMa} \\end{eqnarray}  There exists rich literature on search methods for ringdown waves. Echeverria~\\cite{ech} investigated the problem of extracting the black hole parameters from gravitational wave data in the case when SNR is large. Finn~\\cite{finn} improved this situation by developing a maximum likelihood analysis method that can deal with any SNR. Flanagan and Hughes then considered the parameter extraction from the three stages of a binary coalescence, i.e., from inspiral, merger and ringdown phases, in their series of papers~\\cite{FH}. For the ringdown phase, they discussed the relation between the energy spectrum of the radiation and SNR. Creighton~\\cite{crei} reported the result of analyzing data of the Caltech 40m by matched filtering, and emphasized the importance of coincidence event searches to discriminate spurious events from real events. But the search was limited to a single ringdown wave template. Recently, Arnaud et al.~\\cite{arna} discussed a tiling method to cover the 2-dimensional template space $\\{f_c,\\,Q\\}$. In our previous paper~\\cite{NTTS}, we proposed a more efficient method for tiling the template space. There, however, we ignored the dependence of the metric of the template space $\\{f_c,\\,Q\\}$ on the initial phase $\\phi_0$. This induces some small but non-negligible decrease in the match for a signal with certain ranges of $\\phi_0$. In this paper, we remove this shortcoming by properly taking into account the initial phase dependence, and develop a similar but substantially improved template spacing which is much more reliable than the previous one.  Here we make a comment on an analysis using real interferometers' data. In this case, we have to deal with non-stationary, non-Gaussian noises, and Tsunesada and Kanda~\\cite{TsuneKanda} found that more fake events are observed than in the case of an inspiraling wave search, because the duration of a ringdown wave is typically much shorter than that of an inspiral wave and it can easily be affected by short bursts. It will be necessary develop a way to remove such fake events without losing real ringdown signals, but we leave this issue for future work.  The paper is organized as follows. In Sec.~\\ref{sec:TDA}, first, we introduce orthonormal template waveforms for ringdown waves. Second, assuming a white noise background, we consider matched filtering in the 4-dimensional template space $\\{t_0,\\,\\phi_0,\\,f_c,\\,Q\\}$. We show that we can effectively reduce the template space to 2-dimensions spanned by $\\{f_c,Q\\}$, but the metric of this reduced template space depends on $\\phi_0$. Then, by carefully taking account of the $\\phi_0$ dependence of the metric, we analytically develop an efficient and reliable tiling method for ringdown wave searches. In Sec.~\\ref{sec:TN}, by using a fitting curve for the TAMA noise spectrum during the Data Taking 8 (DT8) in 2003, we show that our template spacing developed for white noise is valid even in the case of colored noise. Finally, Sec.~\\ref{sec:Dis} is devoted to summary and discussion. In Appendix~\\ref{app:TM}, we recapitulate the tiling method we proposed in~\\cite{NTTS}. In Appendix~\\ref{app:PEE}, we summarize the parameter estimation errors for ringdown signals by using the Fisher information matrix. ",
        "conclusions": "\\label{sec:Dis}  The detection of ringdown waves is a direct confirmation of the existence of a black hole. The ringdown waves have damped sinusoidal waveforms that reflect the mass and spin of a black hole. In our previous paper~\\cite{NTTS}, we proposed an efficient method for tiling the templates for matched filtering in the 2-dimensional $\\{f_c,\\,Q\\}$ space. However, it relied on the template space metric for signals with the initial phase $\\phi_0=0$. In this paper, we took account of the initial phase dependence and developed a new, improved method to search for the ringdown waves.  Since the template metric depends on the initial phase, we first determined the inner envelope curve of all the contours of a fixed maximum distance $ds^2_{\\rm max}$ between signals and templates on the $(f_c,\\,Q)$-plane for all possible values of the initial phase. We then constructed an ellipse that can cover the region inside the inner envelope as much as possible. Finally, we applied our tiling method proposed in~\\cite{NTTS} to obtain an improved, reliable tiling of the template space.  Another change from our previous work is the difference in the definition of the match. In our previous work, we defined the match as the square of the SNR, $\\rho^2$, while here we used the SNR as the match. So, for the same fixed maximum distance, the number of templates needed to cover the template space turns out to be smaller than what we obtained previously, if we ignore the initial phase dependence. As a result, the actual number of templates to cover the template space did not increase much from what was suggested in~\\cite{NTTS}.  We also examined the validity of our tiling method in the case of a colored noise spectrum by a Monte Carlo simulation. As a model of realistic noise power spectrum, we used a fitting curve of the noise power spectrum of TAMA300 during DT8 in 2003. For the pre-assigned maximum allowable SNR loss of $2\\%$, we found that only a few out of 2500 signals had the SNR loss larger than $2\\%$. This means that our template spacing is effective even in the case of colored noise.  Since the real data includes the non-stationary, non-Gaussian noise, it is necessary to check the effectiveness of this method in a realistic situation. Using the real data of TAMA300, this new template spacing is now being tested by Tsunesada et al.~\\cite{Tsune}. They find a lot of fake events due to the non-stationary, non-Gaussian noise. Particularly, the detector has many noise sources that can produce fake ringdown wave signals. Apparently, we need to develop a method to remove these fake events without losing real gravitational ringdown wave signals. Perhaps, the best way is to perform a coincidence analysis if we have plural detectors. We plan to study methods of coincidence or coherent analyses by using several detectors in the future."
    },
    "0403/astro-ph0403504_arXiv.txt": {
        "abstract": "{We explore the prospects for simultaneous, broad-band, multiwavelength polarimetric observations of GRB afterglows. We focus on the role of cosmic dust in GRB host galaxies on the observed percentage polarization of afterglows in the optical/near-infrared bands as a function of redshift.  Our driving point is the afterglow of GRB 030329, for which we obtained polarimetric data in the $R$ band and $K$ band simultaneously $\\sim$1.5 days after the burst. We argue that polarimetric observations can be very sensitive to dust in a GRB host, because dust can  render the polarization of an afterglow wavelength-dependent. We discuss  the consequences for the interpretation of observational data and  emphasize the important role of very early  polarimetric follow-up  observations in all bands, when afterglows are still bright, to study the physical properties of dust and magnetic fields in high-$z$ galaxies. ",
        "introduction": "Various independent studies of GRB afterglows have shown that in the optical bands  the degree of linear polarization  is usually rather low, in the order of 1 to 2\\% (for references see, e.g., Bj\\\"ornsson 2003; Covino et al. 2002, 2003; Greiner et al. 2003a; Lazzati et al. 2003; Masetti et al. 2003). If this turns out to be typical of optical afterglows (of long bursts), then the contribution of the dust in GRB host galaxies  to the observed polarization properties of the afterglows may not be negligible. While dust in a GRB host galaxy will also manifest itself by changing the spectral energy distribution of the afterglow light away from an original power-law (e.g.,  Palazzi et al. 1998; Ramaprakash et al. 1998),  polarization provides another independent method to constrain the existence of a dust component. If the sources responsible for the long bursts are located in star-forming regions, cosmic dust in the GRB environment will imprint a signal on the polarization properties of the afterglow light in the optical/near-infrared (NIR) bands. One might then imagine a situation in which the polarization by intervening dust (which might be time-independent) can dominate over the intrinsic polarization of the afterglow light. Motivated by our simultaneous optical and NIR polarimetric observations of the afterglow of GRB 030329, we explore here how dust in the GRB hosts may affect the polarimetry of the afterglow.  Thereby, we concentrate on the wavelength dependence of the degree of linear polarization $P$. While the intrinsic polarization of the afterglow light  is expected to be due to synchrotron radiation and, thus, wavelength independent, any additional polarization by dust in the GRB host can introduce a wavelength dependence to the measured $P$. Having this in mind, we basically focus on the redshift effect, which modifies this potential wavelength dependence of the degree of polarization (Serkowski 1973; Serkowski et al. 1975).  Recently, Lazzati et al. (2003) have discussed the potential imprint of cosmic dust in GRB hosts on the interpretation of polarization data, with particular emphasis on the afterglow of GRB 021004.  Whereas these authors concentrated on the interpretation of the time evolution of the percentage polarization, $P(t)$, here we focus on the wavelength dependence of $P$ at a fixed time. The latter comes into play if dust dominates the polarimetric properties of the afterglow light. While the time evolution of $P$ is basically an indicator for the GRB outflow geometry (Ghisellini \\& Lazzati 1999; Sari 1999; Rossi et al. 2004; Lazzati et al. 2004; for the cannonball model see Dado et al. 2004), the wavelength dependence of $P$ is a signature for the presence of dust in the GRB host. Also, whereas a measure of $P(t)$ requires an intense monitoring of the rapidly fading afterglow over many days, the wavelength dependence of $P$ can be investigated at any time, in particular when the afterglow is still in its earliest evolutionary phase, i.e., when the afterglow is brightest. Rapid multiwavelength polarimetric observations of GRB afterglows could then allow us to correct for the influence of dust once the time evolution of $P$ is studied. ",
        "conclusions": "Encouraged by our successful simultaneous optical and NIR polarimetric observations of the afterglow of GRB 030329, we have explored here in which manner simultaneous multicolor polarimetric observations can reveal an imprint from cosmic dust in GRB host galaxies on the afterglow light.  We have argued that dust in GRB hosts can strongly affect the polarization properties of an afterglow, depending on its redshift and on the photometric band, in which the observations are performed. In particular, the wavelength dependence of polarization due to dust, which is in strong contrast to the intrinsic wavelength independence of the fireball light, could be used as an indicator for the presence of dust.  With the upcoming $Swift$ satellite mission GRB afterglows are expected to be localized within seconds after burst trigger, in principle allowing for first polarimetric observations within the first hour after a burst.  Since the amount of  maximum polarization, $P_{\\rm max}$, can rapidly increase with increasing visual extinction within the host galaxy (Schmidt-Kaler 1958; Serkowski et al. 1975), in principle large degrees of linear polarization could be detectable for the most  intrinsically extinguished afterglows. Conversely, if a wavelength dependent polarization is never found in an afterglow then either the  host dust properties must be substantially different from those of the Galactic ISM, or there is no dust at all, or the magnetic field in the host galaxy is weak and/or not well ordered.  The position angle of the polarization defines the orientation of the predominant magnetic field in the host galaxy along the line of sight.  This predominant field can either be the large scale magnetic field acting in the GRB host (in particular, if the host is seen edge-on), or the field characterizing the star-forming region in which the burster is placed. The detection of interstellar dust extinction in the host without a substantial polarization signal could hint to a strongly tangled magnetic field structure in the host galaxy, or even indicate the absence of a magnetic field  large enough ($\\sim 10^{-8}$~G) to align the dust grains in the host galaxy. Both effects would indicate very strong evolutionary effects in the galactic magnetic field structure as practically all well studied nearby galaxies are known to have ordered magnetic field on large ($\\sim$~kpc) scales (Kronberg 1994; Scarrott 1996; Wielebinski \\& Krause 1993). In the Milky Way the dust grain alignment along the arbitrary line of sights due to the Galactic magnetic field produces a polarization per magnitude of visual extinction in the range $P/A_{\\rm V}\\approx 0.6-3~\\%\\,{\\rm mag}^{-1}$ (Serkowski et al. 1995). Thus a polarization below this value might indicate special conditions in a GRB host galaxy, including the star-forming region in which the burster is placed."
    },
    "0403/hep-th0403075_arXiv.txt": {
        "abstract": "We derive the $4D$ low energy effective field theory for a closed string gas on a time dependent FRW background. We examine the solutions and find that although the Brandenberger-Vafa mechanism at late times no longer leads to radion stabilization, the radion rolls slowly enough that the scenario is still of interest.  In particular, we find a simple example of the string inspired dark matter recently proposed by Gubser and Peebles. ",
        "introduction": "String theory continues to be our leading candidate for a quantum theory of gravity.  However, it continues to be a challenge to find phenomenological predictions that can be verified by experiment.  One challenge that stymies the effort towards string phenomenology is our lack of understanding of the ground state of string theory.  This has come to be known as the cosmological moduli problem (see e.g. \\cite{dine} and references within). One approach to resolving this problem is to take the point of view that cosmological evolution should be responsible for determining the value of the moduli.  One often finds that, by adopting this view, the moduli relax to special locations in the stringy landscape, which are points of enhanced symmetry. A recent example of such a scenario was presented in \\cite{stanford}. There it was found that the moduli are trapped in orbits around points of enhanced symmetry due to the production of light string modes. Stability then sets in due to the Hubble damping resulting from cosmological evolution. Alternatively, one can also have so-called racetrack models where the moduli continue to roll and do not remain fixed, but could give interesting cosmological consequences \\cite{quevedo}.  In this latter approach an important issue is fine tuning or the cosmic coincidence problem, i.e. why did the moduli start rolling at such a particular time?  The Brandenberger-Vafa scenario (BV scenario) \\cite{bv}, or Brane Gas Cosmology (BGC) \\cite{loitering,branes,stable,isotropy,perturbations,columbia,columbia2,subodh} as it has come to be known, is an example of a cosmological model that incorporates both of these approaches to the moduli problem. In these scenarios one generally works in the background of ten dimensional dilaton gravity\\footnote{Although this has been extended to M-theoretical considerations in \\cite{columbia2}.}, with sources given by the string winding and momentum modes.  It is found that the scale of the extra dimensions (radion) is then stabilized at the self-dual radius, where many of the string modes become massless and the symmetry is enhanced \\cite{stable}\\footnote{We note that the analysis of \\cite{stable} was performed with the case of bosonic strings in mind. Stabilization at the self-dual radius leading to enhanced gauge symmetry is expected for heterotic strings, but this does not occur for Type II strings. We thank Steve Gubser for reminding us of this point.}. A crucial aspect of these findings was the running of the dilaton to weak coupling, which was driven by the winding and momentum modes of the string.  In addition, it has been shown that this model is stable to both anisotropies \\cite{isotropy} and inhomogeneities at the linear level \\cite{perturbations}.  In this paper we would like to extend the BV scenario to better understand its predictions for late-time cosmology.  In particular, we would like to see if the stabilization mechanism is still plausible in the $4D$ effective field theory resulting from dimensional reduction.  Here, one usually assumes that the dilaton is fixed, since otherwise this would lead to unacceptable observational consequences\\footnote{However there have been attempts to keep the dilaton dynamical, see e.g. \\cite{polyakov}}. Given that at late times we are no longer in the regime of dilaton gravity, perhaps one would naively expect that such stabilization would no longer work.  That is, in General Relativity one must generally introduce exotic matter and/or violate the weak energy condition to stabilize extra dimensions \\cite{Carroll}. The string modes that we will consider here do not have such properties. To spare the reader suspense, we do in-fact find that stabilization fails, except in the special case of one extra dimension\\footnote{We mention that the case of $d=1$, from the perspective of the $5D$ Einstein frame, has already been considered in \\cite{subodh}.}. However, this is not as disastrous as one might first imagine. In fact, although the radion is no longer stable in the effective theory, its evolution is slow enough compared to that of the $4D$ background to be observationally acceptable.  In addition, we find that this evolution can lead to interesting phenomenology. As an example, we find an example of a cold dark matter candidate like that recently purposed by Gubser and Peebles \\cite{gubser}.  In Section II, we will briefly review the radion stabilization mechanism as presented in \\cite{stable}.  We present the stress energy tensor for the string modes and the corresponding action for the string modes. In Section III, we consider the evolution in the Einstein frame.  This is not the correct frame at early times when one is interested in the geodesics followed by the strings, but will be important for the late-time cosmology.  In Section IV, we dimensionally reduce the theory and find a form for the string modes in the $4D$ effective theory, which is given by their effective potential. From the effective potential we are able to discuss the stability of the radion, which is shown to exhibit slow rolling behavior.  This leads us to the possibility of closed strings in the extra dimensions behaving as dark matter. ",
        "conclusions": "In this article we have shown that stabilization of the radion in the effective theory at late-times is problematic.  This was not the case with one extra dimension, where after passing to the effective theory the potential retains a minimum at the self-dual radius. In that case, we were able to construct a simple realization of scalar dark matter as discussed in \\cite{gubser}. However, in the case of more than one extra dimension the potential does not possess a local minimum. This led to extra dimensions that are slowly growing, but that remain small compared to the large dimensions.  This case is interesting in many respects, e.g. it could lead to {\\em large} extra dimensions without the need to invoke brane world scenarios. We also note that these results remain valid for both a fixed and evolving dilaton, where in the later case the string modes drive the dilaton towards the region of weak coupling. We find the interpretation of winding and momentum modes as CDM much more involved, but possible.  We leave such considerations for future work."
    },
    "0403/astro-ph0403075_arXiv.txt": {
        "abstract": "Several recent papers have studied lensing of the CMB by large-scale structures, which probes the projected matter distribution from $z=10^3$ to $z\\simeq 0$.  This interest is motivated in part by upcoming high resolution, high sensitivity CMB experiments, such as APEX/SZ, ACT, SPT or Planck, which should be sensitive to lensing.  In this paper we examine the reconstruction of the large-scale dark matter distribution from lensed CMB temperature anisotropies.  We go beyond previous work in using numerical simulations to include higher order, non-Gaussian effects and find that the convergence and its power spectrum are biased, with the bias increasing with the angular resolution. We also study the contamination by the kinetic Sunyaev-Zel'dovich signal, which is spectrally indistinguishable from lensed CMB anisotropies, and find that it leads to an overestimate of the convergence. We finish by estimating the sensitivity of the previously cited experiments and find that all of them could detect the lensing effect, but would be biased at around the 10\\% level. ",
        "introduction": "\\label{sec:intro}  Weak gravitational lensing by large-scale structure has recently become a powerful tool in the cosmologists' toolbox, allowing us to map the mass distribution in the universe.  Lensing measurements using galaxy ellipticities have already begun to constrain cosmological parameters and to test our paradigm for hierarchical structure formation \\citep[see][for a recent review]{vWMe03}.  Even as this effort ramps up, yet another doorway into the weak lensing gold mine will open as a new breed of large surveys, such as the Atacama Cosmology Telescope (ACT\\footnote{http://www.hep.upenn.edu/$\\sim$angelica/act/act.html}), APEX-SZ\\footnote{http://bolo.berkeley.edu/apexsz/}, Planck\\footnote{http://astro.estec.esa.nl/Planck}, and the South Pole Telescope (SPT\\footnote{http://astro.uchicago.edu/spt/}), begin to come on line.  These surveys will probe the millimeter and sub-millimeter wavebands with unprecedented power and resolution, and thus enable us to map the large-scale distribution of mass in the universe by probing the gravity induced deflections of cosmic microwave background (CMB) photons as they journey from primordial times at $z\\sim 10^3$ to the present. In addition to providing the first ever map of the projected dark matter over most of cosmological history, such a measurement may have the potential to provide us with precision constraints on cosmological parameters, such as the neutrino mass and the dark energy equation of state \\citep{Kaplinghat03}, by measuring the matter power spectrum with percent level accuracy.\\\\\\\\ Beginning with the pioneering work of \\cite{Zaldarriaga99}, a considerable effort has been put into developing an accurate and sensitive estimator of the lensing effect.  In this paper, we seek to further the effort to develop methods necessary for the best possible estimation of the projected mass distribution from CMB temperature information.  To this end, we examine several effects which have not been studied in earlier works, both to see how they influence the reconstruction and whether they will provide us an opportunity to enhance the signal-to-noise of the estimator.  We note that although we will concentrate here on the extraction of information on large angular scales, reconstruction at the cluster level \\citep{Seljak00} is another exciting possibility, which we discuss elsewhere \\citep{VAW04}. The current state of the art techniques comprise maximum likelihood estimators \\citep{Hirata03} and the computationally more tractable quadratic estimators \\citep{Hu01a}.  These assume that both the primary CMB and the large scale structure responsible for the lensing are Gaussian random fields, and that noise is both Gaussian and uncorrelated with the signal.  The estimators are optimized to solve for the lensing effect by looking for the non-Gaussianity induced by the mapping of one Gaussian field by another.  Unfortunately, only the first of the four assumptions listed above is true.  The projected mass distribution is non-Gaussian except on large angular scales, and the kinetic Sunyaev-Zel'dovich (kSZ) effect \\citep[][for recent reviews see \\citealt{Reph95,Birk99,Carl02}]{SZ72,SZ80a} is a particularly pernicious source of confusion because it is non-Gaussian, spatially correlated with many of the structures doing the lensing, and spectrally indistinguishable from primary CMB anisotropies. To test the impact of these issues, we create lensing and kSZ fields by using an N-body simulation to model relevant structures, and then apply these fields to random realizations of the CMB (details of our methods are provided in an Appendix).  The maps that result from this are then used to reconstruct the projected mass density and the dark matter power spectrum for various experimental parameters.  We have elected to use the quadratic estimator of \\cite{Hu01a} for our reconstructions, in part because it is computationally more tractable than the maximum likelihood estimator, and because the two methods are predicted to be roughly equivalent for the angular scales and experimental parameters we are considering \\citep{Hirata03}.  More important, we hope to test the validity of the \\emph{assumptions} of the current generation of estimators, and thereby gain insight into the issues facing any method which is based on a Gaussian approximation.\\\\\\\\  We note that while our ability to remove foregrounds and point-sources from the CMB maps may ultimately prove challenging for any reconstruction, we will not focus on these complications here.  Instead, we will restrict our analysis to the observationally irreducible complications discussed above and instrument effects which are roughly consistent with our fiducial surveys.  Also, we will consider only the CMB \\emph{temperature} anisotropies, both because this simplifies the calculations and because it is more relevant for the next generation of wide field instruments with high angular resolution, which will not initially be polarization sensitive.  This subject remains of interest, however, and certainly merits continued study. \\\\\\\\ The outline of our paper is as follows.  In Section \\ref{sec:estimator} we describe the estimator of the projected mass distribution used hereafter, and motivate our choice.  We then use this estimator to reconstruct a Gaussian projected mass map in Section \\ref{sec:gauss}.  In Section \\ref{sec:nongauss} we investigate the effect of non-Gaussianity in the lensing field, and in Section \\ref{sec:ksz} we include contamination from the kSZ and describe our efforts to mitigate the problem.  The effects of these contaminants are considered as a function of instrument resolution in Section \\ref{sec:surveys}, where we also provide an estimate of how well our fiducial surveys will reconstruct convergence maps and power spectra.  We then summarize and discuss our results in Section \\ref{sec:summary}.  Finally, some details of the simulations and the estimator are presented in an Appendix. ",
        "conclusions": "\\label{sec:summary}  Lensing of CMB photons may provide a window into the matter distribution projected from primordial times to the present.  A great deal of work has been put into developing the statistical methods needed to fulfill this promise, and our goal here has been to further these efforts.  To this end, we have extensively checked the power and assumptions of one of the two principal estimators that have emerged, the ``optimal'' quadratic estimator of \\cite{Hu01a}, both for reconstructing projected mass maps and for measuring the matter power spectrum.  Although our results apply specifically to the quadratic estimator, we note that the other principal candidate \\citep[the maximum likelihood estimator of][]{Hirata03} makes the same assumptions about the underlying fields as the quadratic estimator, so it is likely to face the same challenges.  We began our efforts in Section \\ref{sec:gauss}, where we tested the quadratic estimator method under the artificial conditions of a Gaussian gravitational potential, and uncorrelated Gaussian noise, for which the estimator is most likely to succeed.  The noise in the reconstructed maps has $\\kappa$ dependent extra terms which constitute an additive bias in power spectrum measurements. The second order terms were discussed in \\citep{Cooray03}, and we confirm their prediction that these are numerically significant.  In addition, we find for the first time that higher order terms are also likely to be significant, since these second order terms do not fully correct bias for interesting observational parameters.  Although the bias is signal dependent, it may be possible to correct for this fact by using iterative methods or model fits, as suggested by \\cite{Kesden03};  however, high order terms would need to be included to obtain good accuracy.  It is clear that a reconstruction method based on the assumption of a Gaussian lensing field will not optimally treat non-Gaussian features, such as clusters, which contain more information than just the 2-point function.  We investigated this issue in Section \\ref{sec:nongauss} by using a lensing field generated from an N-body simulation.  This imposed both a multiplicative bias on the cross spectrum estimate, and an additional additive bias on the auto spectrum.  The bias increased as the signal-to-noise ratio increased.  We investigated the impact of a second non-Gaussian complication, the kinetic Sunyaev-Zel'dovich effect (kSZ), in Section \\ref{sec:ksz}. This ``noise'' is highly correlated with the lensing signal and, unlike the larger thermal SZ, is impossible to distinguish from the CMB using spectroscopic methods.  We began by showing that, if left untreated, this effect will completely dominate the reconstruction. One obvious method to correct for this is to use a measurement of the thermal SZ (which is correlated with the kSZ) to mask pixels which are likely to have a large contamination, and we were able to substantially reduce the bias caused by the kSZ while excising only a relatively small number of pixels. Although this was encouraging, not all of the kSZ signal is correlated with the thermal SZ and the contamination level remained significant. Further reduction of kSZ noise by simply masking more pixels suffers from steadily decreasing returns in signal-to-noise improvement. We suspect, although we have not checked it explicitly, that contamination {}from the Ostriker-Vishniac effect would behave in a similar manner.  Given the difficulty that one encounters using Gaussian reconstruction methods, it is natural to ask if the non-Gaussian nature of the lensing signal might itself provide the best probe of the lensing field.  This topic has recently been the subject of great interest in the context of galaxy lensing \\citep[e.g.][]{TaJa03,ZaSc03,HoWh03}, but to this point a similar effort is lacking for lensing of the CMB.  There is clearly promise in this idea; however, any measurement of the non-Gaussian signal is likely to be particularly sensitive to the masking technique needed to control the kSZ, which will be highly correlated with this signal. For example, clusters will produce both a large non-Gaussian lensing signal and a large kSZ in the same location, \\emph{at the cluster}. Thus, it may be necessary to mask precisely those pixels which would otherwise have provided the best window into the non-Gaussian signal.  In this paper, we have been primarily interested in the reconstructions that can be accomplished using the unprecedented power and resolution that will be available in the next generation of surveys, and it is in this context that biasing effects emerge as important obstacles.  In the regime of large smoothing and low signal-to-noise, the bias due to non-Gaussian effects is small, and the quadratic estimator performs well in this regime. It is only when the observational parameters are improved that the biasing effects we have been discussing emerge.  We show in Fig. \\ref{fig:biasres} the unfortunate increase in bias that occurs as signal-to-noise is improved, so that a recovery of the power spectrum can have either small error bars or negligible bias, but not both.  Thus, while a detection of the lensing effect is likely to be achieved by APEX/SZ, more ambitious goals that rely on extremely high fidelity reconstructions of the matter power spectrum require significantly more modeling.  We have presented a number of reconstructed maps in this paper, and it is clear that features in these maps do coincide with the those of the original fields at some level.  However, we note that the maps have been smoothed by a $20^\\prime$ FWHM Gaussian window, and even so the reconstructed maps are far from perfect.  We have also performed extensive numerical integrations to compute the corrections to the noise term $\\mathcal{N}_{\\ell}$ outlined in Section \\ref{sec:gauss}. Although these corrections proved useful for for removing some of the bias effects, the computational power needed for this calculation is non-trivial, a point which we discuss in more detail in the Appendix.  Although we have not treated the issue here, we would like to make some comments about the role of polarization in CMB lensing.  It has been emphasized \\citep{Guzik00,Hu02,Cooray03,Hirata03b,Kesden03,Okamoto03} that the inclusion of polarization information might dramatically enhance the prospects for large-scale structure reconstruction from lensing of the CMB.  This is because lensing induces a $B$-mode polarization signal which is otherwise absent for purely scalar, primary fluctuations.  The large intrinsic primary CMB anisotropies, which are a source of ``noise'' for lensing reconstruction, are thus absent.  However, the spatial structure is complicated for polarization, as it is for temperature, and the signal levels are much smaller.  The kSZ effect, which is one of our major contaminants, is also polarized \\citep{SZ80b}.  Thus, although it is certainly reasonable that the addition of polarization information would enhance the prospects for $C_\\ell^{\\kappa\\kappa}$ reconstruction, a detailed calculation is required to determine how much better one can actually do.  Weak lensing is one of the principal tools for cosmologists, and promises to be increasingly significant for many years to come.  Lensing of the CMB is a new addition to the arsenal, which will be facilitated by powerful upcoming surveys such as ACT, APEX-SZ, Planck, and SPT. Even as observers gear up to probe the millimeter and sub-millimeter wavebands with unprecedented power and resolution over large fractions of the sky, theorists continue to improve the methods necessary to extract relevant information from the resulting measurements.  If the pace of development on both fronts continues at its current rate, the future looks bright indeed.  \\noindent {\\bf Acknowledgments:}\\newline A.A. would like to thank the organizers of the workshop ``Cosmology with Sunyaev-Zel'dovich cluster surveys'' held in Chicago in September 2003 for allowing him the chance to present some of this work.  Additionally we would like to thank T. Chang, J. Cohn, M. Zaldarriaga for helpful discussions about these results. The simulations used here were performed on the IBM-SP2 at the National Energy Research Scientific Computing Center. This research was supported by the NSF and NASA.  \\appendix"
    },
    "0403/astro-ph0403596_arXiv.txt": {
        "abstract": "{ We present an \\xmm mosaic observation of the hot ($kT\\sim6.5$~keV) and nearby ($z=0.0881$) relaxed cluster of galaxies A478.  We derive precise gas density, gas temperature, gas mass and total mass profiles up to $12\\arcmin$ (about half of the virial radius  $R_{200}$).  The gas density profile is highly peaked towards the center and the surface brightness profile is well fitted by a sum of three $\\beta$--models. The derived gas density profile is  in excellent agreement, both in shape and in normalization,  with the published Chandra density profile (measured within $5\\arcmin$ of the center).  Projection and PSF effects on the temperature profile determination are thoroughly investigated. The derived radial temperature  structure is as expected for a cluster hosting a cooling  core, with a strong negative gradient at the cluster center.  The temperature rises from $\\sim2$~keV up to a plateau of $\\sim6.5$~keV beyond $2\\arcmin$ (i.e. $r>208\\rm{kpc}=0.1\\, R_{200}$, $R_{200}=2.08$~Mpc being the virial radius). From the temperature profile and the density profile and on the hypothesis of hydrostatic equilibrium, we derived the total mass profile of A478 down to 0.01 and up to 0.5 times the virial radius. We tested different dark matter models against the observed  mass profile.  The Navarro, Frenk \\& White (\\cite{navarro97}) model is significantly preferred to other models. It leads to a total mass of $M_{200}=1.1\\times 10^{15}$~M$_\\odot$ for a concentration parameter of $c=4.2\\pm0.4$.  The gas mass fraction  increases slightly with radius. The gas mass fraction  at a density contrast of $\\delta=2500$ is $\\fgas=0.13\\pm0.02$, consistent with previous results on similar hot and massive clusters.  We confirm the excess of absorption in the direction of A478. The derived absorbing column density   exceeds the 21~cm measurement by a factor of $\\sim2$, this excess extending well beyond the cool core region.  Through the study of this absorbing component and a cross correlation with infrared data, we argue that the absorption excess is of Galactic origin, rather  than intrinsic to the cluster. ",
        "introduction": "As nodes of large scale structure and thus places of dark matter concentration, galaxy clusters can be used as powerful tools to test theories of structure formation.  The basic hierarchical scenarios based on gravitation make the population of galaxy clusters a homologous population of sources.  Their physical properties follow scaling laws depending only on their redshift and mass, and their internal structures are similar.  The exceptional capabilities of \\xmm in terms of sensitivity and of Chandra in term of spatial resolution allow us to characterize the gas density and temperature profiles with  unprecedented accuracy.  For a relaxed cluster, the hydrostatic equations can be used to derive the underlying dark matter distribution, from the very central part of clusters up to nearly the virial radius (David \\etal~\\cite{david01}; Allen \\etal~\\cite{allen01b}; Arabadjis, Bautz \\& Garmire~\\cite{abg02}; Allen, Schmidt \\& Fabian~\\cite{allen02a}; Pratt \\& Arnaud~\\cite{pratt02},\\cite{pratt03}; Lewis \\etal~\\cite{lewis03}; Buote \\& Lewis~\\cite{buote03}). The observed clusters seem to have a cusped dark matter profile as predicted by numerical simulations (Navarro, Frenk \\& White~\\cite{navarro97}; hereafter NFW; Moore \\etal~\\cite{moore99}; hereafter MQGSL).  However, the central slope of the dark matter profile and the possible dispersion of the concentration parameter remain open issues. Larger samples of high quality mass profiles are needed to further assess these points.  In this paper we present the \\xmm spectro-imaging observation of A478, a massive, relaxed nearby cluster ($z=0.0881$ -- \\cite{struble99}). Detected in surveys (UHURU, HEAO-1, Ariel-V), this cluster is well known in X-rays and its physical properties have been carefully studied with previous X-ray observatories: EXOSAT (Edge \\& Stewart~\\cite{edge91}), Einstein and Ginga (Johnstone \\etal~\\cite{johnstone92}), ROSAT (Allen \\etal~\\cite{allen93}; White \\etal~\\cite{white94}) and ASCA (\\cite{markevitch98}, White \\etal \\cite{white00}).  All those previous studies converge for what concerns the overall temperature of the cluster, $kT \\sim 6.8$~keV. Recently, Sun \\etal (\\cite{sun03}) performed a high angular resolution study of the central part of the cluster with Chandra.  They pointed out the presence of an X-ray cavity in the very central part of the cluster which is anti-correlated with the radio lobes.  Here we focus on the characterization of the gas and dark matter distribution of A478.  In a companion paper de Plaa \\etal (in prep.) present a detailed spectroscopic study of the metal abundances and their distribution within A478's core based on EPIC and RGS data.  We present the observation and the different data processing steps in Sect.~\\ref{sec:data}.  In Sect.~\\ref{sec:morpho} we briefly discuss the cluster morphology.  In Sect.~\\ref{sec:ne} we analyze the surface brightness profile and derive the gas density profile. Spatially resolved spectroscopic analysis is presented in Sect.~\\ref{sec:spec}, where we also discuss the temperature and absorption profiles.  In Sect.~\\ref{sec:mass}, we present the resulting total mass and gas mass fraction profiles of A478 and we discuss the shape of the dark matter profile according to our observational results.  Throughout this paper,  we use $H_0=70$~km Mpc$^{-1}$ s$^{-1}$, $\\Omega_m=0.3$ and $\\Omega_{\\Lambda}=0.7$. In such a cosmological framework, at the cluster redshift ($z=0.0881$) $1$\\arcmin~$=99$~kpc. ",
        "conclusions": "In this paper we have analysed imaging and spatially resolved spectral data of the galaxy cluster A478 obtained with the \\xmm satellite.  We obtained well constrained absorption, gas density and gas temperature profiles up to $\\sim 0.5$ virial radius.  As in previous studies, we found an excess of absorption in the direction of A478.  The derived absorbing column density exceeds the 21~cm measurements by a factor of $\\sim2$ in the center and the excess extends well beyond the cooling core region.  This excess of absorption seen in A478 (and other cooling flow clusters) was interpreted in pre-XMM and Chandra studies (e.g. Allen \\etal \\cite{allen93,fabian94}) as the signature of intrinsic cool absorbing material, a consequence of the strong cooling flow in the cluster center.  From the absorption excess extent and a detailed comparison with FIR data, we argue that the absorption excess is rather of Galactic origin. We suggest it could be the effect of a Galactic molecular/cold cloud type structure in the line of sight.  The next generation of FIR space missions will help to clarify this issue with sensitive FIR mapping of the whole cluster area with a high spatial resolution.  We fitted the surface brightness profile with various parametric models of the gas density profile, taking into account  the \\xmm PSF. The gas density profile, derived on scales of  $0.03\\arcmin - 13\\arcmin$,   is highly peaked towards the center and is well fitted by a quadratic sum of three \\betamod.  The derived gas density profile is in excellent agreement both in shape and normalization with the Chandra density profile (measured up to  $5\\arcmin$ of the center). This indicates that the PSF modeling we have used is basically correct and that accurate density profiles in the very center of the cluster can be derived with XMM data, in spite of the PSF blurring.  A raw temperature profile was obtained on scales of  $0.07\\arcmin - 10\\arcmin$  by fitting isothermal models to spectra extracted in $13$ concentric annuli. This profile  shows  a sharp negative gradient measured toward the center ($r<2\\arcmin$), a  signature of a cooling core. Beyond that region the profile is essentially flat. We have thoroughly investigated projection and PSF effects on the temperature profile determination. The PSF effects beyond $0.3\\arcmin$ are much less important than projection effects, whereas both are important in the very center. We discuss the noise introduced by the correction of these effects and a way to overcome this problem.    The derived deprojected PSF-corrected temperature profile ranges  from  $\\sim2$~keV in the center up to an asymptotic value of $\\sim6.5$~keV.  Using this temperature profile and the density profile, we have derived the total mass profile for this cluster from $0.01$ up to $\\sim0.5$ times the virial radius. Systematic uncertainties due to the PSF and projection correction for the temperature profile are taken into account. We have tested different models for the dark matter profile distribution against the observed mass profile. A mass distribution with a cusp in the center, as predicted from numerical simulations, is clearly preferred.  An isothermal sphere model does not provide a good fit to the data. In a second time we tested an MQGSL model and an NFW model.  Those two types of models have similar shapes at large radii (they both scale like $r^{-1.5}$) but differ significantly at small radii. Therefore to discriminate those two models one needs data with a high statistic quality over a wide range of radii (i.e. covering at least two decades).  Our data set nearly fulfilled this requirement and we were indeed able to discriminate between the two models, the NFW model being preferred to the MQGSL model. For the NFW model, we derived a concentration parameter $c=4.2\\pm 0.4$. This value is as expected from numerical simulations: $c \\sim6$ (Navarro, Frenk \\& White \\cite{navarro97}; \\cite{eke98}) with a typical $1\\sigma$ dispersion of $\\Delta(\\log(c))=0.18$ (Bullock \\etal~\\cite{bullock} ). This work can be compared with the similar work on the cluster A1413 by Pratt \\& Arnaud (\\cite{pratt02}). In the case of A1413, if the NFW model was acceptable, the MQGSL model was slightly preferred. Although this cluster is detected out to 0.7 times the virial radius, the data are limited in the center, a shortcoming, as emphasized above, for discriminating between those two models. Moreover, data in the center only are not sufficient (see the work on A1983 by Pratt \\& Arnaud \\cite{pratt03}).  On the other hand, our result agrees with the analysis of A2029 by Lewis et al. (\\cite{lewis03}), which clearly favors an NFW dark matter profile. To our knowledge, this is the only other data set which covers a similar wide radial range (0.001 to 0.1 virial radius) .  The key factors in fitting the mass profiles with different dark matter models, are the resolution in the center as well as the data at large radii. To date \\xmm  is the best satellite to compute total mass profiles, especially through its capability to derive precise temperature profiles. Nevertheless, its spatial resolution limits the investigation at the very center of galaxy clusters. A direct combined analysis of \\xmm and Chandra data of very well relaxed clusters seems to be an ideal path to a full description of the dark matter profile in clusters. However,  one has to keep in mind that this requires an excellent cross calibration between the two satellites, so that the temperature profiles derived at various scales can be combined."
    },
    "0403/astro-ph0403243_arXiv.txt": {
        "abstract": "An empirical relation between the broad line region (BLR) size and optical continuum luminosity is often adopted to estimate the BLR size and then the black hole mass of AGNs. However,  optical luminosity may not be a good indicator of photoionizing luminosity  for extremely radio-loud AGNs because the jets usually contribute significantly to the optical continuum. Therefore, the black hole masses derived for blazar-type AGNs with this method are probably overestimated. Here we first derived a tight empirical relation between the BLR size and the H$_\\beta$ emission line luminosity, $R(\\rm{light-days})= 24.05(L_{H_\\beta}/10^{42} ergs~s^{-1})^{0.68}$, from a sample of 34 AGNs with the BLR size estimated with the reverberation mapping technique. Then we applied this relation to estimate the black hole masses of some AGNs and found that for many extremely radio-loud AGNs the black hole masses obtained with the $R-L_{H_\\beta}$ relation are systematically lower than those derived previously with the $R-L_{5100\\AA}$ relation, while for radio-quiet and slightly radio-loud AGNs the results obtained with these two methods are almost the same. The difference of black hole masses estimated with these two relations increases with the radio-loudness for extremely radio-loud AGNS, which is consistent with the fact that their equivalent widths of H$_{\\beta}$ emission line become smaller at higher radio-loudness. If the small H$_{\\beta}$ equivalent widths of extremely radio-loud AGNs are indeed caused by the beaming effect, we argue that the optical emission line luminosity may be a better tracer of ionizing luminosity for  blazar-type AGNs and the black hole mass derived with the $R-L_{H_\\beta}$ relation  are probably more accurate. ",
        "introduction": "Supermassive black hole is essential  for AGNs activities (Lynden-Bell 1969; Rees 1984). The black hole masses of some nearby AGNs have been recently estimated by the reverberation mapping technique (Wandel, Peterson \\& Malkan 1999; Ho 1999; Kaspi \\etal 2000), with which the size of the broad line region (BLR) can be measured from the time delay between the flux variations of the continuum and the emission lines of AGNs. The black hole mass is then estimated using the Virial theorem from the BLR size and the characteristic velocity (determined by the full width at half-maximum (FWHM) of emission line). So far, the reverberation studies have yielded the black hole masses of about 20 Seyfert 1 galaxies and 17 nearby bright quasars.  An empirical relation between the BLR size ($R$) and the optical continuum luminosity at 5100$\\AA$ ($L_{5100\\AA}$) has been derived by Kaspi et al. (2000) using the observed data of 34 nearby AGNs. Because the  measurement of the BLR size with the reverberation mapping technique needs long-term monitoring of continuum and emission line fluxes, it is impractical for most AGNs. Therefore, the empirical relation has been frequently adopted to estimate the BLR size and then derive the black hole masses  for AGNs in some samples of mostly radio-quiet objects (Laor 2000; McLure \\& Dunlop 2001; Wandel 2002), and of purely radio-loud objects (Lacy et al. 2001; Gu, Cao \\& Jiang 2001; Oshlack, Webster \\& Whiting 2002). However, the optical luminosity of  some radio-loud AGNs  (especially blazars), may not be a good indicator of ionizing luminosity, which is usually related to the UV/optical radiation from the accretion disk around the central black hole. The relativistic jets of blazar-type AGNs not only dominate the radio and high energy X-ray and $\\gamma$ ray radiation, but also significantly contribute to the optical luminosity in some cases (Scarpa \\& Urry 2002).  For example, many optical jets have been discovered recently in AGNs by the HST (Scarpa et al. 1999; Jester 2003; Parma \\etal 2003), which clearly suggests that the jets contribute significantly in the optical band. Furthermore, optical synchrotron radiation has been  detected for many other radio-loud AGNs (Whiting, Webster \\& Francis 2001; Chiaberge, Capetti, \\& Celloti 2002; Cheung \\etal 2003). Therefore, the measured optical continuum luminosity of  some  extremely radio-loud AGNs is significantly contributed by the optical radiation from the jets and may be much larger than the ionizing luminosity required to produce broad emission lines. Using the empirical relation between the BLR size and optical luminosity at 5100$\\AA$, which was obtained based on the sample of mostly radio-quiet AGNs (Kaspi \\etal 2000), one would significantly overestimate the actual BLR size and hence the black hole mass of these radio-loud AGNs. Oshlack et al. (2002) have shown that their estimated black hole masses would be lower if the synchrotron contribution to the optical flux is subtracted. However, it is not easy to make such a correction for a large sample of radio-loud AGNs. In addition, the contribution of the host galaxy to the optical continuum should also be taken into account especially when the host galaxy of AGNs can be resolved optically. Therefore, optical luminosity may not be a good indicator of photoionization luminosity of AGNs in some cases.  In this paper we will first derive an empirical relation between the BLR size and the H$_\\beta$ emission line luminosity for 34 AGNs in the sample of reverberation mapping studies. We  then argue that the BLR size obtained from the H$_\\beta$ luminosity is more reasonable at least for  some extremely radio-loud AGNs. Finally we  apply this new empirical relation to estimate the black hole masses of some quasars and compare them with previous results. ",
        "conclusions": "Using the empirical relation between the BLR size and optical continuum luminosity  possibly induces an overestimation of the BLR size and hence the black hole mass of  some extremely radio-loud AGNs because of the jet contribution to the optical luminosity. We derived another empirical relation between the BLR size and the emission line luminosity, and demonstrated that it can be used to estimate the BLR size and black hole mass of both radio-quiet and radio-loud AGNs. If the relativistic jets and host galaxy have significant contributions to the optical continuum, the emission line luminosity is probably a better tracer of ionizing luminosity. Comparisons of the estimated black hole masses with these two different empirical relations clearly indicate that the difference becomes significant if the radio-loudness of AGNs is larger. Using the  $R-L_{H_\\beta}$ relation may result in   more accurate estimations of black hole masses  of  some blazar-type AGNs.  In our study we focused on the possible effects of beaming on  the optical continuum in radio-loud AGNs and ignored the difference in BLR physics between radio-loud and radio-quiet AGNs. The modest correlation between the EW(H$_{\\beta}$) and the radio-loudness  may indicate the presence of beaming effects, though some other effects such as a lower covering factor of the BLR of radio-loud AGNs can also lead to smaller EW (H$_{\\beta}$) values.  Because currently we know little about the difference of BLR physics between radio-loud and radio-quiet AGNs, to prove the validity of our approach it is necessary to compare our estimated black hole mass with an independent estimate, for example, from the correlations of black hole mass with central velocity dispersion and host galaxy luminosity. Unfortunately, not many measured values of central velocity dispersion or host galaxy luminosity for extremely radio-loud AGNs are available. Although there are 8 objects in the sample of Brotherton (1996) with measured host magnitude (McLure \\& Dunlop 2001), the radio-loudness of these objects are mostly smaller than 1000 and thus the difference estimated with the $R-L_{H_\\beta}$ and $R-L_{5100\\AA}$ relations is rather small. The velocity dispersion measurements for radio-loud quasars are not available and the [OIII] profile in radio-loud AGNs may not be adopted to estimate the central velocity dispersion because of its complexity. Therefore, further imaging studies on the host galaxy and spectroscopic measurements of the central velocity dispersions of a large sample of extremely radio-loud quasars are still desired to confirm our results.    The advantage of using the  $R-L_{H_\\beta}$ relation is that we can estimate the black hole mass of AGNs with only two observed parameters, namely the H$_\\beta$ line luminosity and its FWHM, and it can be applied to a larger sample of AGNs with redshift smaller than 0.8. In principle, one can analogously investigate the relation between the BLR size and the luminosity of some ultraviolet emission lines such as MgII and CIV, which may be used to estimate the black hole mass of some high redshift AGNs. Some recent studies have suggested to use the ultraviolet continuum luminosity and the FWHM of ultraviolet emission lines to estimate the black hole mass of high redshift AGNs (Vestergaard 2002; McLure \\& Jarvis 2002). However, the ultraviolet continuum luminosity can similarly suffer the serious contaminations from jet and Blamer continuum, therefore the luminosity of  ultraviolet emission line again may be a better indicator of ionizing luminosity than the  ultraviolet continuum luminosity.   Finally, one should be cautious to the uncertainties in estimating the black hole mass of AGNs using the  $R-L_{H_\\beta}$ relation. Firstly, the variations of H$_\\beta$ emission line flux and its FWHM are common in AGNs. Estimating the black hole mass with the values of these two parameters in a single spectrum may lead to large errors. Secondly, the different inclination of the BLR may also significantly affect the results (McLure \\& Dunlop 2001; Wu \\& Han 2001).  If the BLR has a flatten geometry and the inclination of BLR is rather small, our derived values of black hole mass may significantly underestimated. However, the ratio of black hole masses estimated with two empirical relations does not depend on the inclination. Better understandings of BLR geometry and dynamics are absolutely needed to diminish the uncertainties in deriving the black hole mass of AGNs (Krolik 2001)."
    },
    "0403/astro-ph0403133_arXiv.txt": {
        "abstract": "To understand the nuclear stellar populations and star formation histories of the nuclei of spiral galaxies, we have obtained K-band nuclear spectra for 41 galaxies and H-band spectra for 20 galaxies in the ISO Atlas of Bright Spiral Galaxies.  In the vast majority of the subsample (80\\%), the near-infrared spectra suggest that evolved red stars completely dominate the nuclear stellar populations and that hot young stars are virtually non-existent.  The signatures of recent star formation activity are only found in 20\\% of the subsample, even though older red stars still dominate the stellar populations in these galaxies.  Given the dominance of evolved stars in most galaxy nuclei and the nature of the emission lines in the galaxies where they were detected, we suggest that nuclear star formation proceeds in the form of instantaneous bursts.  The stars produced by these bursts comprise only $\\sim$2 \\% of the total nuclear stellar mass in these galaxies, but we demonstrate how the nuclear stellar populations of normal spiral galaxies can be built up through a series of these bursts.  The bursts were detected only in Sbc galaxies and later, and both bars and interactions appeared to be sufficient but not necessary triggers for the nuclear star formation activity.  The vast majority of galaxies with nuclear star formation were classified as HII galaxies.  With one exception, LINERs and transition objects were dominated by older red stars, which suggested that star formation was not responsible for generating these galaxies' optical line emission. ",
        "introduction": "}  The nuclear stellar populations of galaxies can reveal important information on the history of nuclear star formation, which could then reveal the mechanisms behind triggering star formation.  Understanding nuclear stellar populations is therefore important to understanding the evolution of galaxies and the enhancement of star formation in the universe.  Near-infrared spectroscopy is ideal for studying stellar populations because it traces not only the presence of hot young stars though near-infrared hydrogen recombination lines, such as the Brackett $\\gamma$ line, but also the presence of red stars, particularly red supergiants, and the presence of shocks from supernovae activity.  Red supergiants and giants can be determined from the presence of various absorption features in the H- and K-bands, such as the CO and Si~{\\small I} absorption features in the H-band and the CO bands longward of 2.29~$\\mu$m \\citep{omo93}. Supernova activity can be inferred from the presence of Fe~{\\small II} emission lines associated with the supernova activity \\citep{omd89, getal91, c93, fw93, getal97, mds02, aetal03} and possibly even molecular hydrogen line emission \\citep{omd89, getal91, fw93} (although molecular hydrogen line emission may be produced by other sources; see \\citet{m94} and \\citet{vr97}). These diagnostics can be used together to characterize the different ages of the various stellar populations within galactic nuclei.  Once the nuclear stellar populations of spiral galaxies are well defined, they can be used to answer a number of specific scientific questions.  To begin with, this will illuminate how star formation proceeds in the nuclei of ``normal'' spiral galaxies, and whether that star formation is either continuous or a series of short bursts. Once we understand the functionality of star formation over time, we can determine how the nuclei of spiral galaxies have evolved.  We can also determine how morphological features, particularly bars, or how environmental influences, such as interactions, influence nuclear star formation activity.  Additionally, we can probe the link between star formation activity and Seyfert or LINER activity.  Much of the work at near-infrared wavelengths has focues on unusual objects with either AGN or very strong star formation activity. Recent surveys include surveys of ultraluminous infrared galaxies \\citep{getal95, metal99, msmetal01, bwd01}, luminous infrared galaxies \\citep{gjdetal97}, starbursts \\citep{e97, i00, cdd01}, Seyferts \\citep{i00, stl01, rkp02, betal02}, LINERs \\citep{letal98, aetal00, stl01}, and interacting galaxies \\citep{var98}.  However, relatively little near-infrared spectroscopic survey has been done with ``normal'' nearby spiral galaxies (see \\citet{mbpetal01} for an example).  Such surveys are necessary, though, for understanding nuclear star formation histories, for identifying the triggers that enhance nuclear star formation, and for setting a baseline for the H- and K-band emission from normal spiral galaxies.  Without such a baseline, the contribution of relatively quiescent stellar populations to H- and K-band emission in starbursts as well as the older stars' contribution to the overall mass and luminosities in these systems will remain unknown, and the enhancement in star formation activity in exotic systems like ultraluminous infrared galaxies will have no context.  Therefore, we have undertaken a near-infrared spectroscopic survey of a subset of the galaxies in the ISO Atlas of Bright Spiral Galaxies (\\cite{betal02a}, henceforth referred to as Paper 1) with the purpose of understanding the nuclear stellar populations and star formation histories of these objects.  The data includes K-band spectroscopy for 41 galaxies, with additional H-band spectroscopy for 20 galaxies.  We first discuss in Section~\\ref{s_data} the sample, the observations, and the data processing.  Next, we divide the galaxies into two groups: quiescent galaxies and non-quiescent galaxies.  In Section~\\ref{s_quies}, the spectra of the quiescent galaxies are described in detail, combined together to make a composite quiescent spectrum, and analyzed using population synthesis models.  The non-quiescent galaxies are also described in detail in Section~\\ref{s_nonquies}, with particular focus on how their nuclear star formation activity may be related to morphology, environment, and AGN activity.  After this, we determine the relative fractions of stars formed from starbursts in these systems.  Finally, in Section~\\ref{s_specphotcomp}, we compare quiescent and non-quiescent galaxies in plots of the $\\frac{f_{12\\mu m}}{f_K}$ ratio to the $f_{12\\mu m}$ or $f_K$ luminosities for the inner $15\\arcsec$ as an effort to justify using mid-infrared fluxes normalized by K-band fluxes (used in \\citet{betal02b}, henceforth referred to as Paper 2) to trace star formation activity. ",
        "conclusions": "\\subsection{Summary of Results}  In this sample, 33 of the 41 galaxies nuclei were found to have quiescent spectra.  These spectra consisted of continuum emission with many metal and CO absorption lines, which are characteristic of evolved red stars.  Using Starburst99, we found that the composite spectrum made from these quiescent galaxies' spectra was best represented by an instantaneous burst that took place $\\sim$180~Myr ago, although we caution that this is not necessarily supposed to reflect the true star formation history.  We found few differences between using Salpeter and other IMFs, but we did find that the continuous star formation scenario failed to reproduce the observed equivalent widths.  The remaining 8 galaxies have nuclear spectra that include emission lines. In many cases, the underlying continuum exhibited the same slope and contained the same absorption features that were present in the composite quiescent spectra.  On top of this quiescent continuum, the galaxy may produce either photoionization lines such as Brackett-$\\gamma$ and He~{\\small I} emission, shock excitation lines such as H$_2$ lines and Fe~{\\small II} emission, or both.  The lines present would reflect the age of the young stellar population, with photoionization lines present up to 8~Myr after a burst and shock excitation lines present from 3.5 to 36~Myr after a burst. Some of the systems showed relatively point-like spectral line emission, whereas others exhibited emission in complex structures.  One galaxy, where the nuclear emission was quiescent, had a strong extranuclear starburst that was detected and included in some parts of the analysis.  These non-quiescent galaxies share many characteristics.  All of them are Sbc galaxies or later, which suggests that these instantaneous nuclear bursts are more likely in the later-type galaxies.  6 of the 8 galaxies were definitely barred, which suggested that bars may play a role in triggering nuclear star formation.  However, 6 of the 8 galaxies either were interacting or had been involved in interactions, which suggests that interactions are just as important as bars in triggering nuclear star formation.  In relation to AGN activity, it appears that most non-quiescent galaxies are H~{\\small II} galaxies and that most LINERs and transition objects are quiescent.  The average ratio of young to old stellar masses for the non-quiescent systems was found to be $\\sim$2~\\%.  Given that these 2~\\% enhancements occur in $\\sim$20~\\% of the galaxies in the sample, these enhancements are found to be consistent with nuclear stellar populations forming through a series of instantaneous bursts.  Finally, we found that the $\\frac{f_{MIR}}{f_K}$ ratios functioned very well in separating non-quiescent galaxies from quiescent galaxies.  The exception was with galaxies that contained only shock excitation lines, but the ratio was expected to decrease signficantly for these galaxies.  The data and results presented here now form a baseline for comparison to other galaxies with more exotic star formation or AGN activity.  Furthermore, the composite quiescent spectra and the conversion factors developed here should provide additional analysis tools for studying enhanced star formation activity.  \\subsection{Future Work}  While the spectroscopy presented here surveys a range of different kinds of galaxies, it is only a limited sample.  A broad range of morphologies were included in the sample, but not a broad range of AGN types; the spectroscopic sample only contained two Seyferts.  Furthermore, as Virgo Cluster galaxies were left out of the original sample, no data are available for comparing field and cluster galaxies.  A larger near-infrared spectroscopic survey of a complete, volume-limited sample of galaxies would provide additional constraints on these stochastic bursts of star formation, such as the frequency of their occurence, their relation to bars, environment, Seyfert activity, and LINER activity, and the range of masses formed within these bursts.  Single slit spectrometers are good for sampling the pointlike nuclei of many nearby galaxies and can return spectroscopic data that accurately represent the stellar populations and total star formation activity within these galaxies.  However, when the nuclear star formation structures become more complex, as is the case for NGC~3556 and NGC~4100, then a single slit could possibly miss much of the spectral line emission coming from within the centers of these galaxies.  Therefore, integral field near-infrared spectroscopic surveys of nearby spiral galaxies would appear to be warranted.  Such a survey would not only provide more accurate data on the stellar populations and total line emission from the nuclei of galaxies but also provide data on the number of galaxies with complex, extended nuclear structures that could lead to additional clues on the phenomenology of nuclear star formation activity."
    },
    "0403/astro-ph0403419_arXiv.txt": {
        "abstract": "The recently discovered $z=10$ galaxy (Pello \\etal\\ 2004) has a strong Ly$\\alpha$ emission line that is consistent with being surprisingly symmetric, even given the relatively poor quality of its spectrum.  The blue wing of a Ly$\\alpha$ line originating at high redshift should be strongly suppressed by resonant hydrogen absorption along the line of sight, an expectation borne out by the observed asymmetric shapes of the existing sample of \\lya\\ emitting sources at lower redshifts ($3< z < 6.7$).  Absorption on the blue side of the line of the Pello \\etal\\ source could be reduced if the intergalactic medium (IGM) in the vicinity of the galaxy is highly ionized, but we show that this requires an unrealistically high ionizing emissivity.  We suggest instead that the Ly$\\alpha$ emitting gas be receding relative to the surrounding gas with a velocity of $\\gsim 35$ km/s, a large velocity that is plausible only if the galaxy is part of a larger system (group of galaxies) with a velocity dispersion $\\gsim 35$ km/s. We find that with this velocity shift, the observed strength and shape of the line is still consistent with the galaxy being surrounded by its own Str\\\"omgren sphere embedded in a fully neutral IGM.  More generally, we predict that at any given redshift, the bright Ly$\\alpha$ emitters with broader lines would exhibit stronger asymmetry than fainter ones.  Bright galaxies with symmetric Ly$\\alpha$ lines may be signposts for groups and clusters of galaxies, within which they can acquire random velocities comparable to or larger than their linewidths. ",
        "introduction": "\\label{sec:introduction}  Three independent observations combine to paint a complex picture of the cosmological reionization process.  First, the recent quasar absorption spectrum observations by the Sloan Digital Sky Survey show strong evidence that the reionization process completes at $z\\sim 6$ (Becker \\etal\\ 2001; Fan \\etal\\ 2002; Cen \\& McDonald 2002). Second, the latest {\\it Wilkinson Microwave Anisotropy Probe (WMAP)} observations detect a high Thomson scattering optical depth, suggesting that the intergalactic medium (IGM) experienced a significantly ionized state at high redshift somewhere between $z=15-25$ (Kogut \\etal\\ 2003) for (at least) a significant redshift interval.  This is somewhat contradicted by the third observational line of evidence of the intergalactic medium having a relatively high temperature at $z\\sim 3-4$, which requires a reionization epoch no earlier than redshift $z=9-10$ (Hui \\& Haiman 2003; Theuns \\etal\\ 2002).  While the overall picture is consistent with a pre-WMAP, physically motivated double reionization model (Cen 2003; Wyithe \\& Loeb 2003), a detailed probe of the ionization state of the IGM at high redshift is sorely wanted.  Ly$\\alpha$ emission lines from high--redshift sources can serve as probes of the ionization state of the IGM. The damping wing of the Gunn--Peterson (GP) absorption from the IGM can cause a characteristic absorption feature (Miralda-Escud\\'e 1998).  For a Ly$\\alpha$ emitting galaxy embedded in a partly neutral IGM, the absorption produces conspicuous effects, i.e. attenuating the emission line, making it asymmetric, and shifting its apparent peak to longer wavelengths (Haiman 2002; Santos 2003).  In practice, the expectation is that strong conclusions cannot be drawn from a single galaxy.  For example, the relatively strong Ly$\\alpha$ line of the $z=6.6$ galaxy discovered by Hu \\etal (2002) is still consistent with being embedded in a neutral IGM, but surrounded by its own ionized Str\\\"omgren sphere (Haiman 2002; Santos 2003).  The recent claim of the detection of a Ly$\\alpha$ emitting galaxy at $z=10$ (Pello \\etal\\ 2004; hereafter P04) provides a new opportunity to study the IGM at $z=10$.  This source is especially interesting, since at its high inferred redshift, absorption by the IGM should increase significantly.  In this {\\it Letter}, we examine both the observed shape and overall attenuation of the detected Ly$\\alpha$ line, in models where the line is processed through the IGM. We find that in order to achieve the observed symmetry of the Ly$\\alpha$ line profile, the emitting gas in the galaxy must be receding faster than the surrounding gas by at least $35$km/s.  Given this required recessional velocity, we find that the P04 source is marginally consistent with being embedded in a fully neutral IGM at $z=10$, with the line suffering an attenuation by a factor of 30.  Throughout this paper, we assume the background cosmology to be flat $\\Lambda$CDM with $(\\Omega_\\Lambda,\\Omega_{\\rm m},\\Omega_{\\rm b}, h)=(0.7,0.3,0.04,0.7)$. ",
        "conclusions": "We considered the implications of the detection of the symmetric, highly attenuated Ly$\\alpha$ emission line from a candidate $z=10$ galaxy.  We find that the observed symmetry of the Ly$\\alpha$ emission line can be accounted for if the emitting galaxy is receding relative to the surrounding absorbing gas by a velocity of at least $35$km/s. Such a relative velocity is mostly plausibly achieved if this detected galaxy is a member of a larger system with a velocity dispersion in excess of $35$km/s.  Thus, while the difficulties and challenges associated with such observations are formidable, it is not a great surprise, in principle, to be able to detect galaxies with such Ly$\\alpha$ emission lines at high redshift.  However, with the required recessional velocity, this galaxy does not place a strong constraint on the ionization state of the IGM, given various uncertainties in the current data and lack of handle of the intrinsic absorption. A fully neutral universe, while not preferred, is still consistent with the observation.  A moderate increase in the sample size of such high redshift galaxies will be highly valuable in statistical inferences for the ionization state of the IGM, based on the systematic dependence of the line properties on redshift and luminosity (Haiman 2002; Rhoads \\& Malhotra 2002). More urgent in the short term is to obtain a higher quality spectrum to better characterize the line profile."
    },
    "0403/astro-ph0403305_arXiv.txt": {
        "abstract": "\\noindent The  Generalized Chaplygin Gas (GCG) with the equation of state $p = - \\frac{A}{{\\rho}^{\\alpha}}$ was recently proposed as a candidate for dark energy in the Universe. In this paper we confront the GCG with SNIa data using avaliable samples. Specifically we have tested the GCG cosmology in three different classes of models with (1) $\\Omega_m= 0.3$, $\\Omega_{Ch}= 0.7$; (2) $\\Omega_m= 0.05$, $\\Omega_{Ch}= 0.95$ and (3) $\\Omega_m = 0$, $\\Omega_{Ch} =  1$, as well as a model without prior assumptions on $\\Omega_m$. The best fitted models are obtained by minimalizing the $\\chi^2$ function. We supplement our analysis with confidence intervals in the $(A_0, \\alpha)$ plane by marginalizing the probability density functions over remaining parameters assuming uniform priors. We have also derived one-dimensional probability distribution functions for $\\Omega_{Ch}$ obtained from joint marginalization over $\\alpha$ and $A_0$. The maximum value of such PDF provides the most probable value of $\\Omega_{Ch}$ within the full class of GCG models. The general conclusion is that SNIa data give support to the Chaplygin gas (with $\\alpha =  1$). However noticeable preference of $A_0$ values close to 1 means that the  $\\alpha$ dependence becomes insignificant. It is reflected on one dimensional PDFs for $\\alpha$ which turned out to be flat meaning that the power of present supernovae data to discriminate between various GCG models (differing by $\\alpha$) is weak. Extending our analysis by relaxing the prior assumption of the flatness of the Universe leads to the result that even though the best fitted values of $\\Omega_k$ are formally non-zero, still they are close to the flat case. Our results show clearly that in GCG cosmology distant (i.e. $z >1$) supernovae should be brighter than in $\\Lambda$CDM model. Therefore one can expect that future supernova experiments (e.g., SNAP) having access to higher redshifts will eventually resolve the issue whether the dark energy content of the Universe could be described as a the Chaplygin gas. Moreover, it would be possible to differentiate between models with various value of $\\alpha$ parameter and/or discriminated between GCG, Cardassian and $\\Lambda$CDM models. This discriminative power of the forthcoming mission has been demonstrated on simulated SNAP data. ",
        "introduction": "For a couple of years two independent observational programs --- the high redshift supernovae surveys (Perlmutter et al. 1999) and CMBR small scale anisotropy measurements (de Bernardis et al. 2000, Benoit et al. 2003, Hinshaw et al. 2003) have brought a new picture of the Universe in the large. While interpreted  ithin the FRW models results of these programs suggest that our Universe is flat (as inferred from the location of acoustic peaks in CMBR power spectrum) and presently accelerates its expansion (as inferred from the SNIa Hubble diagram). Combined with the independent knowledge about the amount of baryons and CDM estimated to be $\\Omega_m =  0.3$ (Turner 2002) it follows that about $\\Omega_X =  0.7$ fraction of critical density $\\rho_{cr} =  \\frac{3 c^2 H_0^2}{8 \\pi G}$ should be contained in a mysterious component called ``dark energy''. The most obvious candidate for this smooth component permeating the Universe is the cosmological constant $\\Lambda$ representing the energy of the vacuum. Well known fine tuning problems led many people to seek beyond the $\\Lambda$ framework, and the concept of the quintessence had been conceived. Usually the quintessence is described in a phenomenological manner, as a scalar  field with an appropriate potential (Ratra \\& Peebles 1988, Caldwell, Dave \\& Steinhardt 1995, Frieman, Stebbins \\& Waga 1995). It turns out, however, that quintessence program also suffers from its own fine tuning problems (Kolda \\& Lyth 1999).  In 1904 Russian physicist Chaplygin introduced the exotic equation of state $p =  - \\frac{A}{\\rho}$ to discribe an adiabatic aerodynamic process (Chaplygin 2004). The attractiveness of this equation of state in the context dark energy models comes mainly from the fact that it gives the unification of both dark energy (postulated in cosmology to explain current aceleration of the Universe) and clustered dark matter which is postulated in astrophysics to explain the flat rotation curves of spiral galaxies. It is interesting that the Chaplygin gas can be derived from the quintessence Lagrangian for the scalar field $\\phi$ with some potential and also from the Born-Infeld form of the Lagrangian (Kamenshchik, Moschella \\& Pasquier 2001). The Chaplygin equation of state has some interesting connections with string theory and it admits the interpretation in the framework of brane cosmologies (Jackiw 2000).  Recently the so called Chaplygin gas (Kamenshchik, Moschella \\& Pasquier 2001, Fabris, Gon{\\c c}alves \\& de Souza 2002, Szyd{\\l}owski \\& Czaja 2004) --- a hypothetical component with the equation of state $p =  - \\frac{A}{\\rho}$ --- was proposed as a challenge to the above mentioned candidates for dark energy. This, also purely phenomenological, entity has interesting connections with string theory (Ogawa 2000). Currently its generalizations admitting the equation of state $p=  - \\frac{A}{\\rho^{\\alpha}}$ where $0\\le \\alpha \\le 1$ have been proposed (Bento, Bertolami \\& Sen 2002, Carturan \\& Finelli 2002a).  In this paper we confront the Generalized Chaplygin Gas with the SNIa data. At this point our choice of Generalized Chaplygin Gas cosmologies deserves a sort of justification. There are two approaches in the literature. First one is phenomenological, namely having no preferred theory of dark energy responsible for acceleration of the Universe one characterizes it as a cosmic fluid  with an equation of state $p_X = w \\rho_X$ where $w \\geq -1$ (see e.g. (Chiba 1998,Turner \\& White 1997) and an immense literature that appeared thereafter). Because, as already mentioned above, a strain of ideas about dark energy is associated with an evolving scalar there are good reasons to expect that cosmic equation of state could be time dependent i.e. $ w = w(t) = w(z)$ (e.g. Weller \\& Albrecht 2001, Maor et al. 2001 and many others thereafter). This approach seems attractive from the perspective of analyzing observational data such like supernovae surveys and indeed this approach was taken while first analyzing the data (Perlmutter et al. 1999, Knop et al. 2003 or Riess et al. 2004). However even though such analysis places constrains on {\\it any} potential theory that might explain the dark energy phenomenon, ultimately one always ends up at testing a {\\it specific} theory. Along this line there appeared attempts to reconstruct the scalar potential, if the scalar field was responsible for dark energy (e.g. Alam et al. 2003 and references therein). Our approach goes along this philosophy but instead is devoted to the Generalized Chaplygin Gas which is being recently considered as candidate to unified dark matter-energy component (i.e. responsible for both clustering and accelerated expansion (Makler, de Oliveira \\& Waga 2003).  The cosmological models with the Generalized Chaplygin Gas have also many special features which make them atractive. In standard cosmological model one can clearly distinguish the epochs of radiation domination followed by (ordinary) matter domination (with decelerated expansion). As mentioned above supernova data suggest that the epoch of decelerated expansion ended and switched to accelerated epoch --- dominated by dark energy. The Generalized Chaplygin Gas models describe smoothly the transition from the decelerated to accelerated epochs. They represent the simplest deformation of concordance $\\Lambda$CDM (Gorini et al. 2004). And moreover, they propose a new unified macroscopic (phenomenological) description of both dark energy and dark matter. This places them in a distinguished position from the point of view of Occam's razor principle. It should be also noted that the Genaralized Chaplygin Gas model allowed us to explain presently observed acceleration of the Universe without the cosmological constant and/or modification of Einstein's equations.  If one takes seriously given dark energy scenario (necessary to explain cosmic acceleration) one should also consider the behaviour of perturbations in such a universe. In the framework of quintessence models with the barotropic equation of state (i.e. $p= w \\rho$ and $w=const$) one faces the problem of instabilities in short scales. This appears because the speed of sound squared (equal here to $w$) is negative (and constant). Calculation of the sound speed in Generalized Chaplygin Gas model (see below) reveal its non-barotropic nature. The perturbations in GCG models are stable in short scales even in an accelerating phase (Carturan \\& Finelli 2002a). Moreover, they behave like dust perturbations when Chaplygin Gas is in dust regime.  Another motivation for studying Generalized Chaplygin Gas models goes from theoretical physics --- specifically from attempts to describe the dark energy in terms of the Lagrangian for a tachyonic field (Garousi 2000, Sen 2002). Of course it would be nice to have a description of dark energy in terms of the non quintessence Lagrangian as it  describes the nature of dark energy while the cosmological constant is only phenomelogical and effective description. One should also note that the Generalized Chaplygin Gas equation of state arises in modern physics in the context of brane models (Bordemann \\& Hoppe 1993, Kamenshchik, Moschella \\& Pasquier 2001, Randall \\& Sundrum 1999) where the Generalized Chaplygin Gas manifests itself as an effect of immersion of our Universe in multidimensional bulk space.  Generalized Chaplygin gas models have been intensively studied in the literature and in particular they have been tested against supernovae data (Makler, de Oliveira \\& Waga 2003, Avelino et al. 2003, Collistete et al. 2003), lensing statistics (Dev, Alcaniz \\& Jain 2003), CMBR measurements (Bento, Bertolami \\& Sen 2003a, 20003b, Carturan \\& Finelli 2003b, Amendola et al. 2003), age-redshift relation (Alcaniz, Jain \\& Dev 2003), x-ray luminosities of galaxy clusters (Cunha, Lima \\& Alcaniz 2003) or from the large scale structure considerations (Bean \\& Dor{\\'e} 2003, Multamaki, Manera \\& Gaztanaga 2003, Bilic et al. 2003). Perspectives to distinguish between Generalized Chaplygin Gas, brane-world scenarios and quintessence in forthcoming gravity wave experiments has been discussed in (Biesiada 2003). Although the results are in general mutually consistent there was no strong convergence to unique values of $A_0$, $\\alpha$ parameters characterizing Chaplygin gas equation of state.  Makler, de Oliveira \\& Waga (2003) have considered the FRW model filled completely with Generalized Chaplygin Gas and concluded that whole class of such models is consistent with current SNIa data although the value of $\\alpha =  0.4$ is favoured. This result has been confirmed by our analysis (class (3) models). However, when the existing knowledge about baryonic matter content of the Universe was incorporated into the study our results were different from Makler, de Oliveira \\& Waga (2003) who found that $\\alpha =  0.15$ was preferred (assuming $\\Omega_m =  0.04$ which is very close to our assumption for class (2)  models).  As noticed by Bean \\& Dor{\\'e} (2003) Generalized Chaplygin Gas models have an inherent degeneracy with cosmological constant models as far as background evolution is concerned, and therefore they have a good fit with SNIa data. These degeneracies disappear at the level of evolution of perturbations and hence confrontation with CMBR spectrum would be decisive. Using available data on the position of CMBR peaks measured by BOOMERANG (de Bernardis et al. 2000) and Archeops (Benoit et al. 2003, Hinshaw et al. 2003, Bento, Bertolami \\& Sen (2002) obtained the following constraints: $0.81 \\leq A_0 \\leq 0.85$ and $0.2 \\leq \\alpha \\leq 0.6$ at $68 \\%$ CL in the model representative of our class (2) (i.e. with $\\Omega_m = 0.05$ assumed). Another estimation of the parameter $\\alpha$ was done by Amendola et al. (2003) with WMAP Data. The obtained the $0\\le \\alpha < 0.2$ at $95\\%$ confidence level.   Using the angular size statistics for extragalactic sources combined with SNIa data it was found in (Alcaniz \\& Lima 2003) that in the the $\\Omega_m =0.3$ and $\\Omega_{Ch}=0.7$ scenario best fitted values of model parameters are $A_0=0.83$ and $\\alpha=1.$ respectively. Recent paper by Bertolami et al. (2004) in which Generalized Chaplygin Gas models have been analyzed against Tonry et al. (2003) supernovae data relaxing the prior assumption on flatness suggests, surprisingly as the authors admit, the preference of $\\alpha > 1$.    \\begin{deluxetable}{@{}p{1.5cm}rrrrrrr} \\tabletypesize{\\scriptsize} \\tablewidth{0pt} \\tablecaption{Results of statistical analysis of Generalized Chaplygin Gas model (with marginalization over ${\\cal M}$) performed on analyzed samples of SNIA (A, C, K6, K3, TBI, TBII, Silver, Gold) as a minimum $\\chi^2$ best-fit (denoted BF) and with the maximum likelihood method (denoted L). First two rows for each sample refer to no prior on $\\Omega_m$. The same analysis was repeated with fixed priors $\\Omega_m=0.0$, $\\Omega_m=0.05$ and $\\Omega_m=0.3$.} \\label{tab:1} \\tablehead{\\colhead{sample} & \\colhead{$\\Omega_m$} & \\colhead{$\\Omega_{Ch}$} & \\colhead{$A_0$} & \\colhead{$\\alpha$} & \\colhead{$\\mathcal{M}$} & \\colhead{$\\chi^2$}& \\colhead{method}} \\startdata A   &  0.00 & 1.00 & 0.77 & 1.00 &-3.39 & 95.4 &  BF  \\\\ &  0.17 & 0.83 & 0.83 & 0.00 &-3.36 & ---  &  L   \\\\ &  0.00 & 1.00 & 0.77 & 1.00 &-3.39 & 95.4 &  BF  \\\\ &  0.00 & 1.00 & 0.73 & 1.00 &-3.38 & ---  &  L   \\\\ &  0.05 & 0.95 & 0.80 & 1.00 &-3.39 & 95.4 &  BF  \\\\ &  0.05 & 0.95 & 0.76 & 1.00 &-3.38 & ---  &  L   \\\\ &  0.30 & 0.70 & 0.96 & 1.00 &-3.39 & 95.8 &  BF  \\\\ &  0.30 & 0.70 & 0.96 & 0.00 &-3.38 & ---  &  L   \\\\ \\tableline % C   &  0.00 & 1.00 & 0.80 & 1.00 &-3.44 & 52.9 &  BF  \\\\ &  0.15 & 0.85 & 0.86 & 0.00 &-3.41 & ---  &  L   \\\\ &  0.00 & 1.00 & 0.80 & 1.00 &-3.44 & 52.9 &  BF  \\\\ &  0.00 & 1.00 & 0.76 & 0.49 &-3.43 & ---  &  L   \\\\ &  0.05 & 0.95 & 0.83 & 1.00 &-3.44 & 53.0 &  BF  \\\\ &  0.05 & 0.95 & 0.79 & 0.11 &-3.43 & ---  &  L   \\\\ &  0.30 & 0.70 & 0.99 & 1.00 &-3.42 & 53.3 &  BF  \\\\ &  0.30 & 0.70 & 0.99 & 0.00 &-3.39 & ---  &  L   \\\\ \\tableline K6  &  0.00 & 1.00 & 0.81 & 1.00 &-3.52 & 55.3 &  BF  \\\\ &  0.10 & 0.90 & 0.88 & 0.00 &-3.51 & ---  &  L   \\\\ &  0.00 & 1.00 & 0.81 & 1.00 &-3.52 & 55.3 &  BF  \\\\ &  0.00 & 1.00 & 0.78 & 0.71 &-3.52 & ---  &  L   \\\\ &  0.05 & 0.95 & 0.84 & 1.00 &-3.52 & 55.4 &  BF  \\\\ &  0.05 & 0.95 & 0.81 & 0.06 &-3.52 & ---  &  L   \\\\ &  0.30 & 0.70 & 1.00 & 1.00 &-3.51 & 55.9 &  BF  \\\\ &  0.30 & 0.70 & 1.00 & 0.00 &-3.49 & ---  &  L   \\\\ \\tableline K3  &  0.00 & 1.00 & 0.85 & 1.00 &-3.48 & 60.4 &  BF  \\\\ &  0.11 & 0.89 & 0.88 & 0.00 &-3.45 & ---  &  L   \\\\ &  0.00 & 1.00 & 0.85 & 1.00 &-3.48 & 60.4 &  BF  \\\\ &  0.00 & 1.00 & 0.80 & 0.30 &-3.47 & ---  &  L   \\\\ &  0.05 & 0.95 & 0.87 & 1.00 &-3.47 & 60.4 &  BF  \\\\ &  0.05 & 0.95 & 0.84 & 0.00 &-3.47 & ---  &  L   \\\\ &  0.30 & 0.70 & 1.00 & 1.00 &-3.44 & 61.4 &  BF  \\\\ &  0.30 & 0.70 & 1.00 & 0.00 &-3.42 & ---  &  L   \\\\ \\tableline TBI &  0.00 & 1.00 & 0.79 & 1.00 &15.895&273.9 &  BF  \\\\ &  0.00 & 1.00 & 0.81 & 1.00 &15.905& ---  &  L   \\\\ &  0.00 & 1.00 & 0.79 & 1.00 &15.895&273.8 &  BF  \\\\ &  0.00 & 1.00 & 0.75 & 1.00 &15.905& ---  &  L   \\\\ &  0.05 & 0.95 & 0.82 & 1.00 &15.895&274.0 &  BF  \\\\ &  0.05 & 0.95 & 0.78 & 1.00 &15.915& ---  &  L   \\\\ &  0.30 & 0.70 & 0.97 & 1.00 &15.915&275.8 &  BF  \\\\ &  0.30 & 0.70 & 0.96 & 0.00 &15.915& ---  &  L   \\\\ \\tableline TBII &  0.00 & 1.00 & 0.78 & 1.00 &15.915&186.5 &  BF  \\\\ &  0.00 & 1.00 & 0.81 & 1.00 &15.925& ---  &  L   \\\\ &  0.00 & 1.00 & 0.78 & 1.00 &15.915&186.5 &  BF  \\\\ &  0.00 & 1.00 & 0.75 & 1.00 &15.915& ---  &  L   \\\\ &  0.05 & 0.95 & 0.81 & 1.00 &15.915&186.6 &  BF  \\\\ &  0.05 & 0.95 & 0.78 & 1.00 &15.925& ---  &  L   \\\\ &  0.30 & 0.70 & 0.97 & 1.00 &15.925&188.4 &  BF  \\\\ &  0.30 & 0.70 & 0.96 & 0.00 &15.935& ---  &  L   \\\\ \\tableline Silver&  0.00 & 1.00 & 0.82 & 1.00 &15.945&229.4 &  BF  \\\\ &  0.00 & 1.00 & 0.84 & 1.00 &15.945& ---  &  L   \\\\ &  0.00 & 1.00 & 0.82 & 1.00 &15.945&229.4 &  BF  \\\\ &  0.00 & 1.00 & 0.79 & 1.00 &15.955& ---  &  L   \\\\ &  0.05 & 0.95 & 0.85 & 1.00 &15.945&229.6 &  BF  \\\\ &  0.05 & 0.95 & 0.81 & 1.00 &15.955& ---  &  L   \\\\ &  0.30 & 0.70 & 0.99 & 1.00 &15.965&232.3 &  BF  \\\\ &  0.30 & 0.70 & 0.99 & 0.00 &15.965& ---  &  L   \\\\ \\tableline Gold &  0.00 & 1.00 & 0.81 & 1.00 &15.945&173.7 &  BF  \\\\ &  0.00 & 1.00 & 0.83 & 1.00 &15.955& ---  &  L   \\\\ &  0.00 & 1.00 & 0.81 & 1.00 &15.945&173.7 &  BF  \\\\ &  0.00 & 1.00 & 0.77 & 1.00 &15.955& ---  &  L   \\\\ &  0.05 & 0.95 & 0.84 & 1.00 &15.945&173.8 &  BF  \\\\ &  0.05 & 0.95 & 0.80 & 1.00 &15.955& ---  &  L   \\\\ &  0.30 & 0.70 & 0.99 & 1.00 &15.965&175.6 &  BF  \\\\ &  0.30 & 0.70 & 0.99 & 0.00 &15.965& ---  &  L   \\\\ \\enddata \\end{deluxetable}   \\begin{table} \\caption{Generalized Chaplygin Gas model parameter values obtained from the marginal probability density functions calculated on Perlmutter, Knop, Tonry/Barris and Riess samples with $\\Omega_m$ prior relaxed.} \\label{tab:2} \\begin{tabular}{@{}p{1.5cm}rrrr} \\hline \\hline sample & $\\Omega_m$ & $\\Omega_{Ch}$ &  $A_0$ & $\\alpha$ \\\\ \\hline A &  $ 0.17^{+0.08}_{-0.17}$ & $0.83^{+0.17}_{-0.08}$   & $0.83^{+0.14}_{-0.09}$ & $-0.0^{+0.67}$ \\\\ C & $0.15^{+0.08}_{-0.15}$ & $0.85^{+0.15}_{-0.08}$   & $0.86^{+0.13}_{-0.10}$ & $0.0^{+0.66}$  \\\\ K6&  $ 0.10^{+0.11}_{-0.10}$ & $0.90^{+0.10}_{-0.11}$   & $0.88^{+0.12}_{-0.08}$ & $-0.0^{+0.66}$ \\\\ K3& $0.11^{+0.07}_{-0.11}$ & $0.89^{+0.11}_{-0.07}$   & $0.88^{+0.11}_{-0.05}$ & $0.0^{+0.66}$\\\\ TBI& $0.00^{+0.21}$ & $1.00_{-0.21}$   & $0.81^{+0.12}_{-0.07}$ & $1.0_{-0.60}$\\\\ TBII&$0.00^{+0.21}$ & $1.00_{-0.21}$   & $0.81^{+0.12}_{-0.07}$ & $1.0_{-0.62}$\\\\ Silver&$0.00^{+0.18}$ & $1.00_{-0.18}$   & $0.84^{+0.09}_{-0.06}$ & $1.0_{-0.59}$\\\\ Gold&  $0.00^{+0.20}$ & $1.00_{-0.20}$   & $0.83^{+0.11}_{-0.07}$ & $1.0_{-0.64}$\\\\ \\hline \\end{tabular} \\end{table}    \\begin{table} \\caption{Generalized Chaplygin Gas model parameter values obtained from the marginal probability density functions calculated on Perlmutter, Knop, Tonry/Barris and Riess samples.  The analysis was done with fixed $\\Omega_m=0.0$, $\\Omega_m=0.05$ and $\\Omega_m=0.3$.} \\label{tab:3} \\begin{tabular}{@{}p{1.5cm}rrrr} \\hline \\hline sample & $\\Omega_m$ & $\\Omega_{Ch}$ &  $A_0$ & $\\alpha$ \\\\ \\hline A     &$ 0.00$ & $1.00$ & $0.73^{+0.08}_{-0.10}$ & $1.0_{-0.63}$ \\\\ &$ 0.05$ & $0.95$ & $0.76^{+0.08}_{-0.09}$ & $1.0_{-0.66}$ \\\\ &$ 0.30$ & $0.70$ & $0.96^{+0.04}_{-0.09}$ & $0.0^{+0.65}$ \\\\ C     &$ 0.00$ & $1.00$ & $0.76^{+0.08}_{-0.10}$ & $0.49^{+0.36}_{-0.35}$ \\\\ &$ 0.05$ & $0.95$ & $0.79^{+0.08}_{-0.11}$ & $0.41^{+0.27}_{-0.41}$ \\\\ &$ 0.30$ & $0.70$ & $0.99^{+0.01}_{-0.11}$ & $0.0^{+0.64}$ \\\\ K6    &$ 0.00$ & $1.00$ & $0.78^{+0.07}_{-0.09}$ & $0.71^{+0.29}_{-0.40}$ \\\\ &$ 0.05$ & $0.95$ & $0.81^{+0.08}_{-0.09}$ & $0.06^{+0.61}_{-0.06}$ \\\\ &$ 0.30$ & $0.70$ & $1.00_{-0.10}        $ & $0.0^{+0.64}$ \\\\ K3    &$ 0.00$ & $1.00$ & $0.80^{+0.06}_{-0.06}$ & $0.30^{+0.39}_{-0.30}$ \\\\ &$ 0.05$ & $0.95$ & $0.84^{+0.05}_{-0.07}$ & $0.0^{+0.67}$ \\\\ &$ 0.30$ & $0.70$ & $1.00_{-0.06}        $ & $0.0^{+0.63}$ \\\\ TBI   &$ 0.00$ & $1.00$ & $0.75^{+0.04}_{-0.05}$ & $1.0_{-0.54}$ \\\\ &$ 0.05$ & $0.95$ & $0.78^{+0.04}_{-0.06}$ & $1.0_{-0.55}$ \\\\ &$ 0.30$ & $0.70$ & $0.96^{+0.04}_{-0.04}$ & $0.0^{+0.67}$ \\\\ TBII  &$ 0.00$ & $1.00$ & $0.75^{+0.04}_{-0.06}$ & $1.0_{-0.54}$ \\\\ &$ 0.05$ & $0.95$ & $0.78^{+0.04}_{-0.06}$ & $1.0_{-0.54}$ \\\\ &$ 0.30$ & $0.70$ & $0.96^{+0.04}_{-0.04}$ & $0.0^{+0.67}$ \\\\ Silver&$ 0.00$ & $1.00$ & $0.79^{+0.03}_{-0.05}$ & $1.0_{-0.52}$ \\\\ &$ 0.05$ & $0.95$ & $0.81^{+0.04}_{-0.04}$ & $1.0_{-0.54}$ \\\\ &$ 0.30$ & $0.70$ & $0.99^{+0.01}_{-0.03}$ & $0.0^{+0.64}$ \\\\ Gold  &$ 0.00$ & $1.00$ & $0.77^{+0.04}_{-0.05}$ & $1.0_{-0.58}$ \\\\ &$ 0.05$ & $0.95$ & $0.80^{+0.04}_{-0.05}$ & $1.0_{-0.59}$ \\\\ &$ 0.30$ & $0.70$ & $0.99^{+0.01}_{-0.04}$ & $0.0^{+0.64}$ \\\\ \\hline \\end{tabular} \\end{table}  \\begin{table} \\noindent \\caption{Results of statistical analysis of Generalized Chaplygin Gas models with flat prior relaxed and with marginalization over ${\\cal M}$ performed on Knop Samples K3 as a minimum $\\chi^2$ best-fit(denoted BF) and with the maximum likelihood method (denoted L). First two rows refer to no prior on $\\Omega_m$. The same analysis was repeated with fixed $\\Omega_m=0.0$, $\\Omega_m=0.05$ and $\\Omega_m=0.3$.} \\label{tab:4} \\begin{tabular}{@{}p{1.5cm}rrrrrrrr} \\hline \\hline sample & $\\Omega_k$ & $\\Omega_m$ & $\\Omega_{Ch}$ &  $A_0$ & $\\alpha$ & $\\mathcal{M}$ & $\\chi^2$& method \\\\ \\hline K3  & -0.19  &  0.00 & 1.19 & 0.82 & 1.00 &-3.48 & 60.3 &  BF  \\\\ & -0.60  &  0.00 & 1.26 & 0.89 & 0.00 &-3.46 & ---  &  L   \\\\ & -0.25  &  0.00 & 1.25 & 0.82 & 1.00 &-3.49 & 60.3 &  BF  \\\\ &  0.10  &  0.00 & 0.90 & 0.76 & 0.00 &-3.46 & ---  &  L   \\\\ & -0.28  &  0.05 & 1.23 & 0.84 & 1.00 &-3.49 & 60.3 &  BF  \\\\ &  0.05  &  0.05 & 0.90 & 0.78 & 0.00 &-3.47 & ---  &  L   \\\\ & -0.48  &  0.30 & 1.18 & 0.93 & 0.97 &-3.49 & 60.3 &  BF  \\\\ & -0.35  &  0.30 & 1.05 & 0.88 & 0.00 &-3.47 & ---  &  L   \\\\ Gold& -0.12  &  0.00 & 1.12 & 0.80 & 0.99 &15.945&173.4 &  BF  \\\\ & -0.32  &  0.00 & 1.06 & 0.82 & 0.00 &15.945& ---  &  L   \\\\ & -0.13  &  0.00 & 1.13 & 0.81 & 1.00 &15.935&173.4 &  BF  \\\\ & -0.19  &  0.00 & 1.19 & 0.76 & 0.85 &15.945& ---  &  L   \\\\ & -0.17  &  0.05 & 1.12 & 0.83 & 1.00 &15.935&173.4 &  BF  \\\\ & -0.20  &  0.05 & 1.15 & 0.78 & 0.54 &15.945& ---  &  L   \\\\ & -0.31  &  0.30 & 1.01 & 0.94 & 1.00 &15.955&173.6 &  BF  \\\\ & -0.30  &  0.30 & 1.00 & 0.91 & 0.00 &15.945& ---  &  L   \\\\ \\hline \\end{tabular} \\end{table}   \\begin{table} \\caption{Results of statistical analysis of Generalized Chaplygin Gas models with flat prior relaxed and with marginalization over ${\\cal M}$ performed on Knop Samples K3. Model parameter values are obtained from the marginal probability density functions. First row refer to no prior on $\\Omega_m$. The same analysis was repeated with fixed $\\Omega_m=0.0$, $\\Omega_m=0.05$ and $\\Omega_m=0.3$.} \\label{tab:5} \\begin{tabular}{@{}p{1.5cm}rrrrr} \\hline \\hline sample & $\\Omega_k$ & $\\Omega_m$ & $\\Omega_{Ch}$ &  $A_0$ & $\\alpha$ \\\\ \\hline K3  & $-0.60^{+0.38}_{-0.38}$ & $ 0.00^{+0.29}$ & $1.26^{+0.25}_{-0.39}$ & $0.89^{+0.11}_{-0.07}$ & $ 0.0^{+0.64}$ \\\\ & $ 0.10^{+0.37}_{-0.60}$ & $ 0.00$         & $0.90^{+0.59}_{-0.37}$ & $0.76^{+0.10}_{-0.07}$ & $ 0.0^{+0.66}$  \\\\ & $ 0.05^{+0.31}_{-0.58}$ & $ 0.05$         & $0.90^{+0.58}_{-0.31}$ & $0.78^{+0.10}_{-0.06}$ & $ 0.0^{+0.66}$ \\\\ & $-0.35^{+0.17}_{-0.40}$ & $ 0.30$         & $1.05^{+0.41}_{-0.17}$ & $0.88^{+0.09}_{-0.05}$ & $ 0.0^{+0.63}$  \\\\ Gold& $-0.32^{+0.25}_{-0.25}$ & $ 0.00^{+0.28}$ & $1.06^{+0.24}_{-0.22}$ & $0.82^{+0.13}_{-0.05}$ & $ 0.0^{+0.64}$ \\\\ & $-0.19^{+0.29}_{-0.28}$ & $ 0.00$         & $1.19^{+0.28}_{-0.29}$ & $0.76^{+0.03}_{-0.05}$ & $ 0.85^{+0.15}_{-0.52}$  \\\\ & $-0.20^{+0.28}_{-0.29}$ & $ 0.05$         & $1.15^{+0.29}_{-0.28}$ & $0.78^{+0.05}_{-0.06}$ & $ 0.54^{+0.36}_{-0.32}$ \\\\ & $-0.30^{+0.21}_{-0.23}$ & $ 0.30$         & $1.00^{+0.23}_{-0.21}$ & $0.91^{+0.04}_{-0.05}$ & $ 0.00^{+0.60}$  \\\\ \\hline \\end{tabular} \\end{table}   \\begin{figure} \\includegraphics[width=0.8\\textwidth]{f1.eps} \\caption{Residuals (in mag) between the Einstein-de Sitter model (zero line), the $\\Lambda$CDM model (upper curve) and the best-fitted Generalized Chaplygin Gas model  with $\\Omega_m= 0.3, \\Omega_{Ch}= 0.7$ (middle curve), sample K3.} \\label{fig:1} \\end{figure}  \\begin{figure} \\includegraphics[width=0.8\\textwidth]{f2.eps} \\caption{Levels of constant $\\chi^{2}$ on the plane $(A_0,\\alpha)$ for Generalized Chaplygin Gas model  with $\\Omega_m= 0.3, \\Omega_{Ch}= 0.7$, sample K3, marginalized over ${\\cal M}$. The figure shows preferred values of $A_0$ and $\\alpha$.} \\label{fig:2} \\end{figure}  \\begin{figure} \\includegraphics[width=0.8\\textwidth]{f3.eps} \\caption{Confidence  levels on the plane $(A_0,\\alpha)$ for Generalized Chaplygin Gas model  with $\\Omega_m= 0.3, \\Omega_{Ch}= 0.7$, sample K3, marginalized over ${\\cal M}$. The figure shows the ellipses of preferred values of $A_0$ and $\\alpha$.} \\label{fig:3} \\end{figure}  \\begin{figure} \\includegraphics[width=0.8\\textwidth]{f4.eps} \\caption{Residuals (in mag) between the Einstein-de Sitter model (zero line), the flat $\\Lambda$CDM model (upper curve) and the best-fitted Generalized Chaplygin Gas model with $\\Omega_m= 0.05, \\Omega_{Ch}= 0.95$ (middle curve), sample K3.} \\label{fig:4} \\end{figure}  \\begin{figure} \\includegraphics[width=0.8\\textwidth]{f5.eps} \\caption{Levels of constant $\\chi^{2}$ on the plane $(A_0,\\alpha)$ for Generalized Chaplygin Gas model  with $\\Omega_m= 0.05, \\Omega_{Ch}= 0.95$, sample K3, marginalized over ${\\cal M}$. The figure shows preferred values of $A_0$ and $\\alpha$.} \\label{fig:5} \\end{figure}  \\begin{figure} \\includegraphics[width=0.8\\textwidth]{f6.eps} \\caption{Confidence  levels on the plane $(A_0,\\alpha)$ for Generalized Chaplygin Gas model with $\\Omega_m= 0.05, \\Omega_{Ch}= 0.95$, sample K3, marginalized over ${\\cal M}$. The figure shows the ellipses of preferred values of $A_0$ and $\\alpha$.} \\label{fig:6} \\end{figure}  \\begin{figure} \\includegraphics[width=0.8\\textwidth]{f7.eps} \\caption{Residuals (in mag) between the Einstein-de Sitter model (zero line), the flat $\\Lambda$CDM model (upper curve) and the best-fitted Generalized Chaplygin Gas model  with $\\Omega_m= 0 , \\Omega_{Ch}= 1$ (middle curve), sample K3.} \\label{fig:7} \\end{figure} \\begin{figure} \\includegraphics[width=0.8\\textwidth]{f8.eps} \\caption{Levels of constant $\\chi^{2}$ on the plane $(A_0,\\alpha)$ for Generalized Chaplygin Gas model  with $\\Omega_m= 0, \\Omega_{Ch}= 1$, sample K3, marginalized over ${\\cal M}$. The figure shows preferred values of $A_0$ and $\\alpha$.} \\label{fig:8} \\end{figure}  \\begin{figure} \\includegraphics[width=0.8\\textwidth]{f9.eps} \\caption{Confidence  levels on the plane $(A_0,\\alpha)$ for Generalized Chaplygin Gas model with $\\Omega_m= 0, \\Omega_{Ch}= 1$, sample K3,  marginalized over ${\\cal M}$. The figure shows the ellipses of preferred values of $A_0$ and $\\alpha$.} \\label{fig:9} \\end{figure}  \\begin{figure} \\includegraphics[width=0.8\\textwidth]{f10.eps} \\caption{The density distribution (one dimensional PDF) for $\\Omega_{Ch}$ obtained from sample K3 by marginalization over remaining parameters of the model. We obtain the limit $\\Omega_{Ch} > 0.70$ at the confidence level $95.4 \\%$.} \\label{fig:10} \\end{figure} ",
        "conclusions": "It is apparent that Generalized Chaplygin Gas models have brighter supernovae at redshifts $z>1$. Indeed one can see on respective figures (Figs 1, 4, 7, 13) that systematic deviation from the baseline Einstein de Sitter model gets larger at higher redshifts. This prediction seems to be independent of analysed sample.  We obtained that the estimated value of $A_0$ is close to $0.8$ in all considered models with exepction the model class (1) ($\\Omega_m=0.3$) when $A_0>0.95$. Extending our analysis by relaxing the flat prior lead to the result that even though the best fitted values of $\\Omega_k$ are formally non-zero, yet they are close to the flat case. It should be viewed as an advantage of the GCG model since in similar analysis of $\\Lambda$CDM model in Perlmutter et al. (1999) high negative value of $\\Omega_{k}$ were found to be best fitted to the data and independent inspiration from CMBR and extragalactic astronomy has been invoked to fix the curvature problem. Another advantage of GCG model is that in a natural way we obtained the conclusion that matter (baryonic) component should be small what is in agreement with prediction from BBN (Big Bang Nucleosynthesis). Both estimations of $A_0$, $\\Omega_k$ and $\\Omega_m$ are independent of the sample used in our analysis.   Our results suggest that SNIa data support the Chaplygin gas (i.e. $\\alpha = 1$) scenario when the $\\chi^2$ best fitting procedure is used. The minimization procedure performed on Tonry/Barris  and Riess data gives also $\\alpha=1$ (only except the model with fixed $\\Omega_m=0.3$). However, the maximum likelihood fitting with Knop et al.'s sample prefers, quite unexpectedly, a small value of $\\alpha$ or even $\\alpha = 0$, i.e. the $\\Lambda$CDM scenario. Please note that small value of $\\alpha$ is in agreement with the results obtained from CMBR (de Bernardis et al. 2000 Benoit et al. 2003, Hinshaw et al. 2003, Bento, Bertolami \\& Sen 2002, Amendola et al. 2003) and with the recent analysis of Zhu (2004) who, using combined data of X-ray gas mass fraction of the galaxy cluster, FR IIB radiogalaxies and combined sample Perlmutter et al. (1998) and Riess et al. (1998, 2001), sugested that $\\alpha$ could be even less than 0. The results are dependent both on the sample chosen and on the prior knowledge of $\\cal M$ in which the Hubble constant and intrinsic luminosity of SNIa are entangled. Moreover the observed preference of $A_0$ values close to 1 means that the $\\alpha$ dependence becomes insignificant (see equation (\\ref{Hubble})). It is reflected on one dimensional PDFs for $\\alpha$ which turned out to be flat meaning that the power of the present supernovae data to discriminate between various Generalized Chaplygin Gas models (differing by $\\alpha$) is weak.  However, we argue that with future SNAP data it would be possible to differentiate between models with various value of $\\alpha$ parameter. Residual plots indicate the differences between $\\Lambda$CDM and Generalized Chaplygin Gas cosmologies at high redshifts. Therefore one can expect that future supernova experiments (e.g. SNAP) having access to higher redshifts will eventually resolve the issue whether the dark energy content of the Universe could be described as a Chaplygin gas. The discriminative power of forthcoming SNAP data has been illustrated on respective figures (Fig.26-29) obtained from the analysis on simulated SNAP data.  \\begin{figure} \\includegraphics[width=0.8\\textwidth]{f21.eps} \\caption{The density distribution (one dimensional PDF) for $\\Omega_{Ch}$ obtained from sample K3 by marginalization over remaining parameters of the model. We obtain the limit $\\Omega_{Ch} \\epsilon (0.61, 1.79)$ at the confidence level $95.4 \\%$. (Non-flat GCG model.)} \\label{fig:21} \\end{figure}  \\begin{figure} \\includegraphics[width=0.8\\textwidth]{f22.eps} \\caption{The density distribution (one dimensional PDF) for $A_0$ obtained from sample K3 by marginalization over remaining parameters of the model. We obtain the limit $A_0 \\epsilon (0.73,1)$ at the confidence level $95.4 \\%$.(Non-flat GCG model.)} \\label{fig:22} \\end{figure}  \\begin{figure} \\includegraphics[width=0.8\\textwidth]{f23.eps} \\caption{The density distribution (one dimensional PDF) for $\\alpha$ obtained from sample K3 by marginalization over remaining parameters of the model. We obtain the limit $\\alpha < 0.95$ at the confidence level $95.4 \\%$.(Non-flat GCG model.)} \\label{fig:23} \\end{figure}  \\begin{figure} \\includegraphics[width=0.8\\textwidth]{f24.eps} \\caption{The density distribution (one dimensional PDF) for $\\Omega_{k}$ obtained from sample K3 by marginalization over remaining parameters of the model. We obtain the limit $\\Omega_{k} \\epsilon (-1,0.23)$ at the confidence level $95.4 \\%$. (Non-flat GCG model.)} \\label{fig:24} \\end{figure}  \\begin{figure} \\includegraphics[width=0.8\\textwidth]{f25.eps} \\caption{The density distribution (one dimensional PDF) for $\\Omega_{m}$ obtained from sample K3 by marginalization over remaining parameters of the model. We obtain the limit $\\Omega_{m} \\epsilon( 0,0.53)$ at the confidence level $95.4 \\%$. (Non-flat GCG model.)} \\label{fig:25} \\end{figure}  \\begin{figure} \\includegraphics[width=0.8\\textwidth]{f26.eps} \\caption{Confidence  levels on the plane $(A_0,\\alpha)$ for sample X1 ($\\Lambda$CDM model) of simulated SNAP data, marginalized over $\\Omega_m$. The figure shows the ellipses of preferred values of $A_0$ and $\\alpha$.} \\label{fig:26} \\end{figure}  \\begin{figure} \\includegraphics[width=0.8\\textwidth]{f27.eps} \\caption{Confidence levels on the plane $(A_0,\\alpha)$ for sample X2 (Cardassian model) of simulated SNAP data, marginalized over $\\Omega_m$. The figure shows the ellipses of preferred values of $A_0$ and $\\alpha$.} \\label{fig:27} \\end{figure}   \\begin{figure} \\includegraphics[width=0.8\\textwidth]{f28.eps} \\caption{Confidence  levels on the plane $(A_0,\\alpha)$ for sample X3a (Generalized Chaplygin Gas model) of simulated SNAP data ($\\Omega_m=0$, $A_0=0.85$, $\\alpha=1.$), marginalized over $\\Omega_m$. The figure shows the ellipses of preferred values of $A_0$ and $\\alpha$.} \\label{fig:28} \\end{figure}  \\begin{figure} \\includegraphics[width=0.8\\textwidth]{f29.eps} \\caption{Confidence  levels on the plane $(A_0,\\alpha)$ for sample X3b (Generalized Chaplygin Gas model) of simulated SNAP data ($\\Omega_m=0$, $A_0=0.76$, $\\alpha=0.49$), marginalized over $\\Omega_m$. The figure shows the ellipses of preferred values of $A_0$ and $\\alpha$.} \\label{fig:29} \\end{figure}"
    },
    "0403/astro-ph0403149_arXiv.txt": {
        "abstract": "We studied the X-ray and optical absorption properties of 13 Gamma Ray Burst afterglows observed by BeppoSAX. We found that X-ray absorption in addition to the Galactic one along the line of sight is highly statistically significant in the two cases with the best statistics (probability $>99.9\\%$). In three other cases the presence of X-ray absorption is marginally significant (probability $\\sim 97\\%$). Measured rest frame absorbing column densities of hydrogen, $N_H$, range from 0.1 to 10.0 $\\times10^{22}$ cm$^{-2}$ (at 90\\% confidence level) assuming a solar metal abundance. X-ray absorption may be common, although the quality of present data does not allow us to reach a firm conclusion. We found that the rest frame column densities derived from XMM and Chrandra data as quoted in the literature are in good agreement with the BeppoSAX estimated rest frame $N_H$ range, supporting our result. For the same GRB afterglow sample we evaluated the rest frame visual extinction $A_{Vr}$. We fitted the optical-NIR afterglow photometry with a power law model corrected at short wavelengths by four different extinction curves. By comparing X-ray absorptions and optical extinction, we found that if a Galactic-like dust grain size distribution is assumed, a dust to gas ratio lower than the one observed in the Galaxy is required by the data. A dust to gas ratio $\\sim$ 1/10 than the Galactic one, as in the Small Magellanic Cloud (SMC) environment, has been tested using the SMC extinction curve, which produces good agreement between the best fit $N_H$ and $A_{Vr}$. We note, however, that the best fit $N_H$ values have been obtained by assuming solar metal abundances, while the metallicity of the SMC ISM is $\\sim$ 1/8 the solar one (Pei 1992). If such low metallicity were assumed, the best fit $N_H$ values would be higher by a factor of $\\sim7$, providing a significant increase of the $\\chi^2$. Alternative scenarios to explain simultaneously the optical and X-ray data involve dust with grain size distributions biased toward large grains. Possible mechanisms that can bring to such a grain size distribution are discussed. ",
        "introduction": "Soon after the $\\gamma$-ray flash, optical and X-ray afterglows of Gamma Ray Bursts (GRB) are among the brightest sources in the sky at cosmological redshifts.  More than thirty GRB redshifts have been measured to date and their distribution ranges from 0.168 (Greiner et al. 2003) to 4.5 (Andersen et al. 2000) with a median of z $\\sim 1.1$ (excluding GRB 980425 if at z=0.0085). Follow-up observations of the GRB localized by BeppoSAX, the IPN and by the RXTE and HETE2 satellites, show that tens of minutes after the GRB the optical afterglow can be as bright as R=14-16 mag; a few hours later it can still be as bright as R=17-19 mag. An exceptionally bright example is the case of GRB 030329 for which the optical afterglow reached R=12.7 mag at 1.5 hours from the GRB and it decreased down to R=19 mag after $\\sim 10$ days (e.g. Price et al. 2003; Stanek et al. 2003). The X-rays afterglow can be as bright as the Crab Nebula a few minutes after the GRB, while 5-8 hours later it can be as bright as a bright AGN, i.e. $\\approx$mCrabs (see Frontera et al. 2000, also Fiore et al. 2000 for an estimate of the GRB afterglows logN-logF in the range 0.5-2.0 keV, where fluxes are integrated from minutes up to hours after the GRB event). This opens up the perspective for gathering spectra of high quality of sources at cosmological redshifts, provided that the afterglow can be observed in such short time scales.  Spectral studies of GRB afterglows can provide crucial data on the environments in which GRB occurs. This can give us both important constraints on the nature of the GRB progenitor and on the interstellar matter (ISM) in the GRB host galaxies.  Candidate GRB progenitor include: 'collapsars' (Woosley 1993, Paczynski 1998; MacFadyen \\& Woosley 1999), 'supranovae' (Vietri \\& Stella 1998), mergers of two neutron stars or a neutron star and a black hole (e.g. Eichler et al. 1989; Paczynski 1990; Narayan, Paczynski \\& Piran 1992; Meszaros \\& Rees 1992).  Multiwavelength afterglow observations suggest the association of ``long GRB'' (for which the duration of the prompt event peaks at $\\sim20$s, e.g. Norris, Scargle \\& Bonnell 2000) progenitor with massive stars. Some indications are, among others: i) the GRB locations relative to their host galaxy (Bloom, Kulkarni \\& Djorgovski 2002a); ii) the detection of iron lines in the X-ray afterglow spectrum of five GRB (Piro et al. 1999, 2000; Yoshida et al. 1999; Antonelli et al. 2000; Amati et al. 2000); iii) the presence of late re-brightening in the light curves of some afterglows, interpreted as the evidence of an underlying supernova (e.g. Bloom et al. 2002b); iv) the presence of supernova features emerging in the optical afterglow spectra of GRB 030329 (Stanek et al. 2003, Hjorth et al. 2003; Kawabata et al. 2003) and of GRB 021211 (Della Valle et al. 2003). Because of the short lifetimes of massive stars, it is likely that their collapse happens close to their star forming region, that is, close to a dense and dusty environment (e.g. Fryer, Woosley \\& Hartman 1999; Perna \\& Belczynski 2002a). Therefore, one important element in favor of massive star progenitor would be the evidence of strong dust extinction. Indeed, the non detection of the $\\sim 60\\%$ of X-ray afterglows at optical wavelengths (dark GRB), favors this hypothesis. However, the optical spectra of GRB afterglows, do not generally show strong reddening (Simon et al. 2001; Galama \\& Wijers 2001). Moreover, Lazzati, Covino \\& Ghisellini (2002) show that even if dark GRB are located in the innermost regions of Galactic-like molecular clouds, in several cases the corresponding visual extinction is not enough to hide their optical afterglows. Therefore, a more accurate study of the GRB environment is needed.  Prompt observations of GRB afterglows offer a new and distinctive path for the study of the matter in the immediate surroundings of the GRB, r$\\sim$1-10~pc (Perna \\& Loeb 1998; B\\\"ottcher et al. 1999; Perna et al. 2002b; Perna \\& Lazzati 2002c; Fox et al. 2003) and in the GRB host galaxy (Ciardi \\& Loeb 2000; Fiore 2001; Bloom et al. 2002a; Savaglio, Fall \\& Fiore 2003) X-ray and optical-UV spectroscopy can tell us about gas density and ionization status, metal abundances, dust content and kinematics of the matter in the GRB environment.  This can be done by using both emission features and absorption features (see e.g. Kumar \\& Narayan 2003; B\\\"ottcher et al. 1999; Piro et al. 2000; Ghisellini et al. 2002). Although absorption features are more difficult to study than emission lines, they carry with them unbeatable information, because they probe matter along a single beam, i.e. along the line of sight to the background beacon.  This greatly reduces the complications related to the matter geometry and dynamics, which strongly affects emission line studies. Even when such emission or absorption features are not detected or are not resolved, X-ray spectra can provide useful information from measures of low energy cut-offs due to photoelectric absorption  Moreover, GRB afterglows can be used to study the ISM of their host galaxies. High redshifts galaxies have so far been studied mainly via `Lyman-break galaxies' (e.g. Steidel et al. 1999), and via galaxies which happen to lie along the line of sight to bright quasars, notably the `Damped Lyman-alpha' systems (DLA, e.g. Pettini et al. 1997, 1999). However, probably neither of these methods gives an undistorted view of the bulk of the typical high redshifts galaxy. Lyman-break galaxies are characterized by pronounced star-formation and their inferred chemical abundances may be related to these regions rather than being representative of typical high z galaxies (see e.g. Steidel et al. 1999).  DLA studies are likely to be biased against dusty systems, since these would hide the background quasars in color-based surveys. Furthermore, DLA are often identified with dwarf or Low Surface Brightness galaxies and therefore may not be representative of the full high z galaxy population.  On the other hand, GRB seem to occur well within the main body of their host galaxies, not in the outer halos (Bloom et al. 2002a), therefore, GRB afterglows can provide a powerful, independent tool to study the ISM of high redshifts galaxies (Castro et al. 2003, Mirabal et al. 2002, Fiore 2001; Savaglio et al. 2003). The latter authors compared metal column density in GRB host galaxies with those of DLAs. They find that while the column densities of iron and magnesium are similar or slightly higher than that in DLAs, the column of zinc are significantly higher.  Since iron and magnesium, unlike zinc, tend to be depleted in dust, this finding indicates that a significant fraction of these elements is in dust since not observed through absorption lines, thus GRB are probably probing denser and dustier regions, than DLAs.  We present in this paper the results of a systematic spectral analysis of 13 GRB afterglows observed by BeppoSAX in X-rays, and by several ground based telescopes in optical-UV, aimed at measuring and constraining the X-ray absorption along the line of sight as well as the extinction in the optical-UV bands.  The comparison of the results obtained in the two spectral band will provide constraints on the dust to gas ratio of the absorbing matter and on the dust properties, as seen from several hours to a few days from the GRB event.   The paper is organized as follows: \\S 2 shows the BeppoSAX X-ray afterglow selected sample; \\S 3 presents the X-ray data reduction and spectra extraction procedures, along with a discussion on the systematics affecting the column density measurements; \\S 4 presents the optical photometry of 9 GRB afterglows of our sample with a detected optical counterpart; results are presented and discussed in \\S 5 and \\S 6 respectively. ",
        "conclusions": ""
    },
    "0403/astro-ph0403463_arXiv.txt": {
        "abstract": "We examine the X-ray spectra and variability of the sample of X-ray sources with $L_{\\rm X} \\approx  10^{31} - 10^{33}$~\\ergsec\\ identified within the inner 9\\arcmin\\ of the Galaxy by \\citet{mun03}. Very few of the sources exhibit intra-day or inter-month variations. We find that the spectra of the point sources near the Galactic center are very hard between 2--8~keV, even after accounting for absorption. When modeled as power laws the median photon index is $\\Gamma = 0.7$, while when modeled as thermal plasma we can only obtain lower limits to the temperature of $kT > 8$~keV. The combined spectra of the point sources is similarly hard, with a photon index of $\\Gamma = 0.8$. Strong line emission is observed from low-ionization, He-like, and H-like Fe, both in the average spectra and in the brightest individual sources. The line ratios of the highly-ionized Fe in the average spectra are consistent with emission from a plasma in thermal equilibrium. This line emission is observed whether average spectra are examined as a function of the count rate from the source, or as a function of the hardness ratios of individual sources. This suggests that the hardness of the spectra may in fact be to due local absorption that partially-covers the X-ray emitting regions in the Galactic center systems. We suggest that most of these sources are intermediate polars, which (1) often exhibit hard spectra with prominent Fe lines, (2) rarely exhibit either flares on short time scales or changes in their mean X-ray flux on long time scales, and (3) are the most numerous hard X-ray sources with comparable luminosities in the Galaxy. ",
        "introduction": "}  Recent deep \\chandra\\ observations of the inner 9\\arcmin\\ around the super-massive black hole at the Galactic center have revealed over 2000 individual point-like X-ray sources \\citep{mun03}. The sources have luminosities between $10^{31}$ and $10^{33}$~\\ergsec \\citep[for a distance of 8 kpc to the Galactic center; see][]{mcn00}, and thus they probably represent some combination of young stellar objects, Wolf-Rayet and early O stars, interacting binaries (RS CVns), cataclysmic variables (CVs), young pulsars, and black holes and neutron stars accreting from binary companions (low- and high-mass X-ray binaries; LMXBs and HMXBs). However, the spectra of the Galactic center sources are very hard in the energy range of 2--8 keV. Spectra that are similarly hard have only been observed previously from magnetically accreting CVs (mCVs) and HMXB pulsars. Moreover, seven of the hard sources exhibit X-ray modulations with periods between 300~s and 4.5~h, which also suggests that they are magnetized white dwarfs or neutron stars \\citep{mun03c}. These basic observations are a good first step toward determining the natures of the point sources. However, if their natures can be determined conclusively, the large number of sources in the field would make it possible to study two important pieces of astrophysics: (1) the history of star formation at the Galactic center, and (2) the physics of X-ray production in accreting stellar remnants.  How stars form at the Galactic center is still a mystery, because the strong tidal forces and milliGauss magnetic fields there should prevent all but the most massive molecular clouds from collapsing. Nonetheless, it appears that star formation has occurred recently, because three massive stellar clusters younger than $10^{7}$ years old lie within $\\approx 30$~pc of the Galactic center: IRS 16, the Arches, and the Quintuplet \\citep{kra95,pau01,fig99}. However, it is still a matter of debate as to whether the star formation is continuous or episodic, and whether it occurs only in localized regions or is relatively uniform throughout the Galactic center. \\citet{fig03} addressed this question by modeling the evolution of the population of luminous infrared stars, and concluded that the star formation is probably continuous. The X-ray sources at the Galactic center could provide an additional, independent constraint on the star formation history there, because they should be dominated by accreting stellar remnants.  The size of the sample of X-ray sources --- an order of magnitude larger than the numbers of known LMXBs, HMXBs, and magnetic CVs --- also makes it a valuable database for studying the physics of X-ray production. Several outstanding questions could be addressed with the current data. If the sample contains large numbers of magnetic CVs, it could be used to determine the duty cycle of bright accretion states \\citep[e.g.,][]{gs88} and the fraction of such systems that exhibit hard spectral components \\citep[e.g.][]{ram03}. If there is a significant number of neutron star HMXBs, it may be possible to determine whether material accreted at rates far below the Eddington limit can penetrate the neutron star's magnetosphere and reach its surface \\citep{neg00,cam02,orl03}. Finally, the large sample of Galactic center X-ray sources would be useful for identifying systems with unusual properties. Previous hard X-ray surveys of the Galactic plane have identified several slowly-rotating accreting neutron stars and white dwarfs \\citep{kin98,tor99,oos99,sak00,sug00}, magnetic CVs with extremely strong emission lines from He-like Fe \\citep{mis96,ish98,ter99}, and accreting stellar remnants with high intrinsic absorption \\citep{pat03,mg03,wal03}. These systems could represent resting points for stellar remnants that have not been observed previously, and are therefore important for calculating the formation rate of such remnants in the Galaxy.  In this paper, we take a further step toward the above goals by using the properties of the X-ray emission from the point sources near the Galactic center \\citep{mun03} to constrain better their natures. In Sections 2.1--2.3, we examine the spectra of the point sources both individually and averaged together, in order to determine the temperatures of the emitting regions. In Section 2.4, we search for short-term variability, which is often seen from coronal X-ray sources, and long-term variability, which is common in some accreting X-ray sources. In Section~3, we compare the properties of the observed sources with those of known classes of X-ray source. Finally, in Section~4, we briefly explore the future prospects for definitively identifying the natures of these sources. ",
        "conclusions": "We have established that, on average, the X-ray sources detected in 626~ks of \\chandra\\ ACIS-I observations of the field around \\sgrastar\\ have hard, $\\Gamma < 1$ spectra with prominent emission from He-like Fe at 6.7~keV (Figure~\\ref{fig:psmod} and Table~\\ref{tab:psint}). They also generally do not vary by more than factors of a few on time scales of hours or months. The best candidates for these hard X-ray sources are intermediate polars, which represent the most luminous and spectrally hardest 5\\% of all CVs. Therefore, the Galactic center X-ray sources are likely to be only a sub-sample of a population of $\\sim 10^{4}$ CVs located near the Galactic center.  Although a single population of sources may dominate the image, there are certainly many classes of objects present in smaller numbers in the field. Determining the numbers of rare objects is particularly important. For instance, the numbers of massive Wolf-Rayet and O stars and faint neutron star high-mass X-ray binaries can constrain the recent rate of massive star formation near the Galactic center, while the numbers of LMXBS provide direct tests of the validity of unusual pathways for binary stellar evolution. For this reason, we are carrying out deep infrared observations of the Galactic center to identify counterparts to the X-ray sources. These observations will be useful for distinguishing CVs from, for example, HMXBs and WR/O stars. At at a distance of 8~kpc and with an extinction of $A_K \\approx 5$ \\citep{td03}, CVs should have $K$ magnitudes of 22--25, and therefore would be among the faintest detectable sources at the Galactic center \\citep{war95,hoa02}. In contrast, HMXBs and WR/O stars should have $K$ magnitudes brighter than 15 \\citep{zom90,weg94} and will be very easy to detect. Therefore, the prospects for identifying the natures of the Galactic center X-ray sources are promising."
    },
    "0403/astro-ph0403655_arXiv.txt": {
        "abstract": "We present an analysis of rotation measure (RM) fluctuations from the Test Region of the Southern Galactic Plane Survey (SGPS), along with emission measure (EM) fluctuations in the same field taken from the Southern H-Alpha Sky Survey Atlas. The structure function of RM fluctuations shows a relatively steep slope at small scales (1~--~5 arcmin), a break in slope to a flatter structure function at intermediate scales (5~--~60 arcmin), and a systematic variation of the strength of fluctuations as a function of position angle on the sky at the largest scales (60~--~200 arcmin). The structure function of EM fluctuations shows similar behavior, although the lower resolution of the data prevents detection of a possible break in the spectrum. We interpret the anisotropy in RM/EM structure on large scales as resulting from a large-scale gradient in electron density (and possibly magnetic field) across the region. The break in the slope of the RM structure function at scales of $\\sim5$~arcmin can be explained by contributions from two spatially distinct magneto-ionized screens, most likely in the Local and Carina spiral arms. The observed structure function then implies that the outer scale of RM fluctuations in these screens is $\\sim2$~pc. Such behavior is in striking contrast to the expectation that interstellar turbulence forms an unbroken spectrum from kpc down to AU scales.  We conclude that we have identified an additional source of enhanced turbulence, injected on scales of a few pc, possibly seen only in the Galactic plane.  The most likely source of such turbulence is individual H\\,{\\sc ii} regions from relatively low-mass stars, whose characteristic scale size is similar to the outer scale of turbulence inferred here. These sources may be the dominant source of density and velocity fluctuations in warm ionized gas in the Galactic plane. ",
        "introduction": "The evidence for the presence of turbulence in the interstellar medium (ISM) is overwhelming. Although some studies of the characteristics of this turbulence (viz.\\ the inner and outer scales, shape of the spectrum and power law spectral index) indicate the presence of standard incompressible Kolmogorov (1941) turbulence, many observations indicate different kinds of turbulence, drivers and environments. Armstrong, Rickett \\& Spangler (1995) compiled observations of (among others) interstellar scattering of pulsars and extragalactic sources, dispersion measures of pulsars and rotation measures (RM) of extragalactic sources in one power spectrum. The result was the so-called ``big power law in the sky'', a Kolmogorov-like power spectrum of electron density fluctuations over 12 orders of magnitude, from a fraction of an AU to kiloparsecs. Although the big power law in the sky is very well-determined on small scales (up to $\\sim 0.001$~pc), the extension towards larger scales is based only on two types of measurements: the RM of extragalactic sources, which gives an upper limit to the amount of structure in the electron density because the contribution of the magnetic field is unknown; and from velocity measurements of H\\,{\\sc i}, making assumptions about the connection between neutral and ionized material.  The dominant source for the large-scale energy input in the Galaxy is generally assumed to be supernova explosions (Spitzer 1978; Vollmer \\& Beckert 2002; Korpi et al.\\ 1999). Other possible sources are superbubbles, massive \\HII\\ regions and massive stellar winds \\citep{nf96}, Galactic fountains, chimneys, or gravitational scattering by transient spiral waves (see Sellwood \\& Balbus (1999) and references therein), gravitational instabilities in a shearing disk (Elmegreen, Elmegreen \\& Leitner 2003), or magneto-rotational instabilities \\citep{mk04}. The big power law in the sky certainly suggests the classical scenario of turbulent energy input on large scales, which cascades down to smaller scales until the energy is dissipated on the smallest scale \\citep{k41}.  However, there are indications of other types of turbulence in the ISM as well. \\cite{ms96} found a break in the structure function of RMs of extragalactic sources and emission measures (EM) of the warm ionized gas on scales of a few parsec, which they interpreted as a transition from three-dimensional to two-dimensional turbulence as one moves from small to large scales. Analytic theory (Goldreich \\& Sridhar 1995, 1997) and simulations (e.g.\\ Cho, Lazarian \\& Vishniac 2002; Maron \\& Goldreich 2001) show that magnetic fields can lead to anisotropic turbulence, which is predicted to exhibit the same power law spectral index as Kolmogorov turbulence. Furthermore, observations of RMs of extragalactic sources (Simonetti \\& Cordes 1986; Spangler \\& Reynolds 1990; Clegg \\etal\\ 1992) show higher amplitudes of structure in RM in the Galactic plane than out of the plane, suggesting the existence of an additional source of structure on parsec scales.  Results from interstellar scattering show a similar enhancement of structure in the inner Galaxy, but on much smaller scales (Rao \\& Ananthakrisnan 1984; Dennison \\etal\\ 1984; Anantharamaiah \\& Narayan 1988).  Structure in the ionized ISM can be studied by way of structure functions (SFs). Earlier determinations of SFs of RM using extragalactic sources or the diffuse synchrotron background yielded results that varied widely with scale and area in the sky. Flat structure functions (indicating no structure on the probed scales) found by Simonetti, Cordes \\& Spangler (1984) on scales larger than $4\\degr$ towards the Galactic pole were interpreted as showing structure intrinsic to the extragalactic sources with a negligible Galactic contribution. However, this cannot explain the flat SFs from low-latitude extragalactic sources found by \\citet{scs84} on scales $\\ga 4\\degr$ and by \\citet{ccs92} on scales $\\ga 1\\degr$. Furthermore, Sun \\& Han (2004) have recently found shallow SFs of RM in the Galactic plane, and a flat SF at the North Galactic Pole, also from extragalactic sources. Haverkorn, Katgert \\& de Bruyn (2003a) studied the structure function of RM from diffuse radio emission, and found very shallow slopes for two fields at intermediate latitudes.  In this paper we study the turbulent structure in ionized gas in the inner Galactic plane, by means of SFs of RM from the diffuse synchrotron background, and of SFs of EM from H$\\alpha$ emission, both in a region in the inner Galactic plane.  In Section~\\ref{s:obs} we present our radio polarimetric observations of the Galactic synchrotron background in the Galactic plane. Section~\\ref{s:sfdet} discusses the computation of the structure function, while in Section~\\ref{s:sfint} we interpret the structure function as arising from Faraday screens in two spiral arms along the line of sight, both exhibiting turbulent structure. In Section~\\ref{s:hii}, we discuss evidence for enhanced density fluctuations in the Galactic plane, and propose that the structure is dominated by discrete \\HII\\ regions. ",
        "conclusions": "\\label{s:con}  We have analyzed structure functions (SF) of RM and EM data in the SGPS Test Region. These show very consistent results: \\begin{itemize} \\item the SFs of both RM and EM, \\sfrm\\ and \\sfem, exhibit a linear slope in log-log space, with \\sfrm\\ showing a break at $\\sim4\\arcmin$ to a flatter slope at larger angular scales; \\item the slope at angular scales $\\ga4\\arcmin$ is shallow though non-zero in both \\sfrm\\ and \\sfem, with a SF spectral index $b \\approx 0.2$; \\item at scales larger than about a degree, the spectral index of \\sfrm\\ and \\sfem\\ is anisotropic, forming a quasi-sinusoidal dependence on position angle. \\end{itemize} The anisotropy in one-dimensional SFs at large scales is explained by a large-scale gradient in electron density, possibly accompanied by magnetic field, across the field of view due to a foreground structure. The break in \\sfrm\\ and shallowness of the slope in both \\sfrm\\ and \\sfem\\ at smaller scales can be explained by the superposition of two contributions to EM and RM of the Local and Carina spiral arms. Within such a model, we infer the outer scale of structure in the spiral arms to be about 2~pc. As Kolmogorov-like turbulence is observed in the ISM on scales much larger than a few pc, these results constitute evidence for an additional contribution to turbulent fluctuations in the Galactic plane. The inferred outer scale of fluctuations agrees with the size of Str\\\"omgren spheres around the most abundant late-type B~stars. Therefore, \\HII\\ regions could provide the dominant source of structure on pc scales in the Galactic spiral arms.  The suggestion that the break in \\sfrm\\ and the shallowness of both \\sfrm\\ and \\sfem\\ is due to two separate Faraday screens will be tested in a forthcoming paper by studying SFs in different regions of the sky. This can also shed light on the occurrence of any spatial variations of the turbulent spectrum. Furthermore, the combined RM and EM data should enable at least partial decoupling of magnetic field and electron density, and the amplitudes of the SF will yield information on this additional component of structure in the ionized ISM in the Galactic plane."
    },
    "0403/astro-ph0403525_arXiv.txt": {
        "abstract": "We present deep CCD photometry in the $VI$ passbands using the WIYN 3.5m telescope of a field located approximately 20' southeast of the center of M33; this field includes the region studied by Mould \\& Kristian in their 1986 paper. The color-magnitude diagram (CMD) extends to I$\\sim$25 and shows a prominent red giant branch (RGB), along with significant numbers of asymptotic giant branch and young main sequence stars. The red clump of core helium burning stars is also discernable near the limit of our CMD. The I-band apparent magnitude of the red giant branch tip implies  a distance modulus of $(m-M)_I = 24.77 \\pm 0.06$, which combined with an adopted reddening of $E(V-I)=0.06\\pm0.02$ yields an absolute modulus of $(m-M)_0 = 24.69 \\pm 0.07$ (867$\\pm$28 kpc) for M33. Over the range of deprojected radii covered by our field ($\\sim$8.5 to $\\sim$12.5 kpc), we find a significant age gradient with an upper limit of $\\sim$1 Gyr ($\\sim$0.25 Gyr/kpc). Comparison of the RGB photometry to empirical giant branch sequences for Galactic globulars allows us to use the dereddened color of these stars to construct a metallicity distribution function (MDF). The primary peak in the MDF is at a metallicity of $[Fe/H]$$\\sim$--1.0 with a tail to lower abundances. The peak does show radial variation with a slope of $\\Delta$[Fe/H]/$\\Delta$$R_{deproj}$ = --0.06 $\\pm$ 0.01 dex/kpc. This gradient is consistent with the variation seen in the inner {\\it disk} regions of M33. As such, we conclude that the vast majority of stars in this field belong to the disk of M33, not the halo as previously thought. ",
        "introduction": "In recent years M33 has been the target of many systematic studies.  From searches for variable stars to pin down one of the fundamental distance estimates in the Cepheid distance latter (e.g. Macri et al. 2001), to kinematical studies of the stellar populations (e.g. Chandar et al. 2002), to X-ray surveys (e.g. Haberl \\& Pietsch 2001).  However, in spite of M33 being the second closest spiral galaxy after M31, close enough for the brighter members of its stellar population to be resolved, the properties of its field halo stars have received comparatively little attention.  The first attempt to study the field halo stars in M33 using CCDs was that of Mould \\& Kristian (1986, hereafter MK86). In this classic paper, the authors observed an approximately 5x5 arcmin field (see Fig. 1) with the Palomar 5m telescope along with one of the first generation science-grade CCD detectors, a Texas Instruments 800 x 800 pixel array. Their color-magnitude diagram (CMD), based on aperture photometry of 215 stars transformed onto the Cousins system (Cousins 1976a,b), extends from above the first ascent red giant branch (RGB) tip down to $\\sim$1.5 mags below the tip (I$\\sim$22.5). From this diagram, they drew two primary conclusions. First, by comparing the M33 RGB with those of the Galactic globular clusters M92 and 47 Tuc, they found the mean metal abundance of the field stars to be $\\langle$[M/H]$\\rangle$ = --2.2 $\\pm$ 0.8. This result was rather surprising given the much higher metallicity ($\\langle$[M/H]$\\rangle > -0.8$) they found for the halo of M31. Second, from the bolometric magnitude of the RGB tip and an adopted reddening of $E(V-I) = 0.06$, they calculated an absolute distance modulus of 24.8 $\\pm$ 0.2 for M33.  Cuillandre, Lequeux, \\& Lionard (1999) used the UH8k CCD camera at the prime focus of the CFHT 3.6m to image a 28 x 28 arcmin field in a region of M33 that included the MK86 field. Their $(V,V-I)$ CMD reaches as faint as V $\\sim$ 25.5. Cuillandre et al. (1999) adopt an M33 distance modulus of $(m-M)_0 = 24.82$ and estimate the line-of-sight reddening ($E(V-I) = 0.08$) by averaging values based on the 21-cm line (Hartmann 1994) and from foreground stars (Johnson \\& Joner 1987). Based on these values, a comparison with the RGB sequences of Da Costa \\& Armandroff (1990) yields a mean metal abundance of [Fe/H] = --1.0 with a metallicity spread from $\\sim$--1.5 to $\\sim$--0.6 dex for the field halo stars in M33. This mean abundance is clearly much higher than the value advocated by MK86 and a few tenths of a dex higher than the peak abundance of Milky Way field halo stars (e.g. Norris 1994; Carney et al. 1996;).  Most recently, Davidge (2003) has used the Gemini North 8m telescope equipped with GMOS to image a field located $\\sim$9 kpc in projected distance from M33 ($\\sim$15.5 kpc deprojected distance), approximately twice as far out as the MK86 field. The g', r', i', z' CMDs of Davidge (2003) are based on a field of view of only 5x5 arcmin and are not extraordinarily deep, reaching some 2 magnitudes below the tip of the RGB, but they are notable as being the first such M33 data at these extreme galactocentric distances. Davidge's conclusions are twofold. First, based on the color of the upper RGB, Davidge (2003) determines a mean abundance of [Fe/H] = $-1.0 \\pm 0.3$ (random) $\\pm 0.3$ (systematic) for this region of M33. Second, the presence of a significant number of bright asymptotic giant branch stars suggests that an intermediate-age population exists outside of the `young star-forming' disk of M33. This latter result is consistent with the work of several investigators (e.g. Sarajedini et al. 1998;2000, Chandar et al. 2002) who have suggested that the halo star clusters in M33 are several Gyr younger than similar clusters in the Milky Way's halo.  Within this context, the present series of papers attempts to shed further light on the field halo stellar population(s) of M33. This first paper presents deep ground-based photometry of the MK86 field along with an analysis designed to probe the metallicity distribution function of this region of M33. We begin with a discussion of the observational material in the next section along with details of the photometric reduction procedure. Section 3 describes our artificial star experiments and how they are used to gauge the photometric completeness and errors. We compare our photometry to that of MK86 in Sec. 4. The analysis portion of the paper begins in Sec. 5 in which we use the luminosity function of the RGB to estimate the distance to M33. In Sec. 6, the radial variations of the stellar populations present in our field are analyzed. Section 7 presents the metallicity distribution function and what it implies for this region of M33. Finally, our results are summarized in Sec. 8. ",
        "conclusions": "In this work, we present deep $VI$ photometry of a field in M33, that includes the region studied by Mould \\& Kristian(1986), observed with the WIYN 3.5m telescope. Based on the CMD, which reaches past the helium burning red clump, we draw the following conclusions.  \\noindent 1) The I-band apparent magnitude of the TRGB is measured to be $\\rm I_{\\rm TRGB} = 20.75 \\pm 0.04$, where the error includes the measurement error (0.02 mag) along with random (0.02 mag) and systematic (0.03 mag) errors in the photometry.  \\noindent 2) The I-band apparent magnitude of the TRGB implies  a distance modulus of $(m-M)_I = 24.77 \\pm 0.06$, which combined with an adopted reddening of $E(V-I)=0.06\\pm0.02$ yields an absolute modulus of $(m-M)_0 = 24.69 \\pm 0.07$ (867$\\pm$28 kpc) for M33.  \\noindent 3) By examining the spatial properties of stars in various parts of the CMD, we find that, over the range of deprojected radii covered by our field ($\\sim$8.5 to $\\sim$12.5 kpc), there is a significant age gradient with an upper limit of $\\sim$1 Gyr (i.e., $\\sim$0.25 Gyr/kpc).  \\noindent 4) Using the empirical RGB grid constructed by Saviane et al. (2000), we have converted the dereddened color of each RGB star to a metallicity. The resultant metallicity distribution function (MDF) displays a primary peak at a metallicity of $[Fe/H]$$\\sim$--1.0 with a tail to lower abundances. The peak shows a radial variation with a slope of $\\Delta$[Fe/H]/$\\Delta$$R_{deproj}$ = --0.06 $\\pm$ 0.01 dex/kpc. This gradient is consistent with the variation seen in the inner {\\it disk} regions of M33. Therefore, we conclude that the vast majority of stars in this field belong to the disk of M33, not the halo as previously thought."
    },
    "0403/astro-ph0403239_arXiv.txt": {
        "abstract": "State-of-the-art NLTE model-atmosphere codes have arrived at a high level of numerical sophistication and are now useful tools to analyze high-quality spectra from the infrared to the X-ray wavelength range. The capacity of current computers permit calculations which include line spectra from all elements from hydrogen up to the iron group. The lack of reliable atomic data has become a critical problem for further progress.  We summarize available sources of atomic data, and discuss how these are implemented in the T\\\"ubingen Model-Atmosphere Package \\linebreak \\textsc{tmap}. We describe our Iron Opacity Interface {\\sc IrOnIc} which is used to calculate opacities of iron-group elements from Kurucz's and the Opacity Project's data.  We propose general use of the T\\\"ubingen Model-Atom Database \\linebreak {\\sc tmad}, which would allow an easy exchange of ready-to-use model atoms between all model-atmosphere groups. The comparison of model-atmo\\-sphere calculations would then be much easier, and would save a great deal of manpower that is presently consumed preparing suitable model atoms for spectral analyses. ",
        "introduction": "In the early 80's of the last century, the implementation of approximate lambda-operators (ALO, leading to ``accelerated lambda iteration'', ALI) in the NLTE model-atmosphere codes at Kiel by Werner \\& Husfeld (1985) and Werner (1986) provided an efficient method to calculate synthetic stellar spectra of hot stars. Together with the access to the fourth {\\sc cray} computer (a {\\sc cray-1\\,m}) in Germany (at the Konrad-Zuse-Zentrum f\\\"ur Informationstechnik Berlin), which was installed in February 1984, powerful tools for spectral analysis of hot star spectra were developed. At the end of 1987, the Rechenzentrum der Universit\\\"at Kiel installed a {\\sc cray\\,x-mp}. Access to this machine (in the framework of the Norddeutscher Vektorrechner-Verbund) and the following, even more powerful, {\\sc cray} computers, made all our efforts possible.  The first of our NLTE codes, \\textsc{hymoc} ({\\sc HYd}rogen {\\sc MO}del-atmosphere {\\sc C}ode) (Werner 1986), was able to calculate only pure-hydrogen model atmospheres. However, inasmuch as the atomic data for hydrogen are well known, it was an ideal tool to investigate the numerical approximations, and limitations to the size of model atoms, used in earlier calculations, which were necessary before the ALI method was developed (Rauch \\& Werner 1988).  The next code \\textsc{pro2} ({\\sc pro}gram no.\\,2) (Werner 1988; Werner \\& Dreizler 1999) is much more flexible, and is able to take into account all elements up to the iron group (Dreizler \\& Werner 1993; Rauch 1997). It has been used successfully for the analysis of hot stars (e.g.~Rauch \\& Werner 1991; Rauch 1993; Rauch 2000). Our NLTE group moved from Bamberg (1993) and Potsdam (1995) to T\\\"ubingen (since 1996) where \\textsc{tmap}, the state-of-the-art {\\sc T}\\\"ubingen {\\sc M}odel {\\sc A}tmosphere {\\sc P}ackage was created. With the newly developed \\textsc{IrOnIc} code (\\S 4), it was possible to calculate an extended grid of realistic stellar fluxes from models which take into account the opacities from H -- Ni (Rauch \\& Deetjen 2001). These models are plane-parallel, in hydrostatic and radiative equilibrium, have 350 atomic levels which are treated in NLTE, about 1,000 individual lines from H - Ca, and millions of lines from the iron-group elements.  The state of the field of spectral-analysis of hot stars has completely changed within the last two decades. At the beginning of the 80's, the main obstacles were insufficient numerical methods and computational capacities. Rauch \\& Werner (1991) have shown the enormous progress which came with the ALI method, with examples of very detailed H + He + C + N models in contrast to ``classical'' H + He models. At present, the lack of reliable atomic data for metals, line-broadening tables, etc.~set undesirable limits to highly-developed NLTE codes. This lack often hampers an adequate analysis e.g.~of high-resolution UV spectra provided by the STIS (Space Telescope Imaging Spectrograph) aboard the Hubble Space Telescope (HST) or by the Far Ultraviolet Spectroscopic Explorer (FUSE). ",
        "conclusions": "Present NLTE model-atmosphere codes have reached a very high level of sophistication.  Now, strong efforts to achieve adequate atomic data must be continued in order to be able to analyze reliably the high-quality spectra which are already available from the infrared to the X-ray range. {\\it Everyone who is calculating stellar atmospheres should be aware of the important role that atomic data plays, because one thing is for sure -- even if you use a perfect code, if you put rubbish in, you will get rubbish out.}"
    },
    "0403/astro-ph0403713_arXiv.txt": {
        "abstract": "High ambient interstellar pressure is suggested as a possible factor to explain the ubiquitous observed growth-rate discrepancy for supernova-driven superbubbles and stellar wind bubbles. Pressures of $P/k\\sim 10^5\\ \\ccK$ are plausible for regions with high star formation rates, and these values are intermediate between the estimated Galactic mid-plane pressure and those observed in starburst galaxies.  High-pressure components also are commonly seen in Galactic ISM localizations.  We demonstrate the sensitivity of shell growth to the ambient pressure, and suggest that superbubbles ultimately might serve as ISM barometers. ",
        "introduction": "Mechanical feedback from supernovae (SNe) and stellar winds is an important driver of galactic evolutionary processes.  It affects the phase balance and physical conditions of the interstellar medium (ISM), which in turn determine star-forming conditions, galactic chemical evolution, and properties of the intergalactic medium. The standard paradigm for mechanical feedback is based on the model for adiabatic evolution of the shells and superbubbles (e.g., Pikel'ner 1968; Weaver {\\etal}1977; Mac Low \\& McCray 1988) that are the direct consequence of SN and stellar wind action.  While this model for pressure-driven superbubbles is broadly consistent with observations spanning scales from individual stellar wind bubbles to galactic superwinds, nagging problems persist in comparisons with observations (e.g., see Oey 2004 for a review).  Specifics regarding the energy budgets, fate of the shock-heated $10^6$ K gas, and later-stage evolution are lacking and have profound consequences for galactic evolution.  One problem that is empirically well-established is the result that most superbubbles apparently grow more slowly than expected.  This has been observed in individual stellar wind bubbles such as Wolf-Rayet nebulae (Treffers \\& Chu 1982; Garc\\'\\i a-Segura \\& Mac Low 1995; Drissen {\\etal}1995), as well as in superbubbles powered by OB associations (e.g., Oey 1996$a$; Oey \\& Smedley 1998; Brown {\\etal}1995; Saken {\\etal}1992).  This growth-rate discrepancy has been identified in young, nebular shell systems, in which the parent OB association is still present; thus the input mechanical power is well-constrained. The discrepancy is seen both in objects that show no evidence of previous supernova activity, and in ones where one or two supernovae have already exploded (Oey 1996$a$; hereafter O96).  For constant  input mechanical power $L$ and ambient number density $n$, the evolution of the shell radius is given by (e.g., Castor, McCray, \\& Weaver 1975), \\begin{equation}\\label{eqR} R = 68.9\\ (L_{38}/n)^{1/5}\\ t_6^{3/5}\\ \\ \\rm pc \\quad , \\end{equation} where $L_{38}$ is $L$ in units of $10^{38}\\ \\ergs$, and $t_6$ is age of the bubble in Myr.  The shell expansion velocity $v$ is the time derivative of equation~\\ref{eqR}.  One possible solution to the growth-rate discrepancy suggests that the input parameter $L/n$ is systematically overestimated.  For eight nebular superbubbles with well-constrained $R$, $v$, $L$, and $t$, Oey and collaborators (O96; Oey \\& Massey 1995; Oey \\& Smedley 1998) showed that $L/n$ would need to be reduced by a factor of several, perhaps up to an order of magnitude, to reconcile the observations with prediction.  Since stellar wind power $L$ is sensitive to the stellar mass, a substantial uncertainty in $L$ is not unreasonable.  As shown by multi-wavelength observations of three of the superbubbles (Oey {\\etal}2002), the multi-phase ambient ISM also renders $n$ similarly uncertain. However, the implication of a {\\it systematic} growth-rate discrepancy remains difficult to explain.  Another favorite candidate to solve the problem is cooling of the hot interior whose pressure drives the shell growth.  If this scenario is correct, it implies a departure from the adiabatic evolution.  The mass within the hot bubble interior is dominated by material evaporated from the cool shell walls, and could be supplemented by additional material evaporated and ablated from small clouds that are overrun by the expanding shocks (Cowie \\& McKee 1977; McKee {\\etal}1984; Arthur \\& Henney 1996).  This enhanced interior density would facilitate radiative cooling.  Silich {\\etal}(2001) and Silich \\& Oey (2002) also show that the enhanced metallicity caused by SN explosions and stellar products can further facilitate the cooling, especially for low-metallicity objects.  However, increased X-ray luminosities that are expected from enhanced cooling thus far have not been observed (Chu {\\etal}2003; Chu {\\etal}1995).  The superbubbles studied by Oey and collaborators are all located in the Large Magellanic Cloud (LMC).  For these objects, Silich \\& Franco (1999) suggested that the ambient environment and viewing geometry conspire to yield misleading observed shell dynamics.  They suggest that the superbubbles are more extended perpendicular to the galaxy's plane, as would be expected in the plane-stratified density distribution of disk galaxies (but see also Maciejewski \\& Cox 1999). The elongation of the shells would not be apparent because of the LMC's almost face-on orientation.  While this is an attractive suggestion for the LMC objects, it does not explain the growth-rate discrepancy seen in Galactic (e.g., Brown {\\etal}1995; Saken {\\etal}1992) and M33 (Hunter {\\etal}1995) objects.  Nevertheless, it is apparent that the ambient environment plays a crucial role in the superbubble growth and evolution.  In addition to the work of Silich \\& Franco (1999), other studies have shown that the shell dynamics are sensitive to the ambient density structure (e.g., Oey \\& Smedley 1998; Mac Low {\\etal}1998).  Multi-wavelength observations also show that the ambient multiphase gas distribution is difficult to constrain without direct such observations (Oey {\\etal}2002). ",
        "conclusions": "Our models clearly show that increasing the ambient interstellar pressure by an order of magnitude, from $P_e/k = 1\\times 10^4$ to $1\\times 10^5\\ \\ccK$, can impede the shell growth to a degree that could fully account for the observed growth-rate discrepancy. In \\S 2, we presented arguments that such high interstellar pressures could exist, especially based on the dependence of $P_e$ on star-formation rate.  While other factors mentioned in \\S 1, namely, overestimated $L/n$, elevated radiative cooling, and viewing geometry, could  all be additional factors that contribute to the growth-rate discrepancy, we note that the multi-phase gas morphology is more consistent with high interstellar pressure dominating this effect.  Finally, as noted by Oey \\& Clarke (1997), the assumed global value of $P_e$ plays a critical role in determining the characteristic final sizes of old, SN-dominated superbubbles, and hence, the superbubble size distribution, which is dominated by pressure-confined shells. This, in turn, determines the interstellar porosity and filling factor of the hot, ionized medium in star-forming galaxies.  With adequate clarification in the superbubble evolution process and input parameters, the superbubble sizes, kinematics, and morphologies could potentially provide barometers for the interstellar pressure.  These diagnostics could be especially useful in other galaxies, which have fewer available empirical pressure indicators than the Milky Way."
    },
    "0403/hep-ph0403004_arXiv.txt": {
        "abstract": "\\noindent Assuming that cosmological dark matter consists of weakly interacting massive particles, we use the recent precise measurement of cosmological parameters to predict the guaranteed rates of production of such particles in association with photons at electron-positron colliders. Our approach is based on general physical principles such as detailed balancing and soft/collinear factorization. It leads to predictions that are valid across a broad range of models containing WIMPs, including supersymmetry, universal extra dimensions, and many others. We also discuss the discovery prospects for the predicted experimental signatures. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403707_arXiv.txt": {
        "abstract": "{It is commonly adopted that X-rays from O stars are produced deep inside the stellar wind, and transported outwards through the bulk of the expanding matter which attenuates the radiation and affects the shape of emission line profiles. The ability of the X-ray observatories Chandra and XMM-Newton to resolve these lines spectroscopically provided a stringent test for the theory of the X-ray production. It turned out that none of the existing models was able to fit the observations consistently. The possible caveat of these models was the underlying assumption of a smooth stellar wind. Motivated by the various evidences that the stellar winds are in fact structured, we present a 2-D numerical model of a stochastic, inhomogeneous wind. Small parcels of hot, X-ray emitting gas are permeated with cool, absorbing wind material which is compressed into thin shell fragments. Wind fragmentation alters the radiative transfer drastically, compared to homogeneous models of the same mass-loss rate. X-rays produced deep inside the wind, which would be totally absorbed in a homogeneous flow, can effectively escape from a fragmented wind. The wind absorption becomes wavelength independent if the individual fragments are optically thick. The X-ray line profiles are flat-topped in the blue part and decline steeply in the red part for the winds with short acceleration zone. For the winds where the acceleration extends over significant distances, the lines can appear nearly symmetric and only slightly blueshifted, in contrast to the skewed, triangular line profiles typically obtained from homogeneous wind models of high optical depth. We show that profiles from a fragmented wind model can reproduce the observed line profiles from $\\zeta$\\,Orionis. The present numerical modeling confirms the results from a previous study, where we derived analytical formulae from a statistical treatment. ",
        "introduction": "Hot massive stars possess strong stellar winds, as discovered by the advent of UV spectroscopy (Morton \\cite{Morton67}). One decade later it was detected that these stars emit X-rays (Harnden et al.\\ \\cite{Harnden79}, Seward et al.\\ \\cite{Seward79}). The origin of these X-rays is debated controversially. One of the first suggestions was the existence of a hot corona close to the stellar photosphere. However, the surrounding stellar wind should imprint strong K-shell absorption edges on the X-ray spectrum, which were not observed (Cassinelli \\& Swank \\cite{CS83}). The base-coronal model was finally ruled out because of the lack of characteristic coronal emission lines (Baade \\& Lucy \\cite{Lucy82}). Therefore it became commonly adopted that X-rays from hot stars are produced within their stellar winds, although most of the wind gas is obviously ``cool''.  The driving mechanism for the mass-loss from OB stars has been identified with radiation pressure on spectral lines. A corresponding theory was developed by Castor et al.\\ (\\cite{CAK75}, ``CAK''). After a number of improvements, this theory seems to predict the observed mass-loss rates and wind velocities correctly (Pauldrach et al.\\ \\cite{Pauldrach86}). However, it was pointed out early (Lucy \\& Solomon \\cite{LucySolomon70}), and later further investigated (Carlberg \\cite{Carlberg80}; Owocki \\& Rybicki \\cite{OR84}), that the stationary solution for a line-driven wind is unstable; small perturbation should grow quickly and form shocks. This ``de-shadowing instability'' is implicitly suppressed in the CAK theory.  The most detailed hydrodynamic modeling of this line driven instability was presented by Feldmeier et al.\\ (\\cite{AF97}). These calculations show how initially small perturbations grow and form strong shocks that emit X-rays. This ``wind shock model'' was principally able to reproduce the X-ray flux from the O supergiant $\\zeta$~Ori\\,A, which had been observed at that time only at low spectral resolution.  The ultimate test for models of X-ray production in stellar winds became possible with the launch of the X-ray observatories Chandra and XMM-Newton. Their spectrographs allow for the first time to resolve X-ray emission line profiles. Surprisingly, the shape of the X-ray line profiles turned out to be strikingly different among the few individual O stars observed so far:  \\begin{description} \\item {$\\theta^1$~Ori\\,C (Schulz et al.\\ \\cite{Schulz00}) and $\\zeta$~Ori\\,A (Waldron \\& Cassinelli \\cite{WaldCass01})} show profiles which appear wind-broadened, symmetric and unshifted with respect to the center wavelength;  \\item {$\\delta$~Ori\\,A (Miller et al.\\ \\cite{Milleretal02})} shows profiles which are narrower, symmetric, and also centered at the laboratory wavelength;  \\item {$\\zeta$~Pup (Cassinelli et al.\\ \\cite{Cass01}; Kahn et al.\\ \\cite{Kahnetal01})} shows strongly broadened profiles which are blueshifted against the rest wavelength but otherwise appear to be symmetric.  \\end{description}  If X-rays are produced deep inside the stellar wind, they have to propagate through the absorbing cool wind before they can emerge towards the observer. The opacity at X-rays energies is much larger than in the visual and UV, due to the K-shell absorption of abundant metal ions and the He\\,{\\sc ii} edge.  Red-shifted photons of an emission line originate in the back hemisphere of the wind, and suffer much stronger absorption than photons in the blueshifted part of the line profile. Hence, as soon as the absorbing wind has a significant optical depth, the line profiles should be blueshifted and skewed, which means their shape resembles that of a triangle, with the flux maximum at maximum blueshift (MacFarlane et al.\\ \\cite{MacFarlane91}). None of the observed stars meets this prediction.  Kramer et al.\\ (\\cite{KCO03}) were able to fit eight X-ray lines observed in $\\zeta$\\,Puppis with a homogeneous wind model. However, these fits could only be achieved by assuming much less absorption in the wind than expected from the generally adopted mass-loss rate of this star.  Moreover, Kramer et al.\\ (\\cite{KCO03}) state that there is not a strong trend of absorption with wavelength, in contrast to what is expected from the energy dependence of the mass absorption coefficient. Thus even for the prototype single O star $\\zeta$~Puppis the wind shock model fails to reproduce observed X-ray emission lines consistently. At this point the question arose whether this model is adequate.  We emphasize that the whole interpretation of X-ray spectra had been based so far on over-simplified models, assuming a {\\em homogeneous} distribution of the absorbing material. By contrast, there is strong empirical evidence that stellar winds are in fact strongly clumped (e.g.\\ Hamann \\& Koesterke \\cite{HK98}, Eversberg et al.\\ \\cite{Ev98}, Puls et al.\\ \\cite{Puls03}).  This supports hydrodynamic models (Owocki et al.\\ \\cite{OCR88}, Feldmeier \\cite{AF95}) which show that most of the material in a radiatively driven wind is compressed into a series of dense shells by the instability. As the hydrodynamic models are one-dimensional, they cannot tell anything about the lateral structure of these wind inhomogeneities. The de-shadowing instability acts only in the radial direction; on the other hand, there is no obvious mechanism which could synchronize the wind in different directions. The observed variability in wind lines is not very pronounced, suggesting that the dense shells break up into a large number of fragments.  Wind inhomogeneity alters the radiative transfer significantly. We have studied the effects of wind fragmentation on the X-ray line formation in Feldmeier et al.\\ (\\cite{PaperI}, Paper\\,I), using a statistical approach that holds in the limit of infinitely many fragments. We showed analytically that the continuum wind opacity is greatly reduced in a structured wind, compared to a smooth wind with the same mass-loss rate. The line profiles we obtained are broad, blueshifted and flat-topped, i.e. symmetric, and thus of promising similarity to those observed in $\\zeta$~Pup.  In the present paper we construct a stochastic wind model using a numerical approach. This allows to drop a number of idealizations in favor of a more realistic description.  \\begin{description} \\item {\\it Statistical approach.} While Paper\\,I relied on the statistical limit of very many fragments, we allow now for an arbitrary scale of fragmentation.  \\item {\\it Wind geometry.} We study models with different spatial distributions of fragments, both in the  radial and lateral direction.  \\item {\\it Wind velocity law.} We abandon the restriction to a constant wind velocity and allow for so-called $\\beta$-velocity laws. \\end{description}  Last not least, another purpose of the present paper is to better understand some surprising results obtained in Paper\\,I. Despite of our focus on X-ray emission line formation, we want to emphasize their general importance for radiative transfer in inhomogeneous media.  In the next section we will introduce the stochastic wind model. Section\\,3 provides the formalism for solving the radiative transfer in a fragmented wind, as implemented in our numerical modeling. The basic effects of structured absorption on line profiles are illustrated in Sect.\\,4 using simplified examples. In Sect.\\,5 we introduce the concept of an effective opacity, and compare the line profiles obtained numerically from the stochastic wind model with the results from an analytical, statistical treatment. In Sect.\\,6 we evaluate line profiles using realistic assumptions and explore the parameter space. Conclusions are drawn in the final Sect.\\,7. A forthcoming paper will be devoted to detailed fits of observations. ",
        "conclusions": "A main motivation for our fragmented wind model is to explain the observed X-ray line profiles from O stars, which cannot be reproduced so far. We may check now whether we came closer to this goal. Figure\\,\\ref{fig:zeta_Ori} shows the profile of the Ne\\,{\\sc x} line from $\\zeta$\\,Orionis, as observed with the Chandra satellite (histogram). The smooth line is a tentative comparison with one of our calculations. From the known mass-loss rate of the star, the continuum at the frequency of this line may be optically thick (Waldron et al.\\ \\cite{WaldCass01}).  Therefore we select a suitable profile (with $\\beta = 2$) from the set in Fig.\\,\\ref{fig:bvthick} with $\\tau_\\ast = 10$. The dimensionless frequency is converted into wavelength according to the terminal wind velocity of $\\zeta$\\,Orionis (2100\\,km\\,s$^{-1}$), and the profile is convolved with a Gaussian of 0.023\\,\\AA\\ FWHM to account for instrumental broadening. The good agreement just demonstrates that our fragmented wind model can reproduce emergent line profiles of the observed shape, even if $\\tau_\\ast$ is large. For a systematic fit, all available line profiles should be considered in order to adjust the parameters of the model consistently. This will be subject of a forthcoming paper.  There are two aspects in the problem of fitting the observed line profiles. One is to reproduce the line shape, another is to provide for the observed line flux. Ignace \\& Gayley (\\cite{Rico02}) considered profile shapes for optically thick emission lines. They showed that these lines are blueshifted, and slightly narrower than optically thin lines. The lines have a universal shape which is nearly symmetric and insensitive to the level of continuous absorption. Nevertheless, the flux in the line scales inversely with optical depth of the cool wind, and therefore the problem of missing opacity is still unresolved in their approach. The same problem is highlighted in Kramer et al.\\ (\\cite{KCO03}), which up-to-day is the most consistent attempt to fit the observed lines. Our study shows that incorporation of wind fragmentation can resolve the controversy.  A burning question is, why stars with existing high-resolution X-ray spectra show qualitatively different profile shapes, as listed in the introduction. At present we suggest that only $\\zeta$\\,Pup is a safe prototype for an isolated strong stellar wind.  Indeed, speckle interferometry led to the discovery that $\\theta^1$\\,Ori\\,C and $\\zeta$\\,Ori\\,A, are binary systems (Hummel et al.\\ \\cite{Hummel00}, Weigelt et al.\\ \\cite{Weigelt99}), and $\\delta$\\,Ori\\,A is a well-known multiple system (see e.g. Miller et al.\\ \\cite{Milleretal02}). Moreover, $\\theta^1$\\,Ori\\,C has a strong magnetic field in excess of 1\\,kG (Donati et al.\\ \\cite{Don02}). The physical processes leading to the emission of X-rays can differ between objects. Nevertheless, we claim that the effects of wind fragmentation ought to be included in the modeling of emerging line profiles independent of the particular emission mechanism.  \\begin{figure}[bth] \\centering \\epsfxsize=\\columnwidth \\mbox{\\epsffile{zoriNeX.ps}}  \\caption{Profile of the Ne\\,{\\sc x} line from $\\zeta$\\,Orionis, as observed with the Chandra satellite (histogram). The smooth line is a tentative comparison with one of our calculations, taking the profile for $\\tau_\\ast = 10$ and $\\beta = 2$ from Fig.\\,\\ref{fig:bvthick}. The theoretical profile is scaled to the terminal wind velocity of $\\zeta$ Orionis (2100\\,km\\,s$^{-1}$) and convolved with the instrumental profile, a Gaussian of 0.023\\,\\AA\\ FWHM. The good fit demonstrates that our fragmented wind model can reproduce emergent line profiles of the observed shape.}  \\label{fig:zeta_Ori} \\end{figure}  The scenario chosen as a framework of the model presented here is that of Feldmeier et al.\\ \\cite{AF97}, where X-ray are produced in the collision of fast cloudlets with dense shells. The prominent feature of this hydrodynamic simulation is that the X-ray emitting plasma is always located at the starward face of the cool absorbing fragments. It was already shown in Paper\\,I, and is confirmed here, that this configuration leads to strong depletion of the central part of the line, an effect which is certainly not observed.  The following conclusions can be drawn  from our numerical modeling of the X-ray line emission from inhomogeneous, fragmented stellar wind:  \\begin{enumerate}  \\item Fragmentation drastically reduces the effective opacity of the wind. Therefore X-rays produced deep inside the wind can effectively escape, which would be totally absorbed in a homogeneous wind of the same mass-loss rate.  \\item Absorption in a fragmented wind becomes effectively independent of the mass absorption coefficient and therefore of the wavelength, if the fragments are optically thick.  \\item The line profiles from the fragmented wind model exhibit a variety of shapes, ranging from broad, blueshifted and flat-topped to narrow and nearly symmetric, and are thus of promising similarity to observations.  \\item The possibility of a flat-topped, i.e.\\ almost symmetric blue part of the line profile is due to the model assumption that the fragments are flat, have small thickness, and are aligned perpendicular to the radial flow direction.  \\item The effect of fragmentation is significant when the individual fragments are optically thick. Therefore the effect depends on the average number of fragments per radial direction. This parameter can be empirically restricted by a detailed line fit and has interesting implications for the theoretical understanding of the wind hydrodynamics.  \\item The lateral size of the wind fragments has no influence on the emergent lines.  \\item If the fragments are strictly confined to a pattern of radial cones while moving outwards, which ensures strict mass conservation in each radial direction, the resulting line profiles differ only little from the random fragment model where the absorbing shell fragments are randomly distributed in both radial and angular coordinates.  \\item The numerical modeling featuring stochastically arranged emitting parcels of gas and absorbing fragments confirms the results of Paper\\,I, where analytical formulae have been derived from a statistical treatment.  \\end{enumerate}"
    },
    "0403/astro-ph0403531_arXiv.txt": {
        "abstract": "{Using  spectra of normal emission line galaxies from the First Data Release of the Sloan Digital Sky Survey (SDSS) we have investigated the relations between the  extinction $C({\\rm H}\\beta)$ as derived from the \\Ha/\\Hb\\ emission line ratio and various global parameters of the galaxies. Our main findings are that: 1) $C({\\rm H}\\beta)$  is linked with the galaxy spectral type and colour, decreasing from early- to late-type spirals. 2) $C({\\rm H}\\beta)$ increases with increasing metallicity. 3) $C({\\rm H}\\beta)$ is larger in galaxies with an older stellar population. 4) $C({\\rm H}\\beta)$ is larger for more luminous galaxies. 5) The extinction of the stellar light is correlated with both the extinction of the nebular light and the intrinsic galaxy colours. We propose phenomenological interpretations of our empirical results. We have also cross-correlated our sample of SDSS galaxies with the IRAS data base. Due to the lower redshift limit imposed to our sample and to the detection limit of IRAS, such a procedure selected only luminous infrared galaxies. We found that correlations that were shown by other authors to occur between optical and infrared properties of galaxies disappear when restricting the sample to luminous infrared galaxies. We also found that the optical properties of the luminous infrared galaxies in our SDSS sample are  very similar to those of our entire sample of SDSS galaxies. This may be explained by the IRAS luminosity of the galaxies originating in the regions that formed massive stars less than 1\\,Myr ago, while the opacity of galaxies as derived from the \\Ha/\\Hb\\ emission line ratio is due to diffuse dust. We show some implications of our empirical results on the determination of global star formation rates and total stellar masses in normal galaxies. ",
        "introduction": "Accounting for the presence of dust  is paramount for our understanding of the constitution and evolution of galaxies. Indeed, dust  modifies the light we receive from galaxies by both dimming it and modifying its colour. The determination of the bolometric luminosity of galaxies, the description of their stellar populations using galaxy colours, the estimates of the star formation rates using observed fluxes in either the emission lines or ultraviolet continua, all depend on a correction for the effects of dust. Dust in itself is also an important constituent of galaxies, not so much by its mass but essentially by the effects it has on the thermal balance of the interstellar medium  and on the formation of H$_2$ molecules, which has important consequences on the efficiency of star formation (e.g. Omukai \\cite{O00}, Hirashita \\& Ferrara \\cite{HF02}).  Yet, the question of dust opacity of galaxies is still a subject of strong debate. Earlier, the presence of dust in galaxies was essentially inferred from its impact on the observed distribution of stellar light. The spectacular extinction in the edge-on Sombrero galaxy was attributed to dust located in its plane.  The mottled aspect of the arms in face-on spiral galaxies  was attributed to higher extinction in zones of star formation. Such observations developed the view that dust is a common constituent of spiral galaxies.  In lenticular galaxies, dust lanes have been noticed  (Sandage \\cite{S61}).  Studies using various techniques have led to the conclusion that the average extinction decreases following the sequence late-type spirals -- early-type spirals -- lenticulars and ellipticals (see review by Calzetti \\cite{C01}). Inclination tests measure the dependence of galaxy disk surface brightness on inclination (e.g. Valentijn \\cite{V94} and references therein). This method is statistical by nature and strongly depends on sample selection, as discussed by Valentijn. Another method is to consider multiwavelength data of galaxies and solve simultaneously for the intrinsic colours of the stellar populations and for the reddening. The results of this method strongly rely on the adopted dust distribution models, as emphasized by Witt et al.\\,(\\cite{WTC92}), Bianchi et al.\\,(\\cite{BFG96}), and Witt \\& Gordon\\,(\\cite{WG00}). Infrared data from the IRAS survey provide detection of dust through its emission in the infrared. However, the interpretation of the infrared dust emission in terms of dust content is not straightforward (see e.g. Sauvage \\& Thuan \\cite{ST94}, Dale \\& Helou \\cite{DH02}). A different approach is to study the transparency of galaxies with respect to background light (see Keel \\& White \\cite{KW01} and references therein for methods using one nearby background galaxy, and Gonz{\\'a}lez et al. \\cite{GAD98} for methods using distant galaxies as background sources). Unfortunately, this method suffers from extremely low statistics.  Another extinction indicator that can be used in galaxies presenting emission lines is the observed Balmer decrement. As is known, the intrinsic intensity ratios of hydrogen recombination lines in ionized nebulae have a negligible dependence on the emission conditions (Osterbrock \\cite{O89}) so that the observed values are a consequence of the dust extinction being different at the relevant wavelengths. This has for example been used in the influential work by Calzetti et al.\\,(\\cite{CKS94}) to determine an empirical extinction law for starburst galaxies. However, few studies have used this method to investigate the relation between  extinction and overall galaxy type (e.g. Hubble type). In order to do this, one needs spectra of galaxies spanning the entire range of Hubble types. The atlas of Kennicutt (\\cite{K92}) provided such a data base for a limited sample of galaxies. Among these galaxies, 15 were considered by  Sodr\\'e \\& Stasi\\'nska (\\cite{SS99}, hereinafter SS99) to be normal emission line galaxies, and these authors showed that the extinction at \\Hb\\, as measured by the \\Ha/\\Hb\\ ratio (in the remaining of the paper we will call it the \\emph{Balmer extinction}) decreases steadily from early-type to late-type spirals.  Later, using the Nearby Field Galaxy Survey (NFGS) of Jansen et al. (\\cite{JFFC00a}, \\cite{JFFC00b}), which provided adequate data for about 100 galaxies, Stasi\\'nska \\& Sodr\\'e (\\cite{SS01}, hereinafter SS01) showed that redder galaxies have larger  Balmer extinction. The Sloan Digital Sky Survey (SDSS), which aims at obtaining spectra of $10^{6}$ galaxies in the nearby Universe, provides a wonderful opportunity to study this question in more detail.  The present paper makes use of the observations from the First Data Release (DR1; Abazajian et al. 2003, see also Stoughton et al. 2002) of the SDSS to  study the relation between the  Balmer extinction and other global properties of the galaxies. Section 2 explains the selection of the observational sample and quantities used in the analysis. Section 3 presents  the relation between  Balmer extinction and other galaxy parameters. Section 4 presents an interpretation of our results. In Section 5 we cross-correlate the galaxies of our sample with galaxies detected by IRAS, looking for additional clues on the origin of the observed extinction. Section 6 summarizes the main conclusions of this work and outlines some prospects.  \\begin{figure} \\centerline{\\includegraphics[width=8cm]{stasinska_f1.ps}} \\caption{Examples of two fits using our code to measure the equivalent widths and fluxes of H$\\beta$ (top panels) and H$\\alpha$ (bottom panels). The left and right figures are for spectra with signal-to-noise ratio in $g$-band of 5.9 and 12.4, respectively. We also show the equivalent widths (EW) in \\AA\\ and signal-to-noise ratio (S/N) of the measured emission lines.} \\label{ews1} \\end{figure} ",
        "conclusions": "We have thus found (confirming the results of SS99 and SS01) that the Balmer extinction of galaxies decreases steadily from early- to late-type spirals (whether classified according to their spectral type or according to their colour). We have also shown that there is a direct dependence of Balmer extinction on metallicity, and also on the age of the stellar populations and on the total galaxy luminosity. This last statement is consistent with the finding of Wang \\& Heckman (\\cite{WH96}), based on relations using far-ultraviolet and far-infrared fluxes, that ``the optical depth of normal galactic disks increases with galaxy luminosity''. On the other hand, the dependence we find on galaxy type seems to contradict the general opinion summarized by Calzetti (\\cite{C01}) and the recent result by  Kauffmann et al.  (2003) who conclude that ``galaxies with the youngest stellar populations are the most attenuated by dust''. While previous studies were sometimes based on small samples, the result of Kauffmann et al. (2003) comes from a model fitting of the continua of  $10^{5}$ galaxies from the SDSS. It should be noted that what they measure with this elaborate procedure is the attenuation of the \\emph{stellar} light.  Our work concerns the extinction of \\emph{nebular} light and the determination of the Balmer extinction is straightforward.  The interpretation in terms of global opacity of galaxies is however not simple: with  \\Ha/\\Hb\\  we measure  some sort of average opacity of the zones dimming the light from \\hii\\ regions (excluding the most opaque ones). Still, it is striking to see how well $C({\\rm H}\\beta)$ correlates with the galaxy spectral type. In this section, we present a possible phenomenological interpretation of our results, and show that our findings can actually be reconciled with the results of Kauffmann et al. (2003).  \\begin{figure} \\centerline{\\includegraphics[width=5.cm]{stasinska_f10.ps}} \\caption{Histogram of the values of the Balmer extinction at \\Hb, $C(\\Hb)$ in our sample of galaxies. } \\label{fig8} \\end{figure}   \\subsection{A phenomenological interpretation of the observed Balmer extinction trends in our sample} Let us write $C({\\rm H}\\beta)$ as the following product: \\begin{equation} C({\\rm H}\\beta) = n_{\\rm d} l \\sigma_{\\rm d} = \\frac{n_{\\rm d}}{n_{\\rm M}} \\frac{n_{\\rm M}}{n_{\\rm H}} \\frac{n_{\\rm H}}{n_{*}} n_{*} l \\sigma_{\\rm d} \\end{equation} where $n_{\\rm d}$ is average number of dust particles per unit volume in a galaxy,  $n_{\\rm M}$ is average  number of metallic atoms (in any form) per unit volume, $n_{\\rm H}$  is the average number of hydrogen particles per unit volume, $n_{*}$ is the average number of stars per unit volume, $l$ is the geometrical thickness of the region responsible for the optical extinction,  and $\\sigma_{\\rm d}$ is the typical extinction cross section of dust grains.  The common view that the opacity should increase from early- to late-types is comforted by such arguments that  ``galaxies with young stars contain more gas and hence more dust than galaxies with old stellar populations'' (Kauffmann et al. \\cite{K03}). On the other hand, early-type galaxies are more metal-rich and on average more massive than late-type galaxies, so it is not necessarily surprising that the extinction is actually larger for late-type galaxies. In Sect. 3.2 we have given arguments to say that the Balmer extinction is also determined by the mean age of the stellar population. This could indicate that in late-type galaxies $n_{\\rm d}/n_{\\rm M}$ is larger than in early-type ones. Such a view seems indeed to be supported by recent models for the evolution of dust in galaxies (Hirashita \\cite{H99}) which take into account the processes of formation and destruction of dust (note however that there are presently many assumptions and uncertainties in such models, see e.g. Dwek \\cite{D98} or Edmunds \\cite{E01}). In these models, condensation in cool stellar winds from low-mass stars is an important source of dust production.  We note that 75\\% of the galaxies of our sample have a measured $C({\\rm H}\\beta)$ between 0.3 and 0.9, as seen from the histogram shown in Fig. 10. On the other hand, judging from the metallicities derived for galaxies of luminosities similar to those of our sample (Zaritsky et al. \\cite{ZKR94}, Charlot et al. \\cite{CKL02}), the  metallicity range in our sample is likely higher than just a factor of three. This suggests that factors other than metallicity act to reduce the observed range in $C({\\rm H}\\beta)$. Obviously, $n_{\\rm H}/n_{*}$ is a good candidate, since it decreases from late- to early-type galaxies (Roberts \\& Haynes \\cite{RH94}).  On the other hand $ n_{*} l$, which can be assimilated to the stellar surface density, decreases from early- to late-types.  The last factor in Eq. (3) is  $\\sigma_{\\rm d}$, and one might expect some systematic effects if the grain size distribution depends on the processes for grain growth or destruction that could have different relative importances in galaxies of different types.  In conclusion, dust extinction in galaxies involves many factors, and our finding that $C({\\rm H}\\beta)$ increases from late- to early-types can easily be accounted for, at least qualitatively, within our present-day understanding of galaxies and their constituents.   \\begin{figure} \\centerline{\\includegraphics[width=5.13cm]{stasinska_f11.ps}} \\caption{The relation between \\EWHb and \\EWHa\\ in our sample of galaxies. The layout of the figure is the same as for Fig. 5. } \\label{fig9} \\end{figure}  \\subsection{The extinction of stellar light versus the extinction of nebular light}  As emphasized above, what we measure with $C({\\rm H}\\beta)$ is actually the reddening of the nebular light (called $C_{l}$ in SS01). On the other hand, what is determined by Kauffmann et al. (2003) is  the reddening of the stellar light (called $C_{c}$ in SS01). We now examine the relation between these two quantities in our sample of galaxies.  It had already been noted by SS99 and SS01, on much smaller samples of galaxies, that  \\EWHa\\  and \\EWHb\\ correlate extremely well (see also Kennicutt 1992). This implies that the difference between the extinction of the stellar continuum and that of the nebular emission is strongly linked to the colours of the galaxies, being larger for redder galaxies.  (see Eq. 5 in SS01). In Fig. 11 we show the values of  EW(H$\\beta$) as a function of EW(H$\\alpha$) in our sample of galaxies. The correlation is extremely strong, with $\\tau_{K} = 0.782$ and $r_{S} = 0.931$. Assuming that the relation between these emission line equivalent widths is EW(H$\\beta$) = $A \\times$ EW(H$\\alpha$) (the same model as adopted in SS01), an ordinary least square bisector linear fitting (Isobe et al. 1990) gives $A = 0.185 \\pm 0.001$. Considering only objects for which EW(H$\\alpha$) $>$ 20 \\AA\\ the result is the same, $A = 0.185 \\pm 0.001$. These results may be compared with $A = 0.194 \\pm 0.011$ and $A = 0.245 \\pm 0.007$, presented in SS99 and SS01, respectively. Our value   $A$ is consistent with that obtained by SS99 but not with that obtained by SS01. The cause of this discrepancy is not clear; it might be due to differences in the spectral resolution, because the spectral resolution of the SDSS spectra is almost half that of the NFGS spectra (Jansen et al. \\cite{JFFC00a}, \\cite{JFFC00b}) used by SS01.  From Eq. (5) in SS01 and the value of $A$, we derive that \\begin{equation} C_{c} = C_{l} - 0.81 - 2.99 ~ {\\rm log}\\, \\frac {F_{c}^{o} ({\\rm H}\\alpha)}{F_{c}^{o}({\\rm H}\\beta)}, \\end{equation} where $F_{c}^{o}$(H$\\alpha$) and $F_{c}^{o}$(H$\\beta$) are the intrinsic (i.e. not affected by extinction) stellar fluxes in the continuum adjacent to \\Ha\\ and \\Hb\\ respectively. We can use the dust-free spectrophotometric models of Barbaro \\& Poggianti (1997) to relate the value of $F_{c}^{o}$(H$\\alpha$)/$F_{c}^{o}$(H$\\beta$) to the values of $D(4000)$ and \\EWHa. Using panel $j$ of our Fig. 6 we can relate \\EWHa\\ to log (\\Ha/\\Hb) (by taking the median value of \\Ha/\\Hb\\ for a given \\EWHa) and therefore to  $C_{l}$. With the help of Eq. (4) we thus find that $C_{c}$ goes from  about  0.02 for $D(4000) = 1.7$ to about  0.2 for $D(4000) = 1.3$. This goes in the same direction as the results of  Kauffmann et al. (2003) who find that the median value of $A_{z}$, which is approximately equal to our $C_{c}$, roughly goes from 0.2 at $D(4000) = 1.7$ to 0.6 at $D(4000) = 1.3$. That the numbers are not exactly the same as the ones found by our analysis is not necessarily a worry given the dispersion in the observational points (both here and in the work of Kauffman et al. 2003) and given that the definition of the derived extinction is not exactly the same.  As for $C_{l}$, it can be evaluated using Fig. 6c and Eq. 1. We find $C_{l}$   $\\simeq$ 0.8 for $D(4000) = 1.7$ and $C_{l}$ $\\simeq$ 0.3 for $D(4000) = 1.3$ We note that $C_{c}$ is smaller than  $C_{l}$ at both extremes of the spectral type range, which can be interpreted as due to the fact that dust is more concentrated (and thus more opaque to radiation) in molecular clouds associated with \\hii\\ regions than in the diffuse interstellar medium. One may wonder why $C_{c}$ increases from early- to late-type spirals while $C_{l}$ decreases. The answer to this may be related to the fact that the stellar light from early-type galaxies is dominated by the bulge and to a specific distribution of dust resulting from dynamical effects. Advanced 3D-modelling of the star, dust and gas distribution in galaxies would be needed to test  any interpretation of our empirical result.  \\begin{figure*} \\centerline{\\includegraphics[width=10.9cm]{stasinska_f12.ps}} \\caption{Galaxies from our sample detected by IRAS (see Sect. 5). Plots of log (\\Ha/\\Hb) vs. log $L(\\rm FIR)$ (panel $a$) and log $L(\\Hb)$ (corrected for Balmer extinction)  vs log $L(\\rm FIR)$ (panel b). } \\label{fig10} \\end{figure*}   \\begin{figure*}[t] \\centerline{\\includegraphics[width=16cm]{stasinska_f13.ps}} \\caption{Same diagrams as in Fig. 6, with the galaxies detected by IRAS represented with large circles. } \\label{fig11} \\end{figure*}     We have used the observations from the First Data Release of the SDSS to examine a sample of normal galaxies (as opposed to galaxies with an active nucleus) and to investigate  how the Balmer extinction $C({\\rm H}\\beta)$ (i.e. the extinction at the wavelength of \\Hb\\ derived from the \\Ha/\\Hb\\ emission line ratios) relates with other global properties of the galaxies. Our selection criteria to build up the sample resulted in a data set of 9840 galaxies with adequate data. All these galaxies are at redhifts larger than 0.05 to avoid strong aperture effects.  Our main findings are the following: \\begin{enumerate} \\item $C({\\rm H}\\beta)$  is linked with the galaxy spectral type and colour, decreasing from early- to late-type spirals. \\item $C({\\rm H}\\beta)$ increases with increasing metallicity \\item $C({\\rm H}\\beta)$ is, probably, also affected by the age of the stellar population, being larger in the case of older stellar populations. \\item $C({\\rm H}\\beta)$ depends on galaxy masses. \\item  The extinction of the stellar light is correlated with both the extinction of the nebular light and the intrinsic galaxy colours, resulting in a trend with galaxy colour that may be opposite to the trend of $C({\\rm H}\\beta)$. \\end{enumerate}  The present work thus confirms the conclusions of our previous studies (Sodr\\'e \\& Stasi\\'{n}ska \\cite{SS99} and Stasi\\'{n}ska \\& Sodr\\'e \\cite{SS01}), which were based on much smaller samples and used data with lower spectral resolution (Kewley et al. \\cite{KGJD02}, using the same sample as SS01, also found that early-types in that sample are more heavily reddened than late-types).  Compared to our previous studies, the large number of galaxies in the SDSS sample allows us to investigate issues related to the inclination of galaxies.  The fact that  $C({\\rm H}\\beta)$ correlates so well with other properties of galaxies  is remarkable, given that the extinction, especially in late-types, is known to be not uniform across the face of galaxies (e.g. Beckman et al. \\cite{BPK96}).  We have cross-correlated our sample of SDSS galaxies with the IRAS data base in order to investigate any relationship between $C({\\rm H}\\beta)$ and total infrared luminosity of the galaxies. Due to the lower redshift limit imposed to our sample and to the detection limit of IRAS, such a procedure selected only luminous infrared galaxies. We found that correlations that were shown by other authors to exist between optical and infrared properties of galaxies disappear when restricting to luminous infrared galaxies. We also found that the optical properties of the luminous infrared galaxies in our SDSS sample are very similar to those of our entire sample of SDSS galaxies.  We have proposed a phenomenological interpretation of our findings.  We suggest that the main driver of the  Balmer extinction of galaxies is their mass, combined with their metallicity and presence of old stellar populations. The infrared luminosity of the galaxies as determined by IRAS, which is attributed to radiation from hot stars reprocessed by dust grains, samples the regions with {\\it the most recent episodes} of star formation, and is not connected with the Balmer extinction. Obviously, detailed modelling of the spectral light from galaxies taking into account the effects of dust and using a complete code such as GRASIL (Silva et al. 1998, see also http://web.pd.astro.it/granato/grasil/grasil.html) is needed for a deeper understanding of the empirical relations we have found. This is not an easy task, however, since as noted by Witt at al. (1992) and Witt \\& Gordon (2000), equal amounts of dust in different configurations may produce very different reddening and attenuation effects. In any case, future models of the integrated light from galaxies including the effects of dust should also aim at reproducing the correlations we have shown.  An important outcome of our study is to open the way for an improved correction for extinction in the determination of such parameters as the global star formation rate in galaxies or their total stellar masses. For normal galaxies, the global star formation rate can be obtained from the total \\Ha\\ luminosity corrected for extinction using the extinction derived from the \\Ha/\\Hb\\ emission line ratio (keeping in mind the reservations expressed e.g. by Hirashita et al. 2003). If observations do not allow one to determine the \\emph{Balmer extinction}, one can make use of e.g. the observed galaxy colour or the $D(4000)$ parameter to obtain an estimate of statistical value since we have shown that all these quantities are correlated. On the contrary, the total stellar mass can be estimated from the observed stellar fluxes of the galaxy after correcting for \\emph{stellar extinction}.  This should be done with a proper model fitting of the observed continuum as in Kauffmann et al. (2003). However, the strong correlation that we have found empirically between  stellar extinction, Balmer extinction and  galaxy colours can provide a basis for a statistical method to determine the total masses of galaxies. These aspects will be developed in future work and should be important especially for the study of galaxies at intermediate and high redshifts."
    },
    "0403/astro-ph0403641_arXiv.txt": {
        "abstract": "Evolution of a universe with homogeneous extra dimensions is studied with the benefit of a well-chosen parameter space that provides a systematic, useful, and convenient way for analysis. In this model we find a natural evolution pattern that entails not only stable extra dimensions in the radiation-dominated era, thereby preserving essential predictions in the standard cosmology, but also the present accelerating expansion while satisfying the limit on the variation of Newtonian gravitational constant. In this natural evolution pattern the extra dimensions tend to be stabilized automatically without resorting to artificial mechanisms in both the radiation-dominated and the matter-dominated era, as a wonderful feature for building models with extra dimensions. In addition, the naturalness of this evolution pattern that guarantees the late-time accelerating expansion of a matter-dominated universe presents a solution to the coincidence problem: why the accelerating phase starts at the present epoch. The feasibility of this evolution pattern for describing our universe is discussed. ",
        "introduction": "The observations by Supernova Cosmology Project and Supernova Search Team have suggested in 1998 that the expansion of the present universe is accelerating \\cite{Perlmutter:1999np,Riess:1998cb}. This conclusion is reinforced recently in 2003 by WMAP measurements \\cite{WMAP2003}. One general conclusion from these measurements and the CMB observations in recent years \\cite{Sievers:2002tq,Kuo:2002ua,WMAP2003} is that the universe has the critical density, consisting of $1/3$ of ordinary matter and $2/3$ of dark energy with a negative pressure \\cite{CMB&SN} (such that $p_{\\textsc{x}}/\\rho_{\\textsc{x}} < -0.78$ \\cite{WMAP2003}). Although this acceleration may be driven by the existence of a positive cosmological constant (vacuum energy) \\cite{Lambda models}, there remain other interpretations of the accelerating expansion, such as ``quintessence'' (a slowly evolving scalar field) \\cite{Caldwell:1998ii,ComplexQ} and the presence of extra dimensions \\cite{Gu:2001ni,Gu:2002mz}.   The presence of extra dimensions is required in various theories beyond the standard model of particle physics, especially in the theories for unifying gravity and other forces, such as superstring theory. Extra dimensions should be ``hidden'' (or ``dark'') for consistency with observations. Various scenarios for ``hidden'' extra dimensions have been proposed, for example, a brane world with large compact extra dimensions in factorizable geometry proposed by Arkani-Hamed \\emph{et al.} \\cite{Arkani-Hamed} (see also \\cite{Antoniadis:1990ew}), and a brane world with noncompact extra dimensions in nonfactorizable geometry proposed by Randall and Sundrum \\cite{Randall&Sundrum}. In this paper we will employ the simplest scenario: small compact extra dimensions in factorizable geometry, as introduced in the Kaluza-Klein (KK) theories \\cite{Kaluza&Klein}.   The possibility of generating the accelerating expansion of the present matter-dominated universe via the evolution of homogeneous and isotropic extra dimensions is first pointed out by Gu and Hwang \\cite{Gu:2001ni}. Furthermore, the general idea of unifying dark energy sources (i.e.\\ dark matter and dark energy) and dark geometry (e.g.\\ extra dimensions) into one has been sketched by Gu \\cite{Gu:2002mz}, who pointed out that dark geometry (instead of dark energy) can be an intriguing candidate for generating accelerating expansion. As indicated by Einstein's general relativity, mass (energy) and geometry are two faces of one single nature, therefore it is biased to consider only mass (energy) has dark part, while believing blindly that all geometry is totally ``visible'' to our poor eyes \\cite{Gu:2002mz}.  In this simple scenario making use of a highly symmetric extra space, there are much fewer free parameters (only two additional degrees of freedom, the expansion rate and the curvature of the extra space, in addition to those in the standard cosmological models without dark energy) so that it is much easier to be ruled out, compared with the quintessence models, % by constraints from observations. In particular, an essential difficulty of this model \\cite{Cline:2002mw} stems from the constraint on the variation of the Newtonian gravitational constant, which will be produced along with the evolution of extra dimensions that is the key element for generating accelerating expansion.   In this paper we will study a more general case in which the evolution of our universe in various eras, especially the present accelerating expansion era, is governed not only by the matter contents (excluding dark energy) therein, but also by the curvature of the ordinary 3-space and the evolution of extra dimensions. We will explore the feasibility of this model for generating accelerating expansion while satisfying observational constraints, especially the variation of the Newtonian gravitational constant. % Through these studies we will also present several nice features of this model, such as the automatic stabilization of extra dimensions and the solution to the cosmic coincidence problem --- why the energy densities of dark energy and dark matter are comparable now (i.e.\\ ``why now'' problem), or, more precisely (if dark energy is not a necessary ingredient for the accelerating expansion), why the accelerating phase of our universe starts at the present epoch. % ",
        "conclusions": "We have discussed the evolution of a universe under the influence of flat extra dimensions in various eras via a well-chosen parameter space that provides a systematic, useful, and convenient way for analysis. Since we assume the flatness (zero curvature) of the extra space, the only ingredient of extra dimensions that can affect the evolution of the universe is the evolution of their size. We have studied general features of the evolution, thereby exploring whether the picture is viable, subject to constraints from observations, for describing our world.  Through the above studies we have found a natural evolution pattern from the blazing era to the matter-dominated era in this model. In the blazing era, the ordinary space and the extra space tend to synchronize their expansion rates, i.e., $v/u \\rightarrow 1$, and meanwhile the universe is decelerating, presuming that an extremely tiny curvature contribution $k_a/a^2 u^2$ has been guaranteed initially by inflation.\\footnote{This tendency corresponds to the fixed point at $(k_a/a^2 u^2 \\, , \\, v/u)=(0,1)$ in the parameter space for the blazing era, which is stable in the Y direction but unstable in the X direction.} This tendency provides a natural initial condition, $(k_a/a^2 u^2 \\, , \\, v/u)=(\\pm \\epsilon ,1)$, for the succeeding radiation-dominated era (where `$\\pm \\epsilon$' indicates a tiny curvature contribution).  In the radiation-dominated era, the expansion rate of the extra space, $v$, tends to approach zero from $v/u \\simeq 1$ (as suggested above), meanwhile maintaining the decelerating phase.\\footnote{This tendency corresponds to the fixed point at $(k_a/a^2 u^2 \\, , \\, v/u)=(0,0)$ in the parameter space for the radiation-dominated era, which is stable in the Y direction but unstable in the X direction.} It is a good feature for building models with extra dimensions that the expansion rate of the extra space decreases to zero automatically, that is, extra dimensions are stabilized automatically in this era without resorting to any artificial mechanism. In particular, for small enough $v$ we can recover the standard cosmology (without extra dimensions) and the essential predictions therein (in particular, Big Bang Nucleosynthesis) in the radiation-dominated era. Therefore in our model with extra dimensions the predictions in the standard cosmology in the radiation-dominated era can be preserved in a natural way.  In the matter-dominated era, a natural initial condition, $(k_a/a^2 u^2 \\, , \\, v/u)=(\\pm \\epsilon , 0)$, is provided by the preceding, radiation-dominated era. For the case of a negative curvature, i.e.\\ with the initial condition $(-\\epsilon , 0)$, the universe will eventually change its initial decelerating phase to the accelerating one and approach the attractor at $(k_a/a^2 u^2 \\, , \\, v/u)=(-1,0)$ with stable extra dimension and significant negative curvature contribution. As we have shown, this evolution pattern is marginally consistent with the limit on the variation of the Newtonian gravitational constant. The maximal acceleration associated with this evolution pattern deviates from the one required by SN Ia data together with CMB and LLS observations (within the framework of the standard cosmology) \\cite{Perlmutter:1999jt} with $2.5 \\sigma$ deviation. Thus this evolution pattern is still not ruled out by the present observations.  We note that in this natural evolution pattern the expansion rate of extra dimensions will eventually approach zero, consequently extra dimensions being stabilized automatically without resorting to any artificial mechanism, in both the radiation-dominated and the matter-dominated era. This is the key ingredient for generating accelerating expansion of the present universe as well as satisfying the limit on the variation of the Newtonian gravitational constant (both caused by the evolution of extra dimensions), meanwhile preserving essential predictions in the standard cosmology for the radiation-dominated era. In particular, we emphasize that the naturalness of this evolution pattern, whose behavior is not sensitive to the initial conditions, indicates a solution to the cosmic coincidence problem (``why now'' problem) of dark energy % --- why dark energy starts dominating the universe now, or, more precisely, why the accelerating phase starts at the present epoch. % The existence of this natural evolution pattern implies that the late-time accelerating expansion of an open universe \\footnote{Actually this open universe is nearly flat in the earlier time.} in the matter-dominated era is guaranteed, accordingly the cosmic coincidence problem being solved, in our flat ED model with the help of inflation in the very early time.    Nevertheless, type-(d) trajectories in general entail a too short age of the universe % and negative curvature in the ordinary space rather than almost zero curvature suggested by CMB data within the framework of standard cosmology. These features might eventually rule out this evolution pattern for describing our universe. Further detailed studies about the consistency with observations are in progress. Moreover, in order to complete the picture of the scenario for a universe with extra dimensions and find the best-fit trajectory therein, a more general model with nonzero curvature in the extra space is also under investigation."
    },
    "0403/astro-ph0403194_arXiv.txt": {
        "abstract": "Searches for CO emission in high-redshift objects have traditionally suffered from the accuracy of optically-derived redshifts due to lack of bandwidth in correlators at radio observatories.  This problem has motivated the creation of the new COBRA continuum correlator, with 4~GHz available bandwidth, at the Owens Valley Radio Observatory Millimeter Array.  Presented here are the first scientific results from COBRA.  We report detections of redshifted CO($J=3\\rightarrow2$) emission in the QSOs \\smm  \\ and VCV J140955.5+562827, as well as a probable detection in \\rxj. At redshifts of $z=2.846$, $z=2.585$, and $z=2.796$, we find integrated CO flux densities of $5.4~\\jykms$, $2.4~\\jykms$, and $2.9~\\jykms$ for \\smm, VCV J140955.5+562827, and \\rxj, respectively, over linewidths  of $\\Delta V_{FWHM} \\sim 350~\\kms$. These measurements, when corrected for gravitational lensing, correspond to molecular gas masses of order $M(\\hh) \\sim 10^{9.6-11.1}\\,\\msun$, and are consistent with previous CO observations of high-redshift QSOs.  We also report $3\\sigma$ upper limits on CO(3$\\rightarrow$2) emission in the QSO \\lbqs \\ of $1.3~\\jykms$. We do not detect significant 3~mm continuum emission from any of the QSOs, with the exception of a tentative ($3\\sigma$) detection in \\rxj \\ of $S_{3mm}=0.92~\\mjybeam$. ",
        "introduction": "High-redshift QSOs have been shown to be often associated with dusty host galaxies (e.g.\\ Barvainis \\& Ivison 2002b), and as such offer much potential as powerful probes of the formation and evolution of massive galaxies in the universe.  The masses of high-redshift host galaxies, as inferred by extrapolation from the locally measured correlation between central, supermassive QSO black hole masses and the velocity dispersion of stars in the host galaxy spheroids \\citep{gms2000,fm2000} are large.  Thus, high-redshift QSOs are signposts to the position and redshift of massive galaxies in the early universe. Many high-redshift QSOs have been detected by recent millimeter and submillimeter surveys (e.g.\\ Carilli et al.\\ 2001; Isaak et al.\\ 2002; Barvainis \\& Ivison 2002b), implying that there is a link between QSOs and the massive, dusty, and gas-rich galaxy population discovered by deep submillimeter surveys (e.g.\\ Smail, Ivison, \\& Blain 1997; Cowie, Barger, \\& Kneib 2002). Most of these so-called submillimeter sources are too faint for optical redshifts to be determined, even with the largest telescopes, precluding deep searches for the large molecular gas reservoirs associated with young, massive galaxies. By studying large samples of dusty, high-$z$ QSOs and the gas and dust properties of their host systems, we can evaluate their overlap with the submillimeter galaxy population and investigate their properties and evolution.  Recently, studies of CO emission in high-redshift objects have seen much success, with over 20 separate galaxies detected (e.g.\\ Barvainis, Alloin, \\& Bremer 2002a; Guilloteau et al.\\ 1999), the majority from galaxies with $S_{850\\micron}>10$~mJy (e.g.\\ Figure~1 in Isaak et al.\\ 2002).  From these measurements it is possible to infer the total molecular gas mass in the high-$z$ host galaxies, and in turn establish the fraction of the total mass that is still in gaseous form -- a suggestive indicator of evolutionary state.  CO measurements also provide long-sought evidence of previous metal enrichment, both in these objects and in the early universe, since carbon and oxygen are mainly produced by fusion reactions in stellar cores. These stars must also have produced copious amounts of other heavy elements.  We therefore expect to see metal-rich gas ($Z \\sim Z_{\\sun}$) in massive starbursting systems very soon after the first starbursts (after $\\sim 200$~Myr). The spectacular detections of both CO \\citep{fabian_sdss,berCO}\\ and inferred presence of dust \\citep{berdust,sdssdust}\\ in SDSS~1148+5251 indicate that heavy-element enriched gas was already present in some objects at $z=6.4$, within 800~Myr of the Big Bang.  Star formation must have already been an important process by $z \\sim 6.4$. The growing number of CO detections at lower, yet still cosmologically significant, redshifts, are very important, indicating how widespread star forming systems were at early epochs, and supporting the link between submillimeter galaxies and QSO hosts.  Searches for high-redshift CO have been hampered by the limited spectrometer bandwidths (typically about 500-1500 $\\kms$) available to date on the current generation of millimeter telescopes and interferometers. The lack of bandwidth prevents easy search for CO emission because the optically-determined redshifts for QSOs, typically measured from broad emission lines, are known to be blueshifted from the host galaxy redshifts by as much as 1000-2000 $\\kms$ \\citep{tytlerfan}.  As an example, the QSO APM 08279+5255 has an optical redshift of $z=3.87$ from \\citet{apmoptical}, determined through identification of its rest-frame UV emission lines; by contrast, the redshift of the CO line of its host galaxy, from \\citet{downes99}, is $z_{CO}=3.911$, a discrepancy of approximately 2500 $\\kms$ (blueshifted). At the Owens Valley Radio Observatory (OVRO), a new very broadband cross-correlator system, named COBRA, has been recently implemented to alleviate this problem. With a spectral coverage of 4 GHz, COBRA represents an eight-fold increase over the previous bandwidth at OVRO, and currently provides the largest available bandwidth for CO searches at millimeter wavelengths. This system is particularly advantageous for high-redshift spectral line searches when the redshift is poorly known \\emph{a priori}.  Presented in Figure~1 is a sample COBRA spectrum, showing 3~GHz of the 4~GHz available bandwidth.  As part of the correlator commissioning, we started a program to detect CO emission from high-redshift QSOs. To this end, we compiled a database of submillimeter-bright, $z>2$ QSOs from the literature from which to choose a sample of targets.  To be included in the list, each object was required to meet the following criteria: (1) detection in 2 or more submillimeter/far-infrared (FIR) bands; (2) $T_{dust} \\leq 100$ K, as derived following \\citet{MinSED}, such that the rest-FIR spectral energy distribution includes a significant cold dust contribution; and (3) FIR luminosity greater than or equal to that of the prototypical ULIRG Arp 220 (i.e.\\ $L_{FIR}>10^{12.2}\\,\\lsun$).  We chose a total of 4 QSOs from the total of 23 sources in this database to observe, selecting those with the greatest FIR luminosities but without a published CO detection: \\lbqs, \\smm, \\rxj, and VCV J140955.5+562827.  In the sections that follow, we detail our observations and report the results of our efforts, which are part of an ongoing program of observations at OVRO.  Throughout this paper, we assume an $\\Omega_{\\textrm{M}}=0.3$, $\\Omega_{\\Lambda}=0.7$ cosmology with $H_{0}=70~\\kms~\\mpc^{-1}$.  We also provide a parallel analysis in parentheses in the Einstein-deSitter cosmology $\\Omega_{\\textrm{M}}=1$ and $\\Omega_{\\Lambda}=0$, with $H_{0}=75~\\kms~\\mpc^{-1}$. ",
        "conclusions": "Three more high-$z$ QSOs have been detected in CO emission, \\smm \\ at $z=2.846$, VCV J140955.5+562827 at $z=2.585$, and \\rxj \\ at $z=2.796$, using the new COBRA correlator at the Owens Valley Radio Observatory. All three submillimeter-bright QSOs possess large reservoirs of molecular gas, further linking submillimeter galaxies to their QSO contemporaries, and \\smm \\ may be one of the most massive CO systems known."
    },
    "0403/hep-ph0403248_arXiv.txt": {
        "abstract": "\\widetext We propose a solution to the cosmological monopole problem: Primordial black holes, produced in the early universe, can accrete magnetic monopoles before the relics dominate the energy density of the universe. These small black holes  quickly evaporate and thereby convert most of the monopole energy density into radiation. We estimate the range of parameters for which this solution is possible: under very conservative assumptions we find that the black hole mass must be less than $10^9$gm. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403311_arXiv.txt": {
        "abstract": "{ We investigated the main-line spectral evolution with shell thickness of oxygen rich AGB stars. The study is based on a sample of 30 sources distributed along the IRAS colour-colour diagram. The sources were chosen to trace the Miras with thick shells and the whole range of OH/IR stars. The Miras exhibit a 1665~MHz emission strength comparable to that at 1667~MHz. Even though the Miras of the study have quite thick shells, their spectral characteristics in both main lines attest to a strong heterogeneity in their OH shell with, in particular, the presence of significant turbulence and acceleration. The expansion velocity has been found to be about the same at 1665 and 1667~MHz, taking into account a possible velocity turbulence of 1-2~km~s$^{-1}$ at the location of the main-line maser emission. An increase in the intensity ratio 1667/1665 with shell thickness has been found. A plausible explanation for such a phenomenon is that competitive gain in favour of the 1667~MHz line increases when the shell is getting thicker. There is an evolution in the spectral profile shape with the appearance of a substantial inter-peak signal when the shell is getting thicker. Also, inter-peak components are found and can be as strong as the external standard peaks when the shell is very thick. This trend for an increase of the signal in between the two main peaks is thought to be the result of an increase of the saturation with shell thickness. All sources but two - a Mira and an OH/IR star from the lower part of the colour$-$colour diagram - are weakly polarized. The strong polarization observed for those two particular objects is thought to be the result of perturbations in their shells. ",
        "introduction": "\\label{section Intro.}  The standard model proposed by Reid et~al. (1977) explains convincingly the double-peaked profile usually observed for Miras and OH/IR objects, particularly at 1612~MHz. It also explains the general features of the circumstellar envelopes. Nevertheless, many stars do not have such a simple OH maser profile. In particular, in the main lines, the peaks are often broad and consist of many components. To explain these more complex profiles new models have been suggested.  Alcock \\& Ross (1986) have shown that it is possible to change the profile shape for a partially or fully saturated maser when the masering region becomes sufficiently thick. When that happens, the profiles consist of broader peaks and a signal is seen at the velocity of the star. This is in agreement with the spectral profiles observed for Miras and OH/IR objects, but the inferred internal and external radii disagree with the MERLIN and VLA interferometric observations. To overcome this discrepancy, Alcock \\& Ross suggest that the mass loss is not a continuous and uniform process but rather consists of ejection of blobs of material in random directions such that the symmetrical geometry is preserved. The model calculation results of Collison \\& Nedoluha (1995) agree with the conclusions by Alcock \\& Ross that the observed OH profile cannot be produced by the standard model description of a smooth, spherically symmetric, steady windlike mass loss. The assumption of a non-smooth shell is justified by the interferometric maps themself. Indeed, with the increase in interferometer resolution the more detailed maps revealed that the OH maser emission is located in clumps (Welty, Fix \\& Mutel, 1987; Chapman, Cohen \\& Saikia, 1991) and/or the shell shows evidence for deviation from spherical symmetry (Diamond et~al. 1985; Bowers, Johnston \\& de~Vegt 1989).  Evidence for acceleration in the OH shell has been found for many stars (Chapman et~al. 1994; Etoka \\& Le~Squeren 1996, 1997; Szymczak et~al. 1998; Richards et~al. 1999). The importance of acceleration in the shape of the spectral profile is underlined by the work of Chapman \\& Cohen (1985). According to the value of the logarithmic velocity gradient $\\epsilon$, the shape of the observed line profile largely varies as follows~: (1) for $\\epsilon \\rightarrow 0$ the classical double-peaked profile is observed; (2) for $0 < \\epsilon < 1$ an inter-peak signal can be seen; (3) for $\\epsilon \\simeq 1$ a plateau-shaped profile is obtained and finally (4) for $\\epsilon > 1$ the profile is reduced to a single peak centred on the stellar velocity. Thus, many spectral profiles resembling those observed can be modelled by introducing acceleration in the OH shell.  Based on the IRAS measurements at 12, 25 and 60~$\\mu$m, Olnon et~al. (1984) discovered the existence of a continuous sequence from Miras to OH/IR objects in the so-called colour-colour diagram which plots the colour indices $[60-25]$ vs. $[12-25]$. This sequence has been analysed in terms of shell thickness by van~der~Veen \\& Habing (1988) and van~der~Veen \\& Rugers (1989). These studies lead to a commonly accepted interpretation that it is an evolutionary track followed by intermediate-mass stars. In this scenario, Miras are the progenitors of OH/IR stars involving among other characteristics an increase of their mass-loss rate and expansion velocity with time. Nonetheless, the initial mass of the star could play a role in the setting up of that sequence. Based on an alternative model for the evolution of AGB stars proposed by Epchtein et~al. (1990), L\\'epine et~al. (1995) proposed that the IRAS colour-colour sequence could be a sequence of mass.  The study presented here deals with the spectral profiles in both main lines of Miras and OH/IR stars distributed along the colour-colour diagram. The main aim is to determine whether there is a relation between the spectral profile and the thickness of the circumstellar shell.  The next section presents the observations and selection criteria. Section~\\ref{section Results} and Appendices~\\ref{appendix A},~\\ref{appendix B} and \\ref{appendix C} present the results concerning the individual stars and a comparison with previous observations. The discussion is presented in Sect.~\\ref{section Discus.} and the final conclusions in Sect.~\\ref{section Conclu.}. ",
        "conclusions": ""
    },
    "0403/astro-ph0403127_arXiv.txt": {
        "abstract": "We present 1.4\\,GHz \\hi\\ absorption line observations towards the starburst in NGC\\,2146, made with the VLA and MERLIN. The \\hi\\ gas has a rotating disk/ring structure with column densities between 6 and 18 $\\rm \\times 10^{21}\\:atoms\\, cm^{-2}$. The \\hi\\ absorption has a uniform spatial and velocity distribution, and does not reveal any anomalous material concentration or velocity in the central region of the galaxy which might indicate an encounter with another galaxy or a far--evolved merger. We conclude that the signs of an encounter causing the starburst should be searched for in the outer regions of the galaxy. ",
        "introduction": "Atomic hydrogen absorption measurements on (sub)arc\\-second scales allow us to investigate the distribution and dynamics of neutral gas in the inner regions of starburst galaxies and active galactic nuclei (AGN). While \\hi\\ emission measurements with current instruments are sensitivity--limited to a maximum resolution of a few arcseconds, absorption studies are limited only by the angular resolution of the telescope and the brightness distribution of the background radio continuum. We present \\hi\\ absorption measurements of the peculiar spiral galaxy NGC\\,2146.  It is well established that NGC\\,2146 is undergoing a strong starburst, even stronger than that in M\\,82 (e.g.\\ {\\mbox{\\citealt{kronberg81}}}; \\citealt{tarchi00}; later TNG); however, the origin of this starburst is still unclear. A starburst is often triggered by an interaction with another galaxy, which perturbs the potential equilibrium of the gas, causing a flow of gas towards the center which results in an increase in density and thus fueling of the star--formation process.  It therefore seems plausible that a starburst in a galaxy is triggered by an encounter if the galaxy belongs to a group which contains and still is connected by large \\hi\\ tails, for instance like M\\,81\\,--\\,M\\,82\\,--\\,NGC\\,3077 (\\citealt{yun94}, \\citealt{walter02}) and NGC\\,3627\\,--\\,NGC3628\\,--\\,NGC\\,3623 (the Leo triplet; \\citealt{zhang93}). Such an interaction has been searched for in NGC\\,2146 since the presence of a large \\hi\\ cloud is known since 1976.  Using the NRAO 91--m telescope, \\citet{fisher76} detected a huge \\hi\\ cloud extending out to six Holmberg radii ($\\sim$\\,120\\,kpc) around NGC\\,2146. This cloud could be the consequence of a tidal interaction or of an explosion/ejection in the galaxy. However, there exists no kinematic evidence for the explosion hypothesis, and no companion has been found which may have interacted with NGC\\,2146. The higher resolution observations of the extended \\hi\\ around NGC\\,2146 with the WSRT \\citep{caspers86} and the VLA \\citep{tara96} resolved the cloud into a prominent tail, extending out to 90\\,kpc SE of the body of the galaxy, but no further evidence of an interaction was found at that time. In 1990, \\citet{hutch90} suggested instead that NGC\\,2146 appears to be in the final stage of a fairly gentle far--evolved merger, with the dominant galaxy (NGC\\,2146) now seen close to edge--on and the stripped companion being on a final plunge toward its center. The putative traces of this merger are, however, not particularly compelling. Evidence of a collision with another galaxy which did not remain embedded in NGC\\,2146 has been suggested as an alternative triggering mechanism by \\citet{young88}. They drew especially attention to the 10\\,kpc extended semi--arc of \\hii\\ regions (observed in H$\\alpha$ and [SII]), which is not coplanar with the rotating disk of the galaxy. A more recent analysis by \\citet{tara01} (later TPB) of the \\hi\\ stream suggests today a tidal interaction between NGC\\,2146 and a Low Surface Brightness (LSB) companion of which a remnant is apparently still seen as a 1.5$\\times$10$^{8}$\\,M$_{\\sun}$ \\hi\\ concentration in the southern tail.  In order to study the kinematics and density of gas in the central region of NGC\\,2146 in the light of the proposed merger/encounter hypothesis, we have mapped the \\hi\\ absorption towards the nuclear radio continuum emission using the VLA (1\\farcs8 resolution = 130\\,pc) and MERLIN (0\\farcs2 resolution = 15\\,pc). At the distance of NGC\\,2146 (14.5\\,Mpc, \\citealt{benvenuti75}), 1$''$ is equivalent to 70\\,pc. ",
        "conclusions": "  \\noindent{-- the \\hi\\ absorption is observed over the entire continuum radiation background which coincides with the starburst region;}  \\noindent{-- the \\hi\\ gas towards the starburst region has a column density between 6 and 18 $\\rm \\times 10^{21}\\:atoms\\:cm^{-2}$;}  \\noindent{-- the \\hi\\ velocity field has the structure of a rotating disk and ring, as often found in the inner region of similar starburst systems;}  \\noindent{-- the \\hi\\ absorption has a uniform spatial and velocity distribution without any indication of an encounter with another galaxy or a far-evolved merger; there is no indication of a bar;}  \\noindent{-- the absence of anomalous gas concentration or velocities and the agreement of the \\hi\\ absorption rotation curve with those of different constituents (CO and H$\\alpha$), suggest that the inner gas has had time to relax after the encounter; the signs of an encounter causing the starburst must be searched for in the outer regions of the galaxy.}"
    },
    "0403/astro-ph0403257_arXiv.txt": {
        "abstract": "We show that orbit-superposition dynamical models (Schwarzschild's method) provide reliable estimates of nuclear black hole masses and errors when constructed from adequate orbit libraries and kinematic data.  We thus rebut two recent papers that argue that BH masses obtained from this method are unreliable.  These papers claim to demonstrate that the range of allowable BH masses derived from a given dataset is artificially too narrow as a result of an inadequate number of orbits in the library used to construct dynamical models.  This is an elementary error that is easily avoided.  We describe a method to estimate the number and nature of orbits needed for the library.  We provide an example that shows that this prescription is adequate, in the sense that the range of allowable BH masses is not artificially narrowed by use of too few orbits.  This is illustrated by showing that the $\\chi^2$ versus BH-mass curve does not change beyond a certain point as more orbits are added to the library.  At that point, the phase-space coverage of the orbit library is good enough to estimate the BH mass, and the $\\chi^2$ profile provides a reliable estimate of its errors.  A second point raised by critics is that kinematic data are generally obtained with insufficient spatial resolution (compared to the BH radius of influence) to obtain a reliable mass.  We make the distinction between {\\em unreliable} determinations and {\\em imprecise} ones.  We show that there are several different properties of a kinematic dataset that can lead to {\\em imprecise} BH determinations (insufficient resolution among them), but none of the attributes we have investigated leads to an unreliable determination. In short, the degree to which the BH radius of influence is resolved by spectroscopic observations is already reflected in the BH-mass error envelope, and is not a hidden source of error.  The BH masses published by our group and the Leiden group are reliable. ",
        "introduction": "There are almost twenty detections of massive black holes (BHs) in galaxy centers that employ the technique of orbit superposition modeling \\citep{vdm98,crvdb99, cdzdm99, geb00a, geb03, capp1459, verolme02b}.  The orbit superposition technique is based on a method originally invented by \\citet{sch79}, who noted that the time-averaged orbits in a stationary potential can be summed (at finite spatial resolution) to match the mass distribution that gives rise to the potential, thereby producing an equilibrium dynamical model. The Schwarzschild method was first used to analyze kinematic data for evidence of BHs in galaxy centers by Richstone \\& Tremaine (1985) for M87 and Dressler \\& Richstone (1988) for M31 and M32.  Those early models were spherically symmetric. In modern implementations of this method, large sets of orbits are run in a specified axisymmetric potential (based on the starlight distribution and a central point mass) and then a non-negative linear superposition of orbits is found that best matches the kinematic and photometric observations.  The usual goal of this process is the determination of only two quantities, the stellar mass-to-light ratio $\\Upsilon$ (assumed to be independent of position) and the BH mass $\\mbh$.  \\citet{valluri04} (hereafter VME) and \\cite{cretton04} have challenged the reliability of BH mass determinations obtained via orbit superposition modeling.  Cretton \\& Emsellem also argue that the uncertainties noted by VME can be dealt with via regularization (smoothing the distribution of orbit weights).  VME raise a long list of problems and features of the orbit-superposition analysis technique.  They focus on two main issues: orbit insufficiency and data insufficiency.  First, they contend that the BH mass determination is artificially narrowed through the use of too small an orbit library and imply, but do not show, that we have made this mistake. In addition they suggest that the absence of a flat-bottomed $\\chi^2$ profile as a function of BH mass indicates that this problem has occurred.  Second, they suggest that the BH mass $\\mbh$ can only be determined if the radius of influence $r_i = G\\mbh/\\sigma^2$ where the BH dominates the stellar dynamics is well-resolved spectroscopically (here $\\sigma$ is the line-of-sight velocity dispersion).  In each of their main points they have fastened on one facet of rather complex problems.  We show in \\S 2 that the use of an orbit library which is too small is an easily avoided elementary trap.  In \\S 3 we refer the reader to other work in which we show that we have adequate resolution to determine BH masses reliably, and argue that the presence of extensive radial kinematic coverage mitigates the uncertainties associated with a ``not well-resolved'' radius of influence.  We also show that the ``flat-bottomed'' behavior of $\\chi^2$ is a feature whose presence depends on several features of the quality and extent of the kinematic data, not just the size of the orbit library.  The principal result is that a $\\chi^2$ profile constructed using our method does indicate the range of acceptable BH masses in a manner easily judged by the readers of our papers. ",
        "conclusions": "We have made a point of distinguishing between {\\em imprecise} (the error bars are large) measurements of the BH mass $\\mbh$ and {\\em unreliable} ones (the error bars are wrong).  Figure 1 shows that increasing the number of orbits beyond our standard prescription does not broaden the range of BH masses allowed by our method.  Contrary to the assertion of VME, we must be using enough orbits.  The issue of data sufficiency is far more complex.  In \\S 3 we illustrate (or cite) several ways that data insufficiency can lead to imprecise $\\mbh$ estimates.  We also argue that these estimates are {\\em imprecise}, not unreliable.  Probably the best evidence for this is the fact that our $\\mbh$ estimates obtained with genuine orbit-superposition models lie within their own errors of repeat studies with greatly improved resolution \\citep{geb03}\\footnote {We have not investigated filled kinematic maps in radius and angle, as might be obtained with integral-field devices like SAURON. Especially where these data imply triaxiality (thereby violating the basic assumption of our axisymmetric models), there may be unexpected changes in $\\mbh$. }.  Flat-bottomed profiles---a special case of imprecision---can be produced by kinematic data with limited radial or angular coverage, by sparsely sampled data, or by insufficient resolution. A fourth way to produce a flat bottomed $\\chi^2$ profile is the use of {\\em noiseless} datasets constructed from two-integral models \\citep{cretton04}.  This procedure creates models that match the data perfectly.  Since axisymmetric models have more freedom than two-integral models, a range of three integral distribution functions will---by construction---provide an {\\em exact} match to the data. Since $\\chi^2$ is bounded by zero, a flat bottom in the $\\chi^2$ plot is unavoidable.  \\citet{cretton04} argue that regularization (smoothing the distribution of orbit weights) reduces the uncertainty in $\\mbh$.  Our models use a form of regularization (maximum entropy) but only as a numerical technique to accelerate convergence: we iterate our models while steadily reducing the weight of the entropy to the best-fit criterion until there is no further change in the best-fit model. \\citet{verolme02b} study models with and without regularization for M32 and find no difference in the best fit BH mass.   This paper has examined the two most serious recent criticisms of BH mass estimates from orbit superposition. We conclude that the issues raised by these criticisms have not led us to report erroneous BH masses.  There are a number of other tests of our procedure that also indicate that our method is reliable: \\begin{enumerate} \\item our method has been successfully blind-tested against (spherical) Jaffe models containing BHs; \\item our method yields the same masses as the Leiden code when applied to the same data (for NGC 821 and M32); \\item BHs weighed by this method display the same distribution of residuals from the mass-velocity dispersion relation as BHs weighed via gas dynamics \\citep{dorcarn}. \\end{enumerate} The tests described above are posted on the Nuker Team Website \\newline (http://www.noao.edu/noao/staff/lauer/nuker.html).   The BH masses published by our group (\\citealt{geb00a,bow01,geb03,sio04}) are obtained using an adequate number of orbits.  The plots of $\\chi^2(\\mbh)$ we publish provide the reader with an accurate assessment of the precision of the results. So far as we can tell, BH masses published from the Leiden group (\\citealt{vdm98,crvdb99,verolme02b,capp1459}) also should be reliable."
    },
    "0403/astro-ph0403582_arXiv.txt": {
        "abstract": "We present three-dimensional Monte Carlo radiative transfer models of a very young ($<10^{5}$ years old) low mass ($50\\,\\mathrm{M_{\\sun}}$) stellar cluster containing 23 stars and 27 brown dwarfs. The models use the density and the stellar mass distributions from the large-scale smoothed particle hydrodynamics (SPH) simulation of the formation of a low-mass stellar cluster by Bate, Bonnell and Bromm. Using adaptive mesh refinement, the SPH density is mapped to the radiative transfer grid without loss of resolution. The temperature of the ISM and the circumstellar dust is computed using Lucy's Monte Carlo radiative equilibrium algorithm. Based on this temperature, we compute the spectral energy distributions of the whole cluster and the individual objects.  We also compute simulated far-infrared Spitzer Space Telescope (\\emph{SST}) images (24, 70, and 160~$\\mathrm{\\mathrm{\\mu m}}$ bands) and construct colour-colour diagrams (near-infrared \\emph{HKL} and \\emph{SST} mid-infrared bands). The presence of accretion discs around the light sources influences the morphology of the dust temperature structure on a large scale (up to a several $10^{4}$~au). A considerable fraction of the interstellar dust is underheated compared to a model without the accretion discs because the radiation from the light sources is blocked/shadowed by the discs. The spectral energy distribution (SED) of the model cluster with accretion discs shows excess emission at $\\lambda=$~3--30~$\\mathrm{\\mu m}$ and $\\lambda>500\\,\\mathrm{\\mu m}$, compared to that without accretion discs. While the former is caused by the warm dust present in the discs, the latter is cause by the presence of the underheated (shadowed) dust. Our model with accretion discs around each object shows a similar distribution of spectral index (2.2--20~$\\mathrm{\\mathrm{\\mu m}}$) values (i.e. Class 0--III sources) as seen in the $\\rho$~Ophiuchus cloud. We confirm that the best diagnostics for identifying objects with accretion discs are mid-infrared ($\\lambda=$ 3--10~$\\mathrm{\\mu m}$) colours (e.g. \\emph{SST}~IRAC bands) rather than \\emph{HKL} colours. ",
        "introduction": "Systematic investigations of young stellar objects (YSOs) in a star-forming cloud including comparative studies of theoretical predictions and observations are important for understanding the stellar/substellar formation processes. Examples of well-studied star-forming clouds are the Orion Trapezium Cluster, NGC~2024, and the $\\rho$~Ophiuchus and Taurus-Auriga clouds. Recent observations of young clusters have been used to determine circumstellar disc frequencies, disc mass distributions, initial mass functions (IMFs) and evolutional stages of the objects in the clusters. It has been shown that \\emph{JHKL} and \\emph{HKL} colour-colour diagrams are particularly effective in identifying the presence of circumstellar discs (e.g. \\citealt{kenyon:1995}; \\citealt{McCaughrean:1995}; \\citealt{lada:2000}; \\citealt{haisch:2001}).  The spectral indices \\citep{lada:1987} or the slopes of observed spectral energy distributions (SEDs) of YSOs from near- to mid-infrared wavelengths are used to classify SEDs, and the classification scheme is often related to the evolutionary stages of YSOs (e.g. \\citealt*{adam:1987}; \\citealt{myers:1987}).  The distribution of circumstellar disk masses in young clusters can be measured from millimetre continuum emission (e.g.  \\citealt{eisner:2003} for NGC~2024). Determining IMFs of young clusters requires evolutionary models (e.g. \\citealt{d-antona:1997}; \\citealt{baraffe:1998}) and infrared spectroscopy for spectral type identifications (e.g. \\citealt{luhman:1999} for the $\\rho$~Ophiuchus cloud).   \\citet*{bate:2003a} presented results from a very large three-dimensional (3-D) smoothed particle hydrodynamics (SPH) simulation of the collapse and fragmentation of a 50\\,$\\mathrm{M_{\\sun}}$ turbulent molecular cloud to form a stellar cluster. The calculation resolved circumstellar discs down to $\\sim10$~au in radius and binary stars as close as 1~au.  Although some observational predictions, such as the IMF and binary fraction, may be gleaned directly from a hydrodynamical simulation of stellar cluster formation, the principal observable characteristics (optical, near-IR, IR, and sub-millimetre images and spectra) require further detailed radiative transfer modelling. The density distribution of such hydrodynamical calculations is very complicated, and the corresponding radiative transfer must be also performed in full 3-D.  There are two basic approaches to 3-D radiative transfer problems: \\,grid based methods (e.g. finite differencing, short- and long- characteristic methods), and particle (photon) based methods, i.e. Monte Carlo. Examples of the first kind are \\citet*{stenholm:1991}, \\citet{folini:2003} and \\citet*{steinacker:2002}. Those of the second kind include \\citet{witt:1996}, \\citet{pagani:1998}, \\citet*{wolf:1999}, \\citet{harries:2000} and \\citet{kurosawa:2001a}. The advantages of the second approach are, for example, the flexibility to treat a complex density distribution and a complex scattering function. Readers are referred to \\citet{steinacker:2003} and \\citet{pascucci:2003} for more extensive discussion on the advantages and disadvantages of these two different methods.  The temperature of the interstellar and circumstellar dust in the cluster must be calculated in order to determine the source function of the dust emission. The radiative equilibrium temperature of the dust particles can be found using the Monte Carlo method (e.g. \\citealt*{lefevre:1982}; \\citealt{wolf:1999}; \\citealt{lucy:1999b}; and \\citealt{bjorkman:2001}).  The technique used by \\citet{lucy:1999b} takes into account the fractional photon absorptions between two events of a `photon packet'; hence, it works well even in the limit of low opacity. \\citet{bjorkman:2001} used the immediate re-emission technique, in which radiative equilibrium is forced at each interaction with the dust. A photon is re-emitted immediately after an absorption event using a product of the dust opacity and the difference between the Planck function with a current temperature and that with a new temperature corresponding to the radiative equilibrium of the dust that absorbed the photon. Unlike the method of \\citet{lucy:1999b}, this does not require a temperature iteration if the opacity is independent of temperature. Alternatively, \\citet*{niccolini:2003} used the ray-splitting method showing its effectiveness for both low and high optical depth media.  Here we aim to simultaneously resolve the dust on parsec and sub-stellar-radius spatial scales, whilst including multiple radiation sources. To overcome the resolution problem, we have implemented an adaptive mesh refinement (AMR) scheme in the grid production process of the TORUS radiative transfer \\citep{harries:2000}. See also \\citet{wolf:1999}, \\citet{kurosawa:2001a} and \\citet{steinacker:2002} for a similar gridding scheme in a radiative transfer problem. The method described by \\citet{lucy:1999b} is used to compute dust temperatures in our models.  The objectives of this paper are: 1.\\,to compute observable quantities by solving the radiative transfer problem using a complex density distribution from the SPH calculation by \\citet{bate:2003a}; 2.\\,to analyse the predicted observational properties of the cluster generated in the simulation of \\citet{bate:2003a} at a distance of 140~pc which corresponds to the distance of nearby star-forming regions such as Taurus-Auriga and the $\\rho\\,$Ophiuchus cloud (e.g. \\citealt*{bertout:1999}).  In Section~2, we describe the details of our models. The results of the model calculations are given in Section~3. The conclusions are summarised in Section~4.   \\begin{figure}  \\begin{center}  \\includegraphics[% clip, scale=0.45]{fig01.eps}  \\end{center}  \\caption{Top: The scattering (dotted line), the absorption (dashed line), and the total (solid line) opacities of the grain model, described in Section~\\ref{sub:Dust-Model}, are shown as functions of wavelength. Bottom: The corresponding albedo of the grains is shown as a function of wavelength.}  \\label{fig:albedo}  \\end{figure} ",
        "conclusions": "}  We have presented three-dimensional Monte Carlo radiative transfer models of a very young low-mass stellar cluster with multiple light sources (23 stars and 27 brown dwarfs). The density structure and the stellar distributions from the large-scale SPH simulation of \\citet{bate:2003a} were mapped onto our radiative transfer grid without loss of resolution using an AMR grid. The temperature of the ISM and the circumstellar dust was computed using the Monte Carlo radiative equilibrium method of \\citet{lucy:1999b}. The results have been used to compute the SEDs, the far-infrared \\emph{SST} (MIPS bands) images, and the colour-colour diagrams of this cluster.  We find that the presence of circumstellar discs on scales less than 10~au (Model II) influences the morphology of the temperature structure of the cluster (Figure~\\ref{fig:temperature-map}), and can affect the temperature of the dust on large scales (up to a several $10^{4}$~au).  The dust shadowed by accretion discs has a lower temperature than the model without the discs (Model I). The radiation coming from within a few au of the light sources is anisotropic/bipolar because of the circumstellar discs.  The cluster SEDs (Figure~\\ref{fig:sed_compare}) from Models I and II both show peaks around $2\\,\\mathrm{\\mu m}$ and $200\\,\\mathrm{\\mu m}$. The first peak corresponds to that of stellar emission, and the latter corresponds to emission from ISM dust with $T\\approx20\\,\\mathrm{K}$. The excess emission of Model II between $\\lambda=$~3--30~$\\mathrm{\\mu m}$ is due to warm dust in the accretion discs. The excess emission in the far infrared ($\\lambda>500\\,\\mathrm{\\mu m}$) seen in Model II is caused by the colder ($T<20\\,\\mathrm{K}$) dust.  Assuming the distance to the cluster to be 140~pc, we have constructed simulated images (Figure~\\ref{fig:mips-images-limit}) for a \\emph{SST} MIPS (24, 70, and 160~$\\mathrm{\\mu m}$) observation using the appropriate diffraction limits, the sensitivity limits and the angular sizes of pixels. The emission at 24~$\\mathrm{\\mu m}$ traces the locations of the stars and brown dwarfs very well. No clear distinction can be made between the 70~$\\mathrm{\\mu m}$ image from Model I and Model II. The major structures in the 160~$\\mathrm{\\mu m}$ images appear to be same in both Models, but the emission from the isolated small structure in the lower half of the images is much weaker in Model II. In the simulated 24~$\\mathrm{\\mu m}$ images, 8 out of 50 objects are detected above the flux limit for Model I, and 18 out of 50 objects are detected for Model II. We note that 15 of the 27 brown dwarfs in the cluster (Model II) have $K<20$ and would therefore detectable via deep imaging, while those that are too faint tend to be the youngest, most deeply embedded sources.  Using the flux levels between 2.2~$\\mathrm{\\mu m}$ and 10~$\\mathrm{\\mu m}$, the spectral indices of each SEDs were computed. The objects were then classified according to the spectral index values and the ratio of $\\log\\lambda F_{\\lambda}$ values at $\\lambda=160\\,\\mathrm{\\mu m}$ (\\emph{SST} MIPS) and at $\\lambda=2.2\\,\\mathrm{\\mu m}$ ($K$ band). We have found that 54 per cent of the objects are classified as Class~0, none as Class~I, none as flat spectrum, 22 per cent as Class~II and 24 per cent as Class~III for Model I. For Model II, 56 per cent of objects are classified as Class~0, 18 per cent as Class~I, 8 per cent as flat spectrum, 12 per cent as Class~II, and 6 per cent as Class~III. We also found that the spectral index distribution of Model~II ( with $K<20$ objects) is very similar to that of the $\\rho$~Ophiuchus cloud observation by \\citet{greene:1994}.  Even though our objects are $0.07\\,\\mathrm{Myr}$ old, there are a mixture of Class 0--III objects; hence, the class does not necessary relate to the ages of YSOs.  According to the simulated \\emph{HKL} and mid-infrared (\\emph{SST} IRAC) colour-colour diagrams (Figures~\\ref{fig:color-color-ukirt} and \\ref{fig:color-color-irac}), the disc objects (Model II) tend to be slightly redder than discless objects (Model I) in the \\emph{H-K} vs \\emph{K-L} diagram, but the two populations are not well separated. This can be understood from the model cluster SEDs (Figure~\\ref{fig:sed_compare}) that show no significant difference in the flux levels between $\\lambda=$~1--3~$\\mathrm{\\mu m}$.  On the other hand, the disc objects and discless objects are clearly separated in the colour-colour diagram of the \\emph{SST} IRAC bands (Figure~\\ref{fig:color-color-irac}). As expected, we find that longer wavelengths are more efficient in detecting circumstellar discs. For example mid-infrared ($\\lambda=$~3--10~$\\mathrm{\\mu m}$) colours (e.g. \\emph{SST}~IRAC bands) are superior to \\emph{HKL} colours. We find that the intrinsic colours of the disc objects (Figure~\\ref{fig:color-color-irac}) show a distribution very similar to that found by \\citet{meyer:1997}.  The work presented in this paper can be extended to study how the observable quantities (e.g. colours of stars in a young cluster) evolves with time, by performing radiative transfer model calculations with density structures from a hydrodynamics calculation at different times (ages). Growth of dust grain sizes in accretion discs may become important in predicting how the colours of objects evolve (e.g. see \\citealt{hansen:1974}; \\citealt{wood:2002b}).  One of the obvious next steps in modelling star formation is to combine SPH and radiative-transfer in order to more accurately predict the temperature of the dust. It is interesting to note that the computational expense of the radiative equilibrium calculation performed here on one SPH time-slice (2400 CPU hours on UKAFF) is a significant fraction of that required to perform the complete SPH simulation (95\\,000 CPU hours on the same machine) which involves may thousands of time steps. Although it is not immediately clear how often the temperature would have to be computed during a hydrodynamic simulation (the radiative-transfer calculation would have its own Courant condition), it is apparent that a much more efficient method for calculating the radiation transfer is required before a combined SPH/radiative-transfer simulation becomes tractable."
    },
    "0403/astro-ph0403541_arXiv.txt": {
        "abstract": "Accreting black holes are believed to emit X-rays which then mediate information about strong gravity in the vicinity of the emission region. We report on a set of new routines for the {\\sc{}xspec} package for analysing X-ray spectra of black-hole accretion disks. The new computational tool significantly extends the capabilities of  the currently available fitting procedures that include the effects of strong gravity, and allows one to systematically explore the constraints on more model parameters than previously possible (for example black-hole angular momentum).  Moreover, axial symmetry of the disk intrinsic emissivity is not assumed, although it can be imposed to speed up the computations. The new routines can be used also as a stand-alone and flexible code with the capability of handling time-resolved spectra in the regime of strong gravity. We have used the new code to analyse the mean X-ray  spectrum from the long {\\it{}XMM--Newton} 2001 campaign of the Seyfert~1 galaxy MCG--6-30-15. Consistent with previous findings, we obtained a good fit to the broad Fe~K line profile for a radial line intrinsic emissivity law in the disk which is not a simple power law, and for near maximal value of black hole angular momentum. However, equally good fits can be obtained also for small values of the black hole angular momentum. The code has been developed with the aim of allowing precise modelling of relativistic effects. Although we find that current data cannot constrain the parameters of black-hole/accretion disk system well, the code allows, for a given source or situation, detailed investigations of what features  of the data future studies should be focused on in order to achieve the goal of uniquely isolating the parameters of such systems. ",
        "introduction": "There is now strong evidence that the \\fekalfa line emission in some active galactic nuclei (AGN) and some Galactic X-ray binary black-hole candidates (BHC) originates, at least in part, from an accretion disk in a strong gravitational field. A lively debate is aimed at addressing the question of what the spectral line profiles and the associated continuum can tell us about the  central black-hole, and whether they can be used to constrain parameters of the accretion disk in a nearby zone, about ten gravitational radii or less from the center. For a recent review for AGNs, see Fabian et al.\\ (2000), Reynolds \\& Nowak (2003), and references cited therein. For BHCs, see Miller et al.\\ (2002a), McClintock \\&  Remillard (2003) and references therein. In several sources there is indication of \\fekalfa line emission from within the last stable orbit of a Schwarzschild black hole (e.g.\\ in the Seyfert galaxy MCG--6-30-15; see Iwasawa et al.\\ 1996; Fabian et al.\\ 2000; Wilms et al.\\ 2001; Martocchia, Matt \\& Karas 2002a) while in other cases the emission appears to arise farther from the black hole (e.g.\\ in the microquasar GRS 1915+105, see Martocchia et al.\\ 2002b; for AGNs, see Yaqoob \\& Padmanabhan 2004 and references therein). Often, the results from X-ray line spectroscopy are inconclusive, especially in the case of low spectral resolution data. For example, a spinning  black hole is allowed but not required by the line model of the microquasar V4641 Sgr (Miller et al.\\ 2002b).  The debate still remains open, but  there are good prospects for future X-ray astronomy missions to be able to use the \\fekalfa line to probe the space-time in the vicinity of a black hole, and in particular to measure the angular momentum, or the spin, associated with the metric.  One may also be able to study the `plunge region' (about which very little is known), between the event horizon and the last stable orbit, and to determine if any appreciable contribution to the \\fekalfa line emission originates from there (Reynolds \\& Begelman 1997; Krolik \\& Hawley 2002). In addition to the Fe~K lines, there is some evidence for relativistic {\\it soft X-ray} emission lines due to the Ly$\\alpha$ transitions of oxygen, nitrogen, and carbon (e.g.\\ Mason et al.\\ 2003), although the observational support for this interpretation is still controversial (e.g.\\ Lee et al.\\ 2001).  With the greatly enhanced spectral resolution and throughput of future X-ray astronomy missions, the need arises for realistic theoretical models of the disk emission and computational tools that are powerful enough to deal with complex models and to allow actual fitting of theoretical models to observational data. It is worth noting that some of the current data have been used to address the issue of distinguishing between different space-time metrics around a black hole, however, the current models available for fitting X-ray data are subject to various restrictions which we shall elaborate on in the present paper.  In this paper we describe a generalised scheme and a code which can be used with the standard X-ray spectral fitting package, {\\sc{}xspec} (Arnaud 1996). We have in mind general relativity models for black-hole accretion disks. Apart from a better numerical resolution, the principal innovations compared to currently available schemes are that the new model allows one to  (i)~fit for the black-hole spin, (ii)~study the emission from  the plunge region, and (iii)~specify a more general form of emissivity  as a function of polar coordinates in the disk plane (both for the line and for the continuum). Furthermore, it is also possible to (iv)~study time variability of the observed signal and (v)~compute Stokes parameters of a polarized signal. Items (i)--(iii) are immediately applicable to current data and modelling, while the last two mentioned features are still mainly of theoretical interest at present. Time-resolved analysis and polarimetry of accretion disks are directed towards future applications when the necessary resolution and the ability to do polarimetry are available in X-rays. Thus our code has the advantage that it can be used with time-resolved data for reverberation studies of relativistic accretion disks  (Stella 1990; Reynolds et al.\\ 1999; Ruszkowski 2000; Goyder \\& Lasenby 2004). Also polarimetric analysis can be performed, and this will be extremely useful because it can add very specific information on strong-gravitational field effects (Connors, Stark \\& Piran 1980; Matt, Fabian \\& Ross 1993; Bao et al.\\ 1997). Theoretical spectra with temporal and  polarimetric information can be analysed with  the current version of our code and such analysis should provide tighter constrains on future models than is currently possible.  As mentioned above, Fe~K lines have been reported in the X-ray spectra of numerous AGNs and BHCs.  These sources frequently exhibit remarkable variability patterns which are still difficult to understand. Among these puzzling objects, the Seyfert~1 galaxy MCG--6-30-15 is radically distinctive with its broad and skewed \\fekalfa feature that persists in observations in spite of substantial variablity of the continuum. Interpretation of the Fe~K line in terms of  reflection from a relativistic black-hole disk has been   found to be rather robust, but it has not been possible to fully explore the model parameter space.  In particular, constraints on $a$ have only been obtained by fitting `inverted photon data' which is made from the real data in a model-dependent way that already assumes a certain value of $a$. The puzzling stability of this line has been attributed, amongst other things, to general relativity effects very close to the black hole (Miniutti et al.\\ 2003; Fabian \\& Vaughan 2003). It is thus interesting to know if and how a generalisation of the standard disk-line scheme together with accurate computation of general relativistic effects  can add new pieces of information or whether a more substantial  modification of the whole picture is needed (for further  discussion and references, see e.g.\\ Weaver \\& Yaqoob 1998; Krolik  1999; Hartnoll \\& Blackman 2001; Dumont et al.\\ 2002). Different kinds of such generalisations have been proposed. Previous papers indicated that details of the model do matter. However, a strong gravitational field has  almost invariably been required. In the present paper, we adopt the assumption that the line and continuum emission are produced by an irradiated disk  near a rotating black hole, we relax some of the previous assumptions, and we search for best-fitting parameters of the model.  In the \\S~2 we describe the main features of the new computational tool. We list several variants of the code which address different problems of fitting X-ray spectra of black hole plus accretion disk systems.  In \\S~3 we use the code to analyse time-averaged data for MCG--6-30-15 {\\it{}XMM--Newton} CCD data as detailed in Fabian et al.\\ 2002).  We point out that more complex models may not require a large value of the black-hole  angular momentum, although the current data are consistent with a maximally  rotating black hole. Finally, in \\S~4, we summarise our results and state our conclusions.  \\begin{table*}[tbh] \\caption{Basic features of the new model in comparison with other black hole disk-line models.} \\begin{center} \\begin{small} \\tabletypesize{\\scriptsize} \\begin{tabular}{lccccl} \\tableline \\multicolumn{1}{c}{Model} & \\multicolumn{4}{c}{Effects that are taken into account\\rule[-1.5ex]{0mm}{4.5ex}} & \\multicolumn{1}{c}{Reference}  \\\\ \\cline{2-5} & \\rule[-3.5ex]{0mm}{8ex}\\parbox{29mm}{\\footnotesize{}Energy shift of photons/Lensing effect} & \\parbox{18mm}{\\footnotesize{}Black hole angular momentum} & \\parbox{18mm}{\\footnotesize{}Axisymmetry is assumed} & \\parbox{20mm}{\\footnotesize{}Steady source is assumed} &  \\\\ \\tableline {\\sf{}diskline}   \\rule{0mm}{3ex}  &  yes/no~       & $0$~                 & yes          & yes & Fabian et al.\\ (1989) \\\\ {\\sf{}laor}                        &  yes/yes       & $0.998$~             & yes          & yes & Laor (1991) \\\\ {\\sf{}kerrspec}                    &  yes/yes       & $\\langle0,1\\rangle^{\\dag}$   & no~          & yes & Martocchia et al.\\ (2000) \\\\ {\\sf{}ky}                          &  yes/yes       & $\\langle0,1\\rangle$           & no$^{\\ddag}$ & no  & This paper \\\\ \\tableline \\end{tabular} \\end{small} \\vspace*{2mm}\\par{} {\\parbox{0.9\\textwidth}{\\footnotesize{}$^{\\dag}$~The value of dimension-less $a$ parameter is kept frozen.\\\\ $^{\\ddag}$~A one-dimensional version is available for the case of an axisymmetric disk. In this axisymmetric mode, {\\sf{}ky} still allows $a$ and other relevant parameters to be fitted (in which case the computational speed of {\\sf{}ky} is then comparable to {\\sf{}laor}). The results can be more accurate than those obtained with other routines because of the ability to tune the grid resolution.}} \\label{tab:models} \\end{center} \\end{table*}  \\begin{table}[!tbh] \\caption{Basic versions of the model.} \\begin{center} \\begin{footnotesize} \\tabletypesize{\\footnotesize} \\begin{tabular}{l@{~}p{22mm}@{~}p{36mm}} \\tableline Name \\rule[-1.5ex]{0mm}{4.5ex} & Type$^{\\dag}$~ & Usage \\\\ \\tableline {\\sf{}KYGline} & additive \\rule{0mm}{3ex} & Relativistic spectral line from a black hole disk. \\\\ {\\sf{}KYHrefl} & additive  & Compton reflection with an incident power-law (or a broken power-law) continuum. \\\\ {\\sf{}KYLcr} & additive & Lamp-post Compton reflection model. \\\\ {\\sf{}KYSpot} & To be used as an independent code outside {\\sc{}xspec} & Time-dependent spectrum of a pattern co-orbiting with the disk. \\\\ {\\sf{}KYConv}  & convolution & Convolution of the relativistic kernel with intrinsic emissivity across the disk. \\\\ \\tableline \\end{tabular} \\end{footnotesize} \\vspace*{2mm}\\par{} {\\vspace*{-1mm}\\parbox{0.45\\textwidth}{\\footnotesize{}$^{\\dag}$~Different model types correspond to {\\sc{}xspec} syntax and are defined by the way they act in the overall model and form the final spectrum. According to the usual convention in {\\sc{}xspec}, additive models represent individual {\\it emission} spectral components which may originate e.g.\\ in different regions of the source. Additive models are simply superposed in the total signal. Multiplicative components (e.g.\\ {\\sf{}hrefl}, discussed in the Appendix~\\ref{appb}) multiply the current model by an energy-dependent factor. Convolution models modify the model in a more non-trivial manner. See Arnaud (1996) for details.}} \\label{tab:versions} \\end{center} \\end{table} ",
        "conclusions": "We have presented an extended computational scheme which can be conveniently employed to examine predicted X-ray spectral features from black-hole accretion disks. As mentioned above, a one-dimensional version of the code basically reproduces previous methods, with the addition of the option to vary more  parameters (e.g.\\ $a$). The current two-dimensional  version also accomodates non-axisymmetric and time-evolving emissivity from a geometrically thin disk. The new tool facilitates accurate comparisons  between model spectra and actual data. It provides also  polarimetric information. One can even substitute other tables for the Kerr metric tables to explore different specetime. However, we defer detailed discussion of  these latter capabilities to a future paper. Furthermore, work is in progress to include realistic emissivities relevant to different physical situations, in order to allow for the intrinsically three-dimensional  geometry of the sources (e.g.\\ an optically-thin corona above the disk), and to parallelize the code.  Recently, several research groups have embarked on projects to compare theoretical predictions of intrinsic emissivity of accretion disks with observational data obtained by X-ray satellites. This approach offers a fascinating opportunity to explore processes near black-hole sources, both AGNs and BHCs, and to estimate physical parameters of black holes themselves. In order to limit the number of unknown model parameters and to alleviate the computational burden of the fit procedure, simplifications on the model emissivity (for example, stationarity and/or axial symmetry) have been often imposed.\\footnote{Traditionally, Occam's razor is invoked to justify simplifying assumptions. However, this way of reasoning is not sufficiently substantiated here in view of the fact that accretion flows are normally found to be very turbulent and fluctuating rather than smooth and steady. Put in other words, the assumption of an axisymmetric and stationary disk may provide a successful fit to the time averaged spectrum, but it is not satisfactory because we know that more complicated models {\\it{}must} be involved on physical grounds. (``Things should be made as simple as possible---but not too simple\\ldots'')} Naturally, fitting data with complex models is computationally more demanding.  To summarize, we have illustrated the code by following common practice and employing minimisation of $\\chi^2$ to find the best-fitting parameter values for the highest signal-to-noise Fe K line profile available to date---from the {\\it{}XMM--Newton} campaign of the Seyfert~1 galaxy MCG--6-30-15. The new model can vary more of the parameters than previous variants of the disk-line scheme allowed, which seems to be necessary in view of the fact that simple prescriptions for intrinsic emisivity (axisymmetric, radially monotonic, etc.) are not adequate to reflect realistic simulations of accretion flows. It is without much surprise that a complex model exhibits an intricate $\\chi^2$ space in which ambiguities arise with multiple islands of acceptable parameters. The mean {\\it{}XMM--Newton} spectrum of MCG--6-30-15 analyzed here justifies rejection of the often-used plain power-law emissivity that depends solely on radius from the black hole. However, we cannot arbitrate the issue of non-monotonic emissivity versus a non-axisymmetric one. Due to the inherent degeneracy in time-averaged spectra, both alternatives can reproduce these data with reasonable parameters. In other words, $a$ of the black hole cannot be unambiguously determined if complex (realistic) emissivities are adopted.  It is clear that the relativistic line profiles affected by strong gravity are such that it is very difficult to recover unambiguous information about the key parameters, such as black-hole spin, disk inclination angle, and inner disk radius. This is because even with increased throughput, energy resolution and signal-to-noise, the smooth, featureless tails of the time-averaged line profiles do not retain sufficient information to break the degeneracies between the dependences on the black-hole, accretion disk, and line emissivity parameters.  Higly specific spectral shapes could actually be provided by contributions to the line profile from higher-order images, whose form inherits unique information about space-time of the black hole. It may appear rather unlikely that higher-order images could be discovered in spectra of AGNs, because this would require to see an unobscured inner disk at large (almost edge-on) inclination.  As the data improve, the important measurable quantities that will constrain parameters when compared with the models are: (i)~the slope of the red wing, (ii)~the energy of the cut-off of the blue side of the line profile, and (iii)~the detailed shape of the blue peak (see Figs.\\ \\ref{fig:example0}--\\ref{fig:example3} and further discussion in the Appendix). The extreme part of the red wing will not be as useful because it will always be difficult to tell where the line emission joins the continuum no matter how good the data are, and also because there is often no prominent sharp feature in the red wing since the red Doppler peak is too smeared out when the line emission extends too close to the black hole. In addition, time variability will be extremely important. No matter how, or if, the line emission responds to the continuum, we know that the disk inclination and black-hole spin very likely do not change during different snapshots. Therefore any differences in the snapshot line profiles must be due only to changes in the line emissivity, in terms of its radial and azimuthal distribution and/or its radial extent ($r_{\\rm{}in}$ and $r_{\\rm{}out}$). Combining these approaches and collecting enough snapshots may provide sufficient information to uniquely determine one of the parameters, such as black-hole spin."
    },
    "0403/astro-ph0403588_arXiv.txt": {
        "abstract": "We present a comprehensive space--based study of ten X--ray luminous galaxy clusters ($L_X{\\ge}8{\\times}10^{44}{\\rm erg \\,s^{-1}}$[0.1--2.4\\,keV]) at $z{=}0.2$.  \\emph{Hubble Space Telescope (HST)} observations reveal numerous gravitationally--lensed arcs for which we present four new spectroscopic redshifts, bringing the total to thirteen confirmed arcs in this cluster sample.  The confirmed arcs reside in just half of the clusters; we thus obtain a firm lower limit on the fraction of clusters with a central projected mass density exceeding the critical density required for strong--lensing of $50\\%$. We combine the multiple--image systems with the weakly--sheared background galaxies to model the total mass distribution in the cluster cores ($R{\\le}500{\\rm kpc}$).  These models are complemented by high--resolution X--ray data from \\emph{Chandra} and used to develop quantitative criteria to classify the clusters as relaxed or unrelaxed.  Formally, $(30{\\pm}20)\\%$ of the clusters form a homogeneous sub--sample of relaxed clusters; the remaining $(70{\\pm}20)\\%$ are unrelaxed and are a much more diverse population. Most of the clusters therefore appear to be experiencing a cluster--cluster merger, or relaxing after such an event.  We also study the normalization and scatter of scaling relations between cluster mass, luminosity and temperature.  The scatter in these relations is dominated by the unrelaxed clusters and is typically ${\\sigma}{\\simeq}0.4$.  Most notably, we detect 2--3\\,times more scatter in the mass--temperature relation than theoretical simulations and models predict.  The observed scatter is also asymmetric -- the unrelaxed systems are systematically 40\\% hotter than the relaxed clusters at 2.5--${\\sigma}$ significance.  This structural segregation should be a major concern for experiments designed to constrain cosmological parameters using galaxy clusters.  Overall our results are consistent with a scenario of cluster--cluster merger induced boosts to cluster X--ray luminosities and temperatures. ",
        "introduction": "\\setcounter{footnote}{7}  \\normalsize  Massive galaxy clusters are the largest collapsed structures in the Universe ($M_{\\rm virial}{\\simeq}10^{15}{\\rm M_\\odot}$), containing vast quantities of the putative dark matter (DM), hot intracluster gas (${\\tx}{\\simeq}7{\\rm keV}$), and galaxies ($n_{\\rm gal}{\\sim}10^3$). These rare systems stand at the nodes of the ``cosmic web'' as defined by the large--scale filamentary structure seen in both galaxy redshift surveys (e.g.\\ de~Lapparent, Geller \\& Huchra 1986; Shectman et al.\\ 1996; Vettolani et al.\\ 1997; Peacock et al.\\ 2001; Zehavi et al.\\ 2002) and numerical simulations of structure formation (e.g.\\ Bond et al.\\ 1996; Yoshida et al.\\ 2001; Evrard et al.\\ 2002).  Clusters are inferred to assemble by accreting matter along the filamentary axes, slowly ($t_{\\rm crossing}{\\sim}2$--3\\,Gyr) ingesting DM, gas and stars into their deep gravitational potential wells.  Clusters have long been recognized as cosmological probes.  For example, the evolution of cluster substructure with look--back--time is in principal a powerful diagnostic of the cosmological parameters (Gunn \\& Gott 1972; Peebles 1980; Richstone, Loeb \\& Turner 1992; Evrard et al.\\ 1993).  A complementary probe is to constrain the matter density of the universe and the normalization of the matter power spectrum using the cluster mass function.  However, it is currently not possible to measure the cluster mass function directly. More easily accessible surrogates such as the X--ray luminosity and temperature functions are therefore used in combination with scaling relations between the relevant quantities (e.g.\\ Eke et al.\\ 1996; Reiprich \\& Boehringer 2002; Viana, Nichol \\& Liddle 2002; Allen et al.\\ 2003).  A critical component of such analyses is the precision to which the scaling relations are known.  Samples of X--ray selected clusters are now of a sufficient size that systematic uncertainties may be comparable with the statistical uncertainties, and therefore deserve careful analysis before robust cosmological conclusions may be drawn (e.g.\\ Smith et al.\\ 2003).  Measurements of the Sunyaev--Zeldovich effect (SZE) are also emerging as a powerful cosmological tool (Carlstrom, Holder \\& Reese 2002).  Cosmological SZE surveys will rely on the cluster mass--temperature relationship in a similar manner to cosmological X--ray surveys.  Such experiments may therefore also be compromised if astrophysical systematics are identified and carefully eliminated from the analysis (e.g.\\ Majumdar \\& Mohr 2003).  Detailed study of the assembly and relaxation histories of clusters, and their global scaling relations as a function of redshift are therefore vitally important.  To advance our understanding of the assembly, relaxation and thermodynamics of massive galaxy clusters requires information about the spatial distribution of DM, hot gas and galaxies in clusters. Several baryonic mass tracers are available, for example X--ray emission from the intracluster medium (hereafter ICM -- e.g.\\ Jones \\& Forman 1984; Buote \\& Tsai 1996; Schuecker et al.\\ 2001) and the angular and line--of--sight velocity distribution of cluster galaxies (e.g.\\ Geller \\& Beers 1982; Dressler \\& Shectman 1988; West \\& Bothun 1990).  These diagnostics have often been used as surrogates for a direct tracer of the underlying DM distribution.  The major drawback of this approach is the requirement to assume a relationship between the luminous and dark matter distributions (e.g.\\ that the ICM is in hydrostatic equilibrium with the DM potential) -- it is precisely these assumptions that require detailed testing.  This issue is further aggravated by the expectation that cluster mass distributions are DM dominated on all but the smallest scales (e.g.\\ Smith et al.\\ 2001; Sand, Treu \\& Ellis 2002; Sand et al.\\ 2004).  Gravitational lensing offers a solution to much of this problem, in that the lensing signal is sensitive to the total mass distribution in the lens, regardless of its physical nature and state.  Detailed study of gravitational lensing by massive clusters is therefore an important opportunity to gain an empirical understanding of the distribution of DM in clusters.  Indeed, early comparisons between X--ray and lensing--based mass measurements revealed a factor 2--3 discrepancy between X--ray and strong--lensing--based cluster mass estimates (e.g.\\ Miralda--Escud\\'e \\& Babul 1995; Wu \\& Fang 1997), although the agreement between weak--lensing and X--ray measurements was generally better, albeit within large uncertainties (e.g.\\ Squires et al.\\ 1996, 1997; Smail et al.\\ 1997).  The simplifying assumptions involved in the X--ray analysis were soon identified as the likely dominant source of this discrepancy; this was confirmed by several authors (e.g.\\ Allen 1998; Wu et al.\\ 1998; Wu 2000).  In summary, X--ray and lensing mass measurements for the most relaxed clusters agree well if the multi--phase nature of the ICM in cool cores (e.g.\\ Allen et al.\\ 2001) is incorporated into the X--ray analysis.  The situation is more complex in more dynamically disturbed clusters, with larger discrepancies being found at smaller projected radii.  The origin of the X--ray versus lensing mass discrepancy in clusters that do not contain a cool core is generally attributed to the simplifying equilibrium and symmetry assumptions of the X--ray analysis.  For that reason, most modern X--ray cluster analyses that involve measuring cluster mass using only X--ray data understandably concentrate on relaxed, cool core clusters (e.g.\\ Allen et al.\\ 2001).  An important caveat to adopting lensing as the tool of choice to measure cluster mass is that lensing actually constrains the projected mass distribution along the line--of--sight to the cluster.  The addition of three--dimensional information into lensing studies may therefore be important before final conclusions are drawn.  For example Czoske et al.'s (2001; 2002) wide--field redshift survey of Cl\\,0024$+$1654 at $z{=}0.395$ revealed that this previously presumed relaxed strong--lensing cluster (e.g.\\ Smail et al.\\ 1996; Tyson et al.\\ 1998) cluster is not relaxed, and appears to have suffered a recent merger along the line--of--sight (see also Kneib et al.\\ 2003).  Early gravitational lensing studies of galaxy clusters concentrated on individual clusters selected because of their prominent arcs (e.g.\\ Mellier et al.\\ 1993; Kneib et al.\\ 1994, 1995, 1996; Smail et al.\\ 1995a, 1996; Allen et al.\\ 1996; Tyson et al.\\ 1998).  This ``prominent arc'' selection function was vital to development of the techniques required to interpret the gravitational lensing signal (e.g.\\ Kneib 1993).  However it also made it difficult to draw conclusions about galaxy clusters as a population of astrophysical systems from these studies.  Smail et al.\\ (1997) used the \\emph{Hubble Space Telescope (HST)} to study a larger sample of optically rich clusters, selected originally for the purpose of studying the cluster galaxies.  As X--ray selected samples became available in the late--1990's, Luppino et al.\\ (1999) also searched for gravitational arcs in ground--based imaging of 38 X--ray luminous clusters.  The broad conclusions to emerge from these pioneering studies were that to use gravitational lensing to learn about clusters as a population, a selection function that mimics mass--selection as closely as possible, and the high spatial resolution available from \\emph{HST} imaging are both key requirements.  We are conducting an \\emph{HST} survey of an objectively selected sample of ten X--ray luminous (and thus massive) clusters at $z{\\simeq}0.2$ (Table~\\ref{sample}, \\S\\ref{design}).  Previous papers in this series have presented (i) a detailed analysis of the density profile of A\\,383 (Smith et al.\\ 2001), (ii) a search for gravitationally--lensed Extremely Red Objects (EROs -- Smith et al.\\ 2002a) and (iii) near--infrared (NIR) spectroscopy of ERO\\,J003707, a multiply--imaged ERO at $z{=}1.6$ behind the foreground cluster A\\,68 (Smith et al.\\ 2002b).  This paper describes the gravitational lensing analysis of all ten clusters observed with \\emph{HST} and uses the resulting models of the cluster cores to measure the mass and structure of the clusters on scales of $R{\\le}500\\,{\\rm kpc}$.  We also exploit archival \\emph{Chandra} observations and NIR photometry of likely cluster galaxies to compare the distribution of total mass in the clusters with the gaseous and stellar components respectively. This combination of strong--lensing, X--ray and NIR diagnostics enable us to quantify the prevalence of dynamical immaturity in the X--ray luminous population at $z{\\simeq}0.2$ and to calibrate the high--mass end of the cluster mass--temperature relationship.  We outline the organization of the paper.  In \\S\\ref{design} we describe the survey design and sample selection.  We then describe the reduction and analysis of the optical data in \\S\\ref{optical}, comprising the \\emph{HST} imaging data (\\S\\ref{hst}) and new spectroscopic redshift measurements for arcs in A\\,68 and A\\,2219 (\\S\\ref{spec}).  The end--point of \\S\\ref{optical} is a definition of the strong-- and weak--lensing constraints available for all ten clusters.  We apply these constraints in \\S\\ref{modelling}, to construct detailed gravitational lens models of the cluster potential wells; the details of the modelling techniques are described in the Appendix, and the process of fitting the constraints in each cluster are described in \\S\\ref{modelling}.  We then complement these gravitational lensing results with observations of the X--ray emission from the clusters' ICM, drawn from the \\emph{Chandra} data archive (\\S\\ref{xray}).  The main results of the paper are then presented in \\S\\ref{results}, including measurements of the mass and maturity of the clusters and a detailed study of the cluster scaling relations. We discuss the interpretation of the results in \\S\\ref{discussion} and briefly assess their impact on attempts to use clusters as cosmological probes.  Finally, we summarize our conclusions in \\S\\ref{conclusions}  We assume a spatially flat universe with $H_0{=}50{\\rm km\\,s^{-1}Mpc^{{-}1}}$ and $q_0{=}0.5$; in this cosmology $1''{\\equiv}4.2\\,{\\rm kpc}$ at $z{=}0.2$.  Our main results are insensitive to this choice of cosmology, for example, the cluster mass measurements would be modified by ${\\ls}10\\%$ if we adopted the currently favored values of $\\Omega_{\\rm M}{=}0.3$, $\\Omega_\\Lambda{=}0.7$, $H_0{=}65{\\rm km\\,s^{-1}Mpc^{{-}1}}$.  We also adopt the complex deformation, $\\vec\\tau{=}\\tau_x{+}i\\tau_y{=}|\\vec\\tau| e^{2i\\theta}$, as our measure of galaxy shape when dealing with the weak lensing aspects of our analysis, where $\\tau{=}(a^2{+}b^2)/2ab$ and $\\theta$ is the position angle of the major axis of the ellipse that describes each galaxy.  We define the terms ``ellipticity'' to mean $\\tau$ and ``orientation'' to mean $\\theta$.  All uncertainties are quoted at the 68\\% confidence level. ",
        "conclusions": "We have undertaken a comprehensive study of the distribution of mass in ten X--ray luminous ($L_X{\\ge}8{\\times}10^{44}{\\rm erg \\,s^{-1}}$[0.1--2.4\\,keV]) galaxy clusters at $z{=}0.21{\\pm}0.04$.  The cornerstone of our analysis is a suite of detailed gravitational lens models that describe the distribution of total mass in the cluster cores.  These models are constrained by the gravitational lensing signal detected in high--resolution \\emph{HST}/WFPC2 imaging of the clusters, including numerous multiply--imaged and weakly--sheared background galaxies.  Analysis of archival \\emph{Chandra} observations complements the lensing analysis and enables us to relate the total mass and structure of the clusters to the thermodynamics of the intra--cluster medium. We re--cap the key results.  \\setcounter{fred}{0} \\begin{list}{(\\roman{fred})}{\\usecounter{fred}\\setlength{\\itemindent}{0mm}\\setlength{\\labelwidth}{15mm}\\setlength{\\labelsep}{2mm}\\setlength{\\leftmargin}{7mm}} \\setlength{\\itemsep}{1.0mm}  \\item Five of the ten clusters contain spectroscopically confirmed strong gravitational lensing, i.e.\\ multiply--imaged background galaxies.  These five clusters comprise: A\\,68 (Smith et al.\\ 2002b; \\S\\ref{spec}), A\\,383 (Smith et al.\\ 2001; Sand et al.\\ 2004), A\\,963 (Ellis et al.\\ 1991), A\\,2218 (Pell\\'o et al.\\ 1992; Ebbels et al.\\ 1998; Ellis et al.\\ 2002) and A\\,2219 (\\S\\ref{spec}).  \\item Of the remaining five clusters, two contain relatively unambiguous examples of strong lensing for which spectroscopic redshifts are not yet available (A\\,267 and A\\,1835).  The other three clusters, A\\,209, A\\,773 and A\\,1763 do not contain any obviously multiply--imaged galaxies, however the optical richness and massive nature of A\\,773 imply that this cluster may well contain strong lensing that has yet to be uncovered.  \\item Based on our search for strong--lensing in these clusters down to a surface brightness limit of $\\mu_{702}{\\simeq}25$\\,mag\\,arcsec$^{{-}2}$, we therefore put a firm lower limit on the fraction of the cluster sample that have a central projected mass density in excess of the critical density required for gravitational lensing of $50\\%$.  Including A\\,267 and A\\,1835 increases this limit to $70\\%$.  \\item We use the strong-- and weak--lensing signals to constrain parametrized models of the cluster potential wells, and from these models compute maps of the total projected mass in the cluster cores.  Spatial analysis of these maps reveals that four of the clusters form a homogeneous sub--sample with very high central mass fractions (${\\Mcen}/{\\Mtot}{>}0.95$).  The remaining six are strongly heterogeneous, with central mass fractions in the range $0.4{\\le}{\\Mcen}/{\\Mtot}{\\le}0.9$.  The central mass fraction of ${\\Mcen}/{\\Mtot}{\\simeq}0.95$ that divides these two populations corresponds to a $K$--band central luminosity fraction of $L_{\\rm K,BCG}/L_{\\rm K,tot}{\\sim}0.5$.  \\item All of the six low central mass fraction clusters have an irregular, but not obviously bi/tri--modal X--ray morphology. Four of the six are constrained by the current lensing data to have a bi/tri--modal mass morphology (A\\,68, A\\,773, A\\,2218, A\\,2219).  The other two (A\\,209 and A\\,1763) may be merging in the plane of the sky and thus any multi--modality in their mass distributions is not well--sampled by our WFPC2 pencil--beam survey of the cluster cores.  \\item Three of the four high central mass fraction clusters also have relaxed X--ray morphologies.  The remaining cluster (A\\,267) has a disturbed X--ray morphology, with a ${\\sim}90$\\,kpc offset between its centers of X--ray emission and mass.  The distribution of mass in this cluster may therefore be more complex than the single dark matter halo (plus cluster galaxies) that the current data are able to constrain.  \\item Combining all of the information available to us, we define ``relaxed'' clusters to be those which appear relaxed in all available diagnostics, with the exception that we do not require a cool core.  Quantitatively relaxed clusters therefore have a single cluster--scale DM halo in their lens model ($N_{\\rm DM}{=}1$), a high central mass fraction (${\\Mcen}/{\\Mtot}{\\ge}0.95$) and central $K$--band luminosity fraction ($L_{K,{\\rm BCG}}/L_{K,{\\rm tot}}{\\gs}0.5$), no evidence for an offset between the X--ray emission and the center of mass ${\\drpeak}{<}4{\\rm kpc}$) and the X--ray morphology is either circular, or mildly elliptical.  The unrelaxed clusters do not meet at least one of these criteria.  \\item Applying these criteria to the cluster sample, we conclude that seven of the ten clusters are dynamically immature, i.e.\\ unrelaxed (A\\,68, A\\,209, A\\,267, A\\,773, A\\,1763, A\\,2218, A\\,2219) and three are relaxed (A\\,383, A\\,963, A\\,1835); thus, formally $70{\\pm}20\\%$ of X--ray luminous cluster cores at $z{=}0.2$ are unrelaxed.  \\item We detect a factor of three more scatter in the mass--temperature plane than predicted by Evrard, Metzler \\& Navarro (1996), implying that great care should be exercised when using such relations to convert cluster temperature functions to mass functions in pursuit of cosmological parameters.  We also consider a number of key uncertainties that may artificially inflate our estimate of the scatter.  This exercise suggests that approximately one third of the scatter detected in this study may be due to issues related to the small field--of--view of our WFPC2 observations.  \\item The scatter in the mass--temperature plane is asymmetric, presenting evidence of structural segregation.  The normalization of the mass--temperature relation for unrelaxed (dynamically immature) clusters is $40\\%$ hotter than for relaxed clusters at 2.5--${\\sigma}$ significance.  This result is consistent with recent simulations of cluster--cluster mergers (Ricker \\& Sarazin 2001; Randall, Sarazin \\& Ricker 2002), implying that merger induced temperature boosts may be the dominant factor behind the hotter normalization of unrelaxed systems.  \\end{list}  In summary, this study is the first of its kind, exploiting detailed strong--lensing constraints on the distribution of mass in X--ray luminous cluster cores, complemented by X--ray spectro--imaging with \\emph{Chandra}.  The high frequency of dynamical immaturity, coupled with the structural segregation of the clusters in the mass--temperature plane have profound implications for our understanding of how clusters form and evolve.  Perhaps of greatest importance at this time is the implications of these results for using clusters to constrain the cosmological parameters, ${\\Omega}_M$, ${\\sigma}_8$ and the dark energy equation--of--state parameter $w$. In a companion paper we demonstrate that inadequate knowledge of the cluster selection function can lead to $20\\%$ systematic errors in cluster--based measurements of ${\\sigma}_8$.  Turning to $w$, forthcoming Sunyaev--Zeldovich Effect experiments designed to detect and measure the mass of clusters out to high redshifts, using mass--temperature scaling relations may be compromised by unidentified and/or poorly calibrated astrophysical systematic uncertainties (see also Majumdar \\& Mohr 2003).  Our future program will build on these results in three ways.  First, we aim to overcome the principal uncertainties in the current work: small number statistics and tiny field--of--view.  Wide--field space--based imaging of a statistically complete sample of clusters would be essential to achieve this goal.  Second, detailed comparison of selection effects and measurement techniques between theoretical and observational studies will enable more detailed and rigorous comparison between observational and synthetic datasets.  Finally, we are gathering \\emph{HST}/ACS imaging of an identically selected sample of clusters at $z{\\simeq}0.55$ drawn from the MACS sample (Ebeling et al.\\ 2001b).  We will combine these data with observations in the X--ray pass--band and compare the results to those found here.  The ${\\sim}3$\\,Gyr difference in look--back--time between $z{=}0.2$ and $z{=}0.55$ will enable us to search for evolutionary trends in the most massive clusters."
    },
    "0403/astro-ph0403294_arXiv.txt": {
        "abstract": "Using the spectroscopic sample of the SDSS DR1 we measure how gas was transformed into stars as a function of time and stellar mass: the baryonic conversion tree (BCT). There is a clear correlation between early star formation activity and present-day stellar mass: the more massive galaxies have formed about 80\\% of their stars at $z>1$, while for the less massive ones the value is only about 20\\%. By comparing the BCT to the dark matter merger tree, we find indications that star formation efficiency at $z>1$ had to be about a factor of two higher than today ($\\sim 10\\%$) in galaxies with present-day stellar mass larger than $2 \\times 10^{11}M_\\odot$, if this early star formation occurred in the main progenitor. Therefore, the LCDM paradigm can accommodate a large number of red objects. On the other hand, in galaxies with present-day stellar mass less than $10^{11}$ M$_{\\odot}$, efficient star formation seems to have been triggered at $z \\sim 0.2$. We show that there is a characteristic mass (M$_* \\sim 10^{10}$ M$_{\\odot}$) for feedback efficiency (or lack of star formation). For galaxies with masses lower than this, feedback (or star formation suppression) is very efficient while for higher masses it is not. The BCT, determined here for the first time, should be an important observable with which to confront theoretical models of galaxy formation. ",
        "introduction": "In the current paradigm for galaxy formation, galaxies form in cold dark matter halos, which evolve from small, primordial, Gaussian fluctuations, by gravitational instability.  This mechanism fits well in the successful LCDM picture which correctly describes the Universe from the Cosmic Microwave Background at $z=1088$ to local galaxy clustering \\citep{SpergelWMAP03,Percival2dF03}.  One of the strong predictions of this galaxy formation paradigm is the typical redshift of dark matter halo formation (i.e.  virialization) as a function of halo mass and cosmological parameters (e.g., \\citet{ST04}).  Given that cosmological parameters have been tightly constrained (e.g., \\cite{SpergelWMAP03}), we can reconstruct the average dark halo formation history as a function of mass (e.g., \\citet{PS74,ST99,ST04}).  Naively, the dark matter halo collapse, i.e. virialization, should trigger baryonic gas transformation into stars; in addition, subsequent dark matter mergers should produce star formation episodes. We will show that this simple model is not in agreement with observations.  All we can observe is the integrated light of galaxies' stellar population; thus, to compare the theory prediction for the dark matter halos with observations, the process of how baryons are transformed into stars needs to be simulated, either through semi-analytical recipes or by means of hydro-dynamic N-body simulations. Since no theory of star formation has been yet established, we do not have a fundamental theory that allows us to compute from basic principles the star formation efficiency. Further, complications arise from other phenomena, such as feedback by stars and AGN that prevent the formation of giant molecular clouds and therefore reduce star formation efficiency. Given the complexity of the task of learning about galaxy formation from numerical simulations, here we take the complementary approach of placing new observational constraints on the stellar assembly history as a function of galaxy mass.  In this paper we use about $10^5$ galaxies from the Sloan Digital Sky Survey Data Release 1 (SDSS DR1) to determine, for the first time, the amount of baryons that have been transformed into stars as a function of total stellar mass and time.  This allows us to build the baryonic conversion tree (BCT), which can then be compared with the dark matter merger tree. We present such comparison  and show how it can be used to compute the star formation efficiency and the relative importance of feedback (or lack of star formation).  Our main findings are: \\begin{enumerate} \\item There is a clear correlation between total stellar mass of the galaxy and the fraction of gas transformed into stars at $z \\geq 1$. The larger the mass the larger the fraction of gas transformed into stars at high redshift.  \\item If large galaxies need to be formed by $z \\sim 1$ in a ``monolithic'' fashion, as observations suggest (e.g., \\citet{Bower+92,PJDWSSDW98,Lilly+98,BE00,Im+02,Renzini03,Gao+03,Glazebrook+04}), high-redshift star formation efficiency needs to be much higher (about 10\\%) in massive galaxies than in less massive ones.  This high star formation activity at early times means the build-up of stellar mass does not follow the hierarchical build-up of total mass. Stars could form in smaller objects (not in the main progenitor only) with lower efficiency, provided that the galaxy formation process has some way to synchronise star formation in disparate pieces.  The reason being that the above mentioned galaxies at $z\\sim 1-2$ have stellar populations with almost no spread in their stellar ages and derived star formation histories consistent with very short times ($< 0.1$ Gyr, see Macarthy et al 2004).  \\item There is an indication that major mergers are not the principal drivers of star formation.  \\item We propose that a threshold for star formation for galaxies with masses $\\sim 10^{10}$ M$_{\\odot}$ can explain the findings above. The existence of such a threshold at low redshift is well documented in the literature \\citep{MK01}, and can be linked to feedback efficiency. Feedback, i.e. suppression of star formation activity, is expected to be very inefficient in massive galaxies, while efficient in less massive galaxies.  \\end{enumerate} ",
        "conclusions": "We have determined for the first time the baryonic conversion tree for galaxies.  We were able to do this with observations at $z<0.3$, using 96545 galaxies from the SDSS DR1 spectroscopic sample.  In the hierarchical structure formation model, massive {\\em dark halos} form (i.e. virialize) later than smaller halos, from mergers of smaller units (e.g. \\citet{lc93,lc94,ljl03}).  This model has been thoroughly tested against numerical cold dark matter simulations. Naively one could expect that the dark matter halo collapse should trigger baryonic gas transformation into stars; in addition subsequent dark matter mergers should produce star formation episodes.  Instead, for the the stellar assembly history we find that: the more massive galaxies have old stellar population and massive, old elliptical galaxies are already in place at $z \\sim 1$.  This has been known for a long time and sometimes it is referred to as ``down-sizing'', \\citep{CSC96}.  We find that massive galaxies have transformed more gas into stars at higher redshift (in agreement with high z observations e.g., \\cite{kodama04}), and then star formation was suppressed, while less massive galaxies transform more gas into stars at low redshift.  So one should not be surprised to see abundant red objects at high redshift in the LCDM paradigm, these objects can form in virialised halos if star formation efficiency is high.  Indeed, \\cite{JFDTPN99} have shown that single-halo hydro-dynamical models would require an increased star formation efficiency for more massive galaxies, higher than the few percent found today in giant molecular clouds, in agreement with the value determined from the ``fossil record''  in the present work. Our findings, based only on observations at $z < 0.35$ (the ``fossil record''), are in agreement with a suite of independent, $z>1$, observations: observations of old elliptical galaxies at high redshift \\citep{D+96,Spinrad+97,NDJH03}, indications that elliptical galaxies are already formed at $z>1$ (e.g. \\citet{Bower+92,PJDWSSDW98,Lilly+98,BE00,Im+02,Renzini03,Gao+03,Glazebrook+04}). It also nicely explains the tightness of the colour-magnitude relation \\citep{Bower+92}.  On the other hand, we find that small galaxies seem to accrete mass passively at early times and see a lot of their gas reservoir transformed into stars at late times.  Since the probability distribution for dark halo mergers peaks at low redshift for massive halos and high redshift for small halos, we conclude that dark matter mergers and star formation are not correlated.  We speculate that one possibility to explain the apparent and illusory anti-hierarchical nature of the stellar assembly history is the existence of a threshold for star formation: once the threshold is crossed all available baryons are turned into stars (``infall model'') and afterward galaxies approximatively evolve passively. The threshold is met at very early time for massive galaxies and a later time for less massive ones (see e.g. \\citet{HJ99}).  A star formation threshold has been observed in disk galaxies by \\citet{MK01} and there has been some recent additional evidence from the formation of dust lanes in disk galaxies \\citep{DYB04} that this threshold may take place at $V_c \\sim 100$ km s$^{-1}$, in agreement with the findings of \\citet{VOJ02} and \\citet{kannappan04} who also found a transition at about $100$ km s$^{-1}$ for star formation efficiency. This $V_c$ value corresponds to the characteristic mass found here ($M_* \\sim 10^{10}$ M$_{\\odot}$), that defines the border between efficient and inefficient star formation.  This characteristic mass has been related to feedback efficiency threshold (e.g., Dekel \\& Silk 1986, Dekel \\& Woo 2003 and references therein).  As we do not yet have a fundamental theory for galaxy formation and given the complexity involved in studying the process with hydrodynamic-N body simulations, we hope that this new determination of the baryonic conversion history will be a useful observable to gauge galaxy formation models against.\\\\"
    },
    "0403/astro-ph0403407_arXiv.txt": {
        "abstract": "We show how the viscous evolution of Keplerian accretion discs can be understood in terms of simple kinetic theory. Although standard physics texts give a simple derivation of momentum transfer in a linear shear flow using kinetic theory, many authors, as detailed by Hayashi \\& Matsuda 2001, have had difficulties applying the same considerations to a circular shear flow. We show here how this may be done, and note that the essential ingredients are to take proper account of, first, isotropy locally in the frame of the fluid and, second, the geometry of the mean flow. ",
        "introduction": "Accretion discs play a central role in a wide range of astronomical environments, mediating the gas flows in the vicinity of object as diverse as AGN, binary stars and protostars (Pringle, 1981). In an accretion disc, the predominant flow is a circular shear flow, with angular velocity $\\Omega(R)$ a function of radius $R$ from the central object. Accretion takes place because of the action of some form of dissipation which releases the free energy of the shear flow as heat, and so allows the disc material to fall deeper into the potential well of the central object. Simple physical energy arguments ({\\it e.g.} Lynden-Bell \\& Pringle, 1974) indicate that the dissipative process must take the form of a stress which transports angular momentum outwards. Because the free energy of a circular shear flow is zero only if $d\\Omega/dR = 0$, it follows that the relevant element of the stress tensor must be of the form \\begin{equation} T_{\\phi R} \\propto - d\\Omega/dR, \\end{equation} where $\\phi$ is the azimuthal coordinate.  This can be deduced from the standard derivation of Navier-Stokes stress to be found in the fluid dynamics textbooks. However, in a recent paper, Hayashi \\& Matsuda (2001) have drawn attention to the fact that attempts to provide a physical explanation of the above result in terms of simple kinetic theory have resulted in failure.  The simple kinetic explanation given in basic physics text books for the effect of viscosity on a simple linear shear flow, of the form ${\\bf u} = (0,U(x))$ in Cartesian coordinates, relies on the fact that the kinetic particles conserve linear momentum between collisions. Thus particles crossing some fiducial plane $x = x_o$ tend to mix up and smooth out the momentum distribution of the fluid and so give rise to a stress of the form $T_{yx} \\propto -dU/dx$.  However, as Hayashi \\& Matsuda (2001) point out, it is in the application of these simple concepts to a circular shear flow that the problems seem to arise. The simple generalisation that, in a circular shear flow, the kinetic particles conserve angular momentum $j = R^2 \\Omega$ between collisions would imply, taken at face value, that the the movement of particles across some fiducial circle $R=R_o$ would tend to try to mix up and smooth out the distribution of angular momentum of the fluid, and thus that the stress would be proportional to $-dj/dR$ (see e.g. Madej \\& Paczynski 1977). From the arguments given above, this is clearly wrong. Not only would this predict a stress in the case when the shear $d\\Omega/dR$ is zero, but for a standard Keplerian accretion disc for which $j \\propto R^{1/2}$ it would transport angular momentum inwards rather than outwards. As detailed by Hayashi \\& Matsuda (2001) attempts to get round this and to produce the 'correct' result have only succeeded by making mathematical errors in the derivation.  From all these problems, Hayashi \\& Matsuda (2001) conclude that although what they call the derivation 'with mathematical rigour' ({\\it i.e.} the usual Navier-Stokes argument) gives the correct answer, in order to obtain the correct answer using kinetic theory one must take account of such complications as Coriolis force.  In this paper, we shall show that, although it is obviously necessary to include Coriolis force if one works in a frame co-rotating with the flow, one can obtain the correct result from straightforward kinetic theory in the inertial frame.  Before we do so, it is instructive to return to the 'mathematical' relationship between stress and strain derived in the standard fluid textbooks (see, for example, Batchelor, 1967, Section 3.3). The basic point we wish to make is that the standard Navier-Stokes expression for momentum transfer ({\\it i.e.} stress) is based on a simple {\\it physical} argument. The argument may involve the use of tensors, which physicists tend to meet in courses on mathematical methods, but the argument itself is not a mathematical one. The point is that stress (momentum transfer) in a fluid (or in a solid) can be expressed as a second order tensor. This is a physical (and not a mathematical) statement, in exactly the same sense that the statement that a velocity is a first order tensor (i.e. a vector) is a physical, and not a mathematical, statement.\\footnote{The mathematical complications come in when one has to specify a particular tensor (or vector) by specifying the coordinate representation of that tensor in a particular coordinate frame, and get worse when one has to specify the representation of that same tensor in some other coordinate frame.} For a fluid, the physical {\\it Ansatz} is simply that the stress tensor must be physically related to the rate of strain tensor (which is a second order tensor which derives from the first derivatives of the velocity field, and thus incorporates information about the shear). The simplest assumption is that relationship between these two tensors is through a (fourth order) tensor which is isotropic. It is this assumption of the isotropy of the relationship, which is based on the physical assumption that the fluid itself is isotropic, which gives rise to the standard Navier-Stokes expression for the viscosity.\\footnote{This simplest assumption is not necessarily the correct one. For example it is possible that the viscosity might depend on the absolute orientation of the shear in some inertial frame ({\\it e.g.}  the stress due to convection in a rotating medium -- see Kumar, Narayan \\& Loeb 1995, and references therein; the stress due to non-isotropic mixing -- Bretherton \\& Turner, 1968; or the stress induced by the magneto-rotational instability -- Torkelsson {\\it et. al.} 2000).}  It is important to realise that these tensors ({\\it i.e.} scalars, vectors, second order tensors) exist as physical quantities, independent of any coordinate system. The statement that there is a relationship between two of them is a physical statement. It is only when one has to calculate a particular element of the stress tensor, for example corresponding to linear momentum transfer in a linear shear flow, or to the angular momentum transfer in a circular shear flow, that one has to evaluate coordinate dependent expressions, which can get mathematically complicated. But when one does this, one finds that the flux of linear momentum in a linear shear flow just depends on $-dU/dx$, and that the angular momentum flux in a circular shear flow just depends on $-d\\Omega/dR$. However, this result also enables us to draw another conclusion. The reason for the difference between the terms $dU/dx$ and $d\\Omega/dR$ is due solely to the difference between the coordinate systems. That is, it comes from geometry alone. This implies that when looking for differences in derivations for simple kinetic theory formulae between the linear and circular shear flows, we need only concern ourselves with geometry, and not with dynamical complications such as Coriolis force. In addition, we also need to take note of the fact that the Navier-Stokes expression does depend critically on assumptions about isotropy of the fluid. Thus we should expect to have to make a similar assumption about the properties of our kinetic particles. ",
        "conclusions": "We have shown that it is possible to obtain the correct expression for the momentum transfer in a both a linear and a circular shear flow using simple kinetic theory.  This correct expression implies that an isotropic viscosity transfers momentum {\\it down} an angular velocity gradient. We have noted that it is essential to take proper account of both the geometry and the fluid's isotropy.  \\begin{figure} \\vspace{2pt} \\epsfig{file=fig3.eps,width=8.5cm,height=7.cm} \\caption{Velocity components for two particles, $1$ and $2$, in a linear shear flow originating at emission points located at $ \\pm \\alpha$. Key as in Figure 1.} \\end{figure}  Our analysis allows us to understand the flaw in the simple heuristic argument applied to a circular shear flow (described in Section 1), which leads to the conclusion that since angular momentum is conserved along particle trajectories, it is the gradient in angular momentum which particle mixing tends to smooth out. Particles arrive at the reference patch, at point S, with momentum along the streamline ($y-$momentum) that derives from two sources -- (i) the $y-$component of the random emission velocity (assumed isotropic in the local rest frame of the emission point) and (ii) the $y-$component of the streaming velocity of the emission point relative to the reference patch.  \\begin{figure} \\vspace{2pt} \\epsfig{file=fig4.eps,width=8.5cm,height=7.cm} \\caption{Velocity components for two particles, $1$ and $2$, in a circular shear flow originating at emission points located at $ \\pm \\alpha$. Key as in Figure 1.In this example, $R \\Omega^\\prime + \\Omega =0 $ (i.e. the speed of the flow is everywhere constant) but nevertheless the total velocity vector (open arrowhead) is larger for particle $1$, implying that there is a net flux of $y$ momentum at S.} \\end{figure}  The usual argument given to explain the momentum transfer in a {\\it linear} shear flow is simply that (ii) has the same sign for all the emission points on one side of the reference patch and so that momentum is added to, or subtracted from, the reference patch just depending on the sign of the  velocity gradient in the mean flow. In fact, there is another contribution, equal in magnitude and sign to that described above, which may be understood by considering the particles that are incident at the reference patch from $\\pm \\alpha$. From Figure 3 it may be seen that the $x-$velocity (and hence flux) of particles on the $+ \\alpha$ side exceeds that on the $- \\alpha$ side ({\\it i.e.} there is a greater flux of particles at the reference patch whose emission is prograde than retrograde). Since the $y-$momentum from (i) far exceeds that from (ii), this slight asymmetry in the fluxes of prograde and retrograde particles produces a net momentum flux that is equal to the more obvious source of momentum flux described above.  We may now apply the same considerations to the {\\it circular} shear flow, where once again there are the two contributions to the $y-$momentum -- (i) and (ii) described above. The important difference now is that the mean streaming velocity is no longer in the $y-$direction. The geometry of the circular arc ensures that this contributes a positive (negative) $x-$velocity for $\\alpha > 0 \\: (\\alpha < 0)$, respectively. As may be seen from Figure 4, particles that are incident from $+ \\alpha$ (prograde particles) have a larger amplitude of both $v_x$ and $v_y$ than those from $- \\alpha$. This ensures that the mean $y-$velocity of particles arriving at the reference patch from a particular emission streamline is not equal to the mean $y-$velocity of particles on that emission streamline. This disparity is because of the relative boost in the arrival flux of particles whose random velocity is prograde in the frame of the emission streamline, compared with those that are retrograde.    This now allows us to understand the behaviour of a Keplerian accretion disc at a qualitative level. The mean angular momentum of particles orbiting at radii less than that of the reference patch is less than that of the reference patch. If the distribution of angular momenta of the particles arriving at the reference patch merely reflected the distribution of angular momenta of particles on their parent emission streamlines, this would imply that the arrival of particles at the reference patch from smaller radii should exert a spindown torque. Instead, the relative boost in the {\\it arrival rate} of particles that are emitted in the prograde direction, ensures that the average angular momentum of the particles arriving at the reference patch exceeds the average at the parent streamline. In the Keplerian case, this relative boost in the arrival rate of the prograde particles is enough to reverse the sign of angular momentum transfer. Thus, particles arriving at the reference patch from smaller radii exert a spin up torque, as required."
    },
    "0403/astro-ph0403631_arXiv.txt": {
        "abstract": "We have used an antenna temperature thresholding algorithm on the Bell Laboratories $\\thco$ Milky Way Survey to create a catalog of 1,400 molecular clouds. Of these, 281 clouds were selected for having well-determined kinematic distances.  The scaleheight, luminosity, internal velocity dispersion, and size of the cloud sample are analyzed to show that clouds smaller than $\\sim  10^{5.5} \\, \\msol$ have a scaleheight which is about 35 pc, roughly independent of cloud mass, while larger clouds, the Giant Molecular Clouds, have a reduced scaleheight which declines with increasing cloud mass. ",
        "introduction": "Since star formation occurs in molecular clouds, the formation of molecular clouds is an essential first step to the star formation process. The overall evolution of galactic metallicity implies that hot ejecta from old stars must somehow return to the cold molecular phase in order to form new stars, and that the timescale for this process is short compared to the age of the Galaxy \\citep[e.g.][]{freeman02}. Giant Molecular Clouds (GMCs) are the largest concentrations of molecular material, amounting to $2 \\times 10^5 \\msol$ or more in a region about 50 pc in size \\citep{stark78}. GMCs are strongly concentrated to spiral arms \\citep{stark79c,lee01}, suggesting that they are transient phenomena: hundreds of thousands of solar masses of molecular material collect during the passage of a spiral density wave, and then dissipate some thirty million years later. Thousands of stars are formed during the lifetime of the GMC, including the generators of the giant \\ion{H}{2} regions that trace spiral structure. How molecular clouds form, and why they form, has been the subject of ongoing theoretical investigation \\citep{elmegreen00,pringle01}. Computer simulations of cloud formation have become increasingly complex and realistic \\citep{hartmann01,zhang02}.  The cloud formation process gathers and concentrates interstellar matter.  However this happens, we expect that the random velocity of the resulting cloud would be less than the velocity dispersion of the precursor material, so that the velocity dispersion of GMCs as a class would be less than that of smaller interstellar clouds. This hypothesis is amenable to observational test. The purpose of this {\\em Letter} is to describe such an observational test, and then to quantify the effect in a way which may make for useful comparisons with theory.  In \\S 2, we take the Bell Laboratories $\\thco$ Survey data and use a brightness temperature thresholding algorithm to create a catalog of molecular clouds. We then use our knowledge of Milky Way structure in order to select a subset of clouds with well-determined kinematic distances.  This allows us to determine the physical size of the cloud and its distance above the galactic plane. These data are analyzed in \\S 3 to show that the scaleheight of large clouds is less than that of smaller clouds. ",
        "conclusions": "The data show that small molecular clouds have a scaleheight which is approximately independent of cloud size. Larger clouds, the GMCs, have a scaleheight which falls off with mass.  This has been demonstrated with a relatively small sample of clouds, but it is a clean sample, chosen from a large-scale survey by an algorithmic method. The implication is that the Giant Molecular Clouds, those clouds which are concentrated in the spiral arms of the Galaxy, are also concentrated to the galactic plane. We can understand this as a manifestation of the molecular cloud formation process. Atomic gas clouds, with a scaleheight of $\\sim 100 \\, \\mathrm{pc}$ and a velocity dispersion $\\sim 12 \\, \\kms$, condense out of the diffuse atomic gas. Their cores become molecular and increasingly more condensed, resulting in a population of high-latitude ($\\sim 70 \\, \\mathrm{pc}$), low mass ($M \\ls 100 \\, \\msol$) partially-molecular clouds with a dispersion $\\sim 8 \\, \\kms$ \\citep{malhotra94b,dame94}. These clouds become larger, more centrally condensed, and more bound, with only a slight reduction in scaleheight and velocity dispersion, to become molecular clouds ($ 100 \\, \\msol \\ls M \\ls 10^5 \\, \\msol$), with dispersion $\\sim 7 \\, \\kms$ \\citep{stark89b}.  These clouds can form stars, but are uniformly distributed throughout the galactic disk, and may survive for many galactic rotations. The largest clouds, the GMCs, are rapidly assembled by the passage of a spiral arm.  This process is dissipative, different from the slow addition of material that forms smaller molecular clouds. It results in a significant loss of random velocity per unit mass, and the resulting GMCs are found at the galactic midplane, in the spiral arm.  The ensuing formation of massive stars destroys the cloud, fragmenting it into stars, ionized gas, and small clouds before the next interarm passage."
    },
    "0403/astro-ph0403361_arXiv.txt": {
        "abstract": "Evolution of first population of massive metal-free binary stars is followed. Due to the low metallicity, the stars are allowed to form with large initial masses and to evolve without significant mass loss. Evolution at zero metallicity, therefore, may lead to the formation of massive remnants. In particular, black holes of intermediate-mass ($\\sim 100-500$\\ M$_\\odot$) are expected to have formed in early Universe, in contrast to the much lower mass stellar black holes ($\\sim 10$\\ M$_\\odot$) being formed at present. Following a natural assumption, that some of these  Population III stars have formed in binaries, the physical properties of first stellar binary black holes are presented. We find that a significant fraction of such binary black holes coalesces within the Hubble time. We point out that burst of gravitational waves from the final coalescences and the following ringdown of these binary black hole mergers can be observed in the interferometric detectors. We estimate that advanced LIGO detection rate of such mergers is at least several events per year with high signal to noise ratio ($\\gtrsim 10$). ",
        "introduction": "The properties of Population III stars have stirred a lot of interest in the recent years. It has been realized that zero metallicity stars with masses up to several hundred of solar masses are stable \\citep{2001ApJ...550..890B}. \\citet{2002ApJ...567..532H} estimated black hole (BH) masses formed in the evolution of high-mass metal-free stars. For initial masses above $40$\\ M$_\\odot$ the remnant is a BH with the mass essentially the same as the progenitor with the exception of stars with initial masses between $140$ and $260$\\ M$_\\odot$ which undergo a pair instability supernova (SN) explosions and leave no remnant at all. \\citet{2003ApJ...591..288H} found that these conclusions hold also for low, non-zero metallicity stars. Numerical studies of collapse and fragmentation of a metal free gas in the early Universe \\citep{1999ApJ...527L...5B,2002ApJ...564...23B,2003ApJ...589..677O}   indicate that the initial mass function (IMF) of Population III stars is top heavy, and might be bimodal with the high mass peak around  $100$\\ M$_\\odot$ \\citep{1998MNRAS.301..569L,2001ApJ...548...19N,2003PASP..115..763C}. Every known stellar system contains a considerable fraction of binaries. Therefore it seems quite natural to allow for a possibility that some fraction of massive Population III stars formed in binaries as well, and then investigate their subsequent evolution. There is a possibility that these systems form massive black hole black hole (BH-BH) binaries which in turn may be observable in gravitational waves by the interferometric detectors.  In this paper we consider the properties of Population III binaries with the initial component masses in the range $100-500$\\ M$_\\odot$. In \\S\\,2 we describe the model of evolution of metal-free systems leading to formation of massive BH-BH binaries. In \\S\\,3 we present characteristic properties of the BH-BH population and discuss the observability of their mergers. Conclusions and summary of results are given in \\S\\,4. ",
        "conclusions": "The principal result of the calculations is presented in Figure~\\ref{rsn}. The expected coalescence rate of Population III intermediate-mass BH-BH binaries is high. We predict that advanced LIGO should observe above a thousand strong ($S/N \\gtrsim 10$) events per year. What are the uncertainties of this prediction?  In our calculation we have assumed that Population III stars were formed at redshifts $z$ from $30$ to $10$ out of $f_{\\rm mass} = 10^{-3}$ of the baryon mass contained in the Universe and that the binary fraction of the initial population was $f_b=0.1$. The exact length and duration of the Population III star formation era does not affect the rate calculation. However, the rate scales linearly with $f_{\\rm mass}$ and $f_b$, and although we have chosen rather conservative values for these two quantities, the rate may decrease if the adopted values were significantly lowered.  The IMF of Population III stars is another unknown. Numerical investigations show that the IMF leans towards massive stars. We assumed rather steep IMF (with the slope $\\alpha=-2$), to assure that we do not overproduce highest mass, and therefore the easiest to detect, BH-BH binaries. As it turns out, the change of IMF slope does not significantly alter the detection rate. Farther steepening of the IMF ($\\alpha=-3$) decreases the expected rate by half, while flatter IMF assumption ($\\alpha=-1$) increases the expected rate by a factor of two. Note that we maintain a flat IMF between $1\\,{\\rm M}_\\odot$ and $100\\,{\\rm M}_\\odot$ when calculating $f_{sim}$ -- the fraction of stars that we simulate out of the total population. The observed rate scales linearly with $f_{sim}$. We also fixed the maximum mass of the Population III stars to $500\\,{\\rm M}_\\odot$. Had we allowed for possibility of star formation with the higher mass the rate would increase.  The estimate of the lifetime of the black hole binaries can be strongly affected if the binaries are hardened by interactions in dense stellar environments. Our simulations show that a large number of black hole binaries should be formed with the lifetimes in excess of the Hubble time. Interactions in dense systems may significantly shorten their lifetimes. We demonstrated above that hardening does not affect strongly the expected rate. Another possibility arises that the interactions disrupt some of BH-BH binaries. This would deplete the population of wide systems. None of the two effects, unless operating on extremely short timescales, can affect much systems with shortest coalescence timescales. Therefore, a combination of the two effects may constrain the mergers of Population III BH-BH binaries to large redshifts.  In order to estimate the $S/N$ for advanced LIGO we have used the formulae of \\citet{1998PhRvD..57.4535F}. The predicted values of $S/N$ may still change when more realistic noise curves and binary gravitational wave form templates are known. Because of large masses of the binaries considered in this work the changes in the low-frequency range are most important. The typical ringdown frequency is $\\nu_{qnr} \\approx 90 (300{\\rm M}_\\odot /M) $\\,Hz, so the rate is most sensitive to the detector performance in the low frequency region. The scaling of the rate with the change of $S/N$ normalization can be read off Figure~\\ref{rsn}.  The influence of the evolutionary model assumptions on detection rate was tested. We changed the CE efficiency (increase/decrease by factor of 2); altered evolution through stable RLOF phases from fully conservative to non-conservative cases; changed the specific angular momentum of the matter leaving the system during RLOF (increase/decrease by factor of 2) and varied the critical mass ratio over which the dynamical instability develops (from $q_{\\rm crit}=2$ in standard model to $q_{\\rm crit}=1-3$). The detection rate was decreased at most by a factor of 3 in the above models. Therefore, the predicted rate does not depend strongly on the binary evolution within the model assumptions.  Summarizing, in our calculations we have tried to use rather conservative assumptions in order not to overestimate the detection rate of Population III BH-BH mergers. If several of the assumptions and values of the model parameters are changed within reasonable limits, we  still obtain a significant detection rate. Even if the rate predicted for our already conservative model is decreased by 2-3 orders of magnitude to allow for different aforementioned uncertainties, we still are left with several strong events per year for advanced LIGO detector. For initial LIGO phase the detection rate falls below one event per year.  We have shown that Population III stars may lead to formation of a large number of binaries containing intermediate-mass black holes. A significant fraction of such systems has coalescence times smaller than the Hubble time. Coalescences of intermediate-mass BH-BH binaries should be detectable by advanced interferometric gravitational wave detectors during the ringdown phase, and possibly also merger phase, provided that accurate templates are available. Given the large expected rate of observed coalescences such events could be the primary targets for LIGO burst \\vspace*{-0.35cm} search."
    },
    "0403/astro-ph0403682_arXiv.txt": {
        "abstract": "Using the Far Ultraviolet Spectroscopic Explorer (FUSE) we have obtained 87 spectra of 57 low-redshift ($z<0.15$) active galactic nuclei (AGN). This sample comprises 53 Type 1 AGN and 4 Type 2. All the Type 1 objects show broad \\ovi\\ $\\lambda 1034$ emission; two of the Type 2s show narrow \\ovi\\ emission. In addition to \\ovi, we also identify emission lines due to \\ciii\\ $\\lambda 977$, \\ciii\\ $\\lambda 991$, \\siv\\ $\\lambda\\lambda 1062,1072$, and \\heii\\ $\\lambda 1085$ in many of the Type-1 AGN. Of the Type 1 objects, 30 show intrinsic absorption by the \\ovi\\ $\\lambda\\lambda 1032,1038$ doublet. Most of these intrinsic absorption systems show multiple components with intrinsic widths of 100 \\kms\\ spread over a blue-shifted velocity range of less than 1000 \\kms. Galaxies in our sample with existing X-ray or longer wavelength UV observations also show \\civ\\ absorption and evidence of a soft X-ray warm absorber. In some cases, a UV absorption component has physical properties similar to the X-ray absorbing gas, but in others there is no clear physical correspondence between the UV and X-ray absorbing components. Models in which a thermally driven wind evaporates material from the obscuring torus naturally produce such inhomogeneous flows. ",
        "introduction": "Roughly 50\\% of all Seyfert galaxies show UV absorption lines, most commonly seen in {\\sc C~iv} and Ly$\\alpha$ \\citep{Crenshaw99}. X-ray ``warm absorbers\" are equally common in Seyferts \\citep{Reynolds97, George98}. All instances of X-ray absorption also exhibit UV absorption \\citep{Crenshaw99}. While \\citet{Mathur94} and \\citet{Mathur95} have suggested that the same gas gives rise to both the X-ray and UV absorption, the spectral complexity of the UV and X-ray absorbers indicates that a wide range of physical conditions are present. Multiple kinematic components with differing physical conditions are seen in both the UV \\citep{Crenshaw99, Kriss00b} and in the X-ray \\citep{Reynolds97, Kriss96, Kaspi01}.  The short wavelength response (912--1187 \\AA) of the {\\it Far Ultraviolet Spectroscopic Explorer (FUSE)} \\citep{Moos00} enables us to make high-resolution spectral measurements ($R \\sim 20,000$) of the high-ionization ion {\\sc O~vi} and the high-order Lyman lines of neutral hydrogen. The \\ovi\\ doublet is a crucial link for establishing a connection between the higher ionization absorption edges seen in the X-ray and the lower ionization absorption lines seen in earlier UV observations. The high-order Lyman lines provide a better constraint on the total neutral hydrogen column density than Ly$\\alpha$ alone. Lower ionization species such as {\\sc C~iii} and {\\sc N~iii} also have strong resonance lines in the \\FUSE\\ band, and these often are useful for setting constraints on the ionization level of any detected absorption. The Lyman and Werner bands of molecular hydrogen also fall in the \\FUSE\\ band, and we have searched for intrinsic $\\rm H_2$ absorption that may be associated with the obscuring torus.  We have been conducting a survey of the $\\sim100$ brightest AGN using {\\it FUSE}. As of November 1, 2002, we have observed a total of 87; of these, 57 have $z < 0.15$, so that the {\\sc O~vi} doublet is visible in the FUSE band. ",
        "conclusions": "The multiple kinematic components frequently seen in the UV absorption spectra of AGN clearly show that the absorbing medium is complex, with separate UV and X-ray dominant zones. In some cases, the UV absorption component corresponding to the X-ray warm absorber can be clearly identified (e.g., Mrk~509) \\citep{Kriss00b}. In others, however, {\\it no} UV absorption component shows physical conditions characteristic of those seen in the X-ray absorber (NGC~3516, NGC~5548) \\citep{Kriss96, Brotherton02}. One potential geometry for this complex absorbing structure is high-density, low-column UV-absorbing clouds embedded in a low-density, high-ionization medium that dominates the X-ray absorption. This is possibly a wind driven off the obscuring torus \\citep{KK95, KK01}. At the critical ionization parameter for evaporation, there is a broad range of temperatures that can coexist in equilibrium at nearly constant pressure; for this reason, the flow is expected to be strongly inhomogeneous. What would this look like in reality? As a nearby analogy, consider the HST images of the pillars of gas in the Eagle Nebula, M16. These show the wealth of detailed structure in gas evaporated from a molecular cloud by the UV radiation of nearby newly formed stars \\citep{Hester96}. Figure 1 shows what this might look like in an AGN---a dense molecular torus surrounded by blobs, wisps, and filaments of gas at various densities. It is plausible that the multiple UV absorption lines seen in AGN with warm absorbers are caused by high-density blobs of gas embedded in a hotter, more tenuous, surrounding medium, which is itself responsible for the X-ray absorption.  Higher density blobs would have lower ionization parameters, and their small size would account for the low overall column densities.  \\begin{figure}[ht] \\centerline{\\epsfxsize=3.7in\\epsfbox{gkriss_agn2003_fig1.ps}} \\caption{ An artist's conception of how a molecular torus surrounding an AGN might appear based on HST images of the Eagle Nebula. Note the complex of wisps and blobs of gas close to the surface of the molecular material. \\label{torus}} \\end{figure}  At sight lines close to the surface of the obscuring torus, one might expect to see some absorption due to molecular hydrogen. Given the dominance of Type 1 AGN in our observations so far, the lack of any intrinsic $\\rm H_2$ absorption is not too surprising since our sight lines are probably far above the obscuring torus. NGC~4151 and NGC~3516 are examples where the inclination may be more favorable since these objects have shown optically thick Lyman limits in the past \\citep{Kriss92, Kriss96}, but our FUSE observations do not show such high levels of neutral hydrogen. Molecular hydrogen will not survive long in an environment with a strong UV flux, and this probably accounts for the lack of $\\rm H_2$ absorption.  In summary, we find that \\ovi\\ absorption is common in low-redshift ($z < 0.15$) AGN. 30 of 53 Type 1 AGN with $z < 0.15$ observed using \\FUSE\\ show multiple, blended \\ovi\\ absorption lines with typical widths of $\\sim 100~\\kms$ that are blueshifted over a velocity range of $\\sim$ 1000 \\kms. Those galaxies in our sample with existing X-ray or longer wavelength UV observations also show {\\sc C~iv} absorption and evidence of a soft X-ray warm absorber. In some cases, a UV absorption component has physical properties similar to the X-ray absorbing gas, but in others there is no clear physical correspondence between the UV and X-ray absorbing components."
    },
    "0403/astro-ph0403227_arXiv.txt": {
        "abstract": "With its excellent spatial resolution, low background, and hard-band response, the {\\em Chandra} X-ray Observatory is ideal for performing exploratory surveys. These efficient, sensitive observations can place constraints on fundamental properties of a quasar continuum including the X-ray luminosity, the ratio of X-ray to UV power, and the X-ray spectral shape. To demonstrate the power of such surveys to provide significant insight, we consider two examples, a Large Bright Quasar Survey sample of broad absorption line quasars and a sample of Sloan Digital Sky Survey (SDSS) quasars with extreme C~{\\sc iv} blueshifts. In both cases, exploratory {\\em Chandra} observations provide important information for a physical understanding of UV spectroscopic differences in quasars. ",
        "introduction": "X-ray emission appears to be a universal signature of quasar spectral energy distributions, confirming expectations from accretion physics. Based on rapid variability of soft X-rays in conjunction with the standard black-hole paradigm, these photons are believed to be emitted from the region immediately surrounding the black hole. Energetically, X-rays are significant, contributing 2--20$\\%$ of the bolometric luminosity.  X-ray observations are thus an important component of any multi-wavelength campaign to probe quasar populations.  The excellent spatial resolution of the {\\em Chandra} High Resolution Mirror Assembly and the effective background rejection of the ACIS instrument make this combination uniquely powerful for quasar surveys. For reference, during a 5~ks observation, the 0.5--8.0~keV background within a $2\\arcsec$-radius source region is typically $\\sim$0.1~ct. Because ACIS is photon-limited even beyond $100$~ks (Alexander et al.\\ 2003), the point-source detection limit scales {\\em linearly} with exposure time, unlike the $\\sqrt{t} $ dependence common in other wavelength bands. In conjunction with sub-arcsec positional accuracy, known optical point sources can be robustly detected with 3--5 photons. In 5 ks, this corresponds to a \\mbox{0.5--8.0~keV} flux of $\\sim7\\times10^{-15}$~erg~cm$^{-2}$~s$^{-1}$ for a typical quasar X-ray spectrum.  A 3--7~ks {\\em Chandra} exposure, the regime of exploratory observations, is generally insufficient for gathering enough X-rays for spectral analysis of a quasar.  However, the strategy of exploratory observations enables the extension of results from spectroscopic observations of individual targets to larger, well-defined samples, and the investigation of connections between X-ray properties and other wavelength regimes. From these datasets, standard X-ray observables are \\mbox{0.5--8.0~keV} flux, hardness ratio,\\footnote{The hardness ratio is defined to be $(h-s)/f$, where $h$=2--8~keV ct, $s$=0.5--2.0~keV ct, and $f$=0.5--8.0~keV ct.} and $\\alpha_{\\rm ox}$.\\footnote{The quantity $\\alpha_{\\rm ox}$ equals $0.384\\log(f_{\\rm X}/f_{\\rm 2500})$ where $f_{\\rm X}$ and $f_{\\rm 2500}$ are the flux densities at rest-frame 2~keV and 2500~\\AA, respectively.}  We briefly describe the initial results from two exploratory {\\em Chandra} quasar surveys to illustrate the utility of this observing strategy.  Other examples in the literature of successful applications of this approach to quasar studies include surveys of high-$z$ (e.g., Brandt et al.\\ 2002; Vignali et al.\\ 2003), red (Wilkes et al. 2002), and X-ray weak (Risaliti et al.\\ 2003) quasars.  \\begin{figure} \\plotone{fig1.eps} \\caption{Hardness ratios versus $\\alpha_{\\rm ox}$ from exploratory {\\em Chandra} surveys for two quasar samples. ($a$) The Large Bright Quasar Survey broad absorption line quasar sample (Gallagher et al.\\ 2003). ($b$) The SDSS C~{\\sc iv} blueshift sample (Richards et al., in prep.).} \\end{figure}  \\vspace{-0.1cm} ",
        "conclusions": ""
    },
    "0403/astro-ph0403233_arXiv.txt": {
        "abstract": "We present a newly developed cosmological hydrodynamics code based on weighted essentially non-oscillatory (WENO) schemes for hyperbolic conservation laws. WENO is a higher order accurate finite difference scheme designed for problems with piecewise smooth solutions containing discontinuities, and has been successfully applied for problems involving both shocks and complicated smooth solution structures. We couple hydrodynamics based on the WENO scheme with standard Poisson solver - particle-mesh (PM) algorithm for evolving the self-gravitating system. The third order low storage total variation diminishing (TVD) Runge-Kutta scheme has been used for the time integration of the system. To test accuracy and convergence rate of the code, we subject it to a number of typical tests including the Sod shock tube in multidimensions, the Sedov blast wave and formation of the Zeldovich pancake. These tests validate the WENO hydrodynamics with fast convergence rate and high accuracy.  We also evolve a low density flat cosmological model ($\\Lambda$CDM) to explore the validity of the code in practical simulations. ",
        "introduction": "Though the universe seems to be dominated by the dark sides of both matter and energy (Turner, 2002), the observed luminous universe has been existing in the form of baryonic matter, whose mass density, constrained by the primordial nucleosynthesis (Walker, et al., 1991), only occupies a small amount of the total density. To account for the observational features revealed by the baryonic matter, i.e., X-ray emitting gas in galaxies and clusters (Mulchaey, 2000), intergalactic medium inferred from Ly$\\alpha$ forest (Rauch, 1998), X-ray background radiation (Giacconi et al. 1962) and distorted spectrum of the cosmic background radiation due to the Sunyaev-Zeldovich effect (Zel'dovich \\& Sunyaev 1969; Ostriker \\& Vishniac 1986) etc., it would be necessary to incorporate hydrodynamics into cosmological investigations. This motivation has stimulated great efforts to apply a variety of gas dynamics algorithms to cosmological simulations. For a general review of the state-of-the-art on this topic, we refer to Bertschinger (1998).  Due to the high non-linearity of gravitational clustering in the universe, there are two significant features emerging in cosmological hydrodynamic flow, which pose more challenges than the typical hydrodynamic simulation without self-gravity. One significant feature is the extremely supersonic motion around the density peaks developed by gravitational instability, which leads to strong shock discontinuities within complex smooth structures. Another feature is the appearance of an enormous dynamic range in space and time as well as in the related gas quantities. For instance, the hierarchical structures in the galaxy distribution span a wide range of length scales from a few kpc resolved by individual galaxy to several tens of Mpc characterizing the largest coherent scale in the universe.  A variety of numerical schemes for solving the coupled system of collisional baryonic matter and collisionless dark matter have been developed in the past decades. They fall into two categories, particle methods and grid based methods.  The particle methods include variants of the smooth-particle hydrodynamics (SPH; Gingold \\& Monagham 1977, Lucy 1977) such as those of Evrard (1988), Hernquist \\& Katz (1989), Navarro \\& White (1993), Couchman, Thomas \\& Pierce (Hydra, 1995), Steinmetz (1996), Owen et al. (1998) and Springel, Yoshida \\& White (Gadget, 2001). The SPH method solves the Lagrangian form of the Euler equations, and could achieve good spatial resolutions in high density regions, but works poorly in low density regions. It also suffers from degraded resolution in shocked regions due to the introduction of sizable artificial viscosity.  The grid based methods are to solve the Euler equations on structured or unstructured grids. The early attempt was made by Cen (1992) using a central difference scheme. It uses artificial viscosity to handle shocks and has first-order accuracy. The modern approaches implemented for high resolution shock capturing are usually based on the Godunov algorithm. The two typical examples are the total-variation diminishing (TVD) scheme (Harten 1983) and the piecewise parabolic method (PPM) (Collella \\& Woodward 1984). Both schemes start from the integral form of conservation laws of Euler equations and compute the flux vector based on cell averages (finite volume scheme). The TVD scheme modifies the flux using an approximate solution of the Riemann problem with corrections added to ensure that there are no postshock oscillations. While in the PPM scheme, the Riemann problem is solved accurately using a quadratic interpolation of the cell-average densities that is constrained to minimize postshock oscillations. In the cosmological setting, the TVD based codes include those of Ryu et al. (1993), the moving-mesh scheme (Pen, 1998), and the smooth Lagrangian method (Gnedin, 1995); and the PPM based codes include those of Stone \\& Norman (Zeus; 1992), Bryan et al. (1995), Sornborger et al. (1996), Ricker, Dodelson \\& Lamb (COSMOS; 2000). The grid-based methods suffer from the limited spatial resolution, but they work extremely well both in low and high density regions as well as in shocks. To reach a large dynamical range, the Godunov methods have also been implemented with adaptive mesh refinement (RAMSES: Teyssier, 2002; ENZO: Norman \\& Bryan, 1999; O'Shea et al., 2004), which is more adequate to explore the fine structures in the hydrodynamic simulation.  We describe in this paper an alternative hydrodynamic solver which discretizes the convection terms in the Euler equations by the fifth order finite difference WENO (weighted essentially non-oscillatory) method, first developed in Jiang \\& Shu (1996), with a low storage third order Runge-Kutta time discretization, which was proven to be nonlinearly stable in Gottlieb \\& Shu (1998). The WENO schemes are based on the essentially non-oscillatory (ENO) schemes first developed by Harten et al. (1987) in the form of finite volume scheme for hyperbolic conservative laws. The ENO scheme generalizes the total variation diminishing (TVD) scheme of Harten (1983). The TVD schemes typically degenerate to first-order accuracy at locations with smooth extrema while the ENO scheme maintains high order accuracy there even in multi-dimensions. WENO schemes further improve upon ENO schemes in robustness and accuracy. Both ENO and WENO schemes use the idea of adaptive stencils in the reconstruction procedure based on the local smoothness of the numerical solution to automatically achieve high order accuracy and non-oscillatory property near discontinuities. For WENO schemes, this is achieved by using a convex combination of a few candidate stencils, each being assigned a nonlinear weight which depends on the local smoothness of the numerical solution based on that stencil. WENO schemes can simultaneously provide a high order resolution for the smooth part of the solution, and a sharp, monotone shock or contact discontinuity transition. WENO schemes are extremely robust and stable for solutions containing strong shocks and complex solution structures. Moreover, a significant advantage of WENO is its ability to have high accuracy on coarser meshes and to achieve better resolution on the largest meshes allowed by available computer memory. We will describe the fifth order WENO scheme employed in this paper briefly in \\S 3. For more details, we refer to Jiang \\& Shu (1996) and the lecture notes by Shu (1998, 1999).  WENO schemes have been widely used in applications. Some of the examples include dynamical response of a stellar atmosphere to pressure perturbations (Zanna, Velli \\& Londrillo, 1998); shock vortex interactions and other gas dynamics problems (Grasso \\& Pirozzoli, 2000a; 2000b); incompressible flow problems (Yang et al., 1998); Hamilton-Jacobi equations (Jiang \\& Peng, 2000); magneto-hydrodynamics (Jiang \\& Wu, 1999); underwater blast-wave focusing (Liang \\& Chen, 1999); the composite schemes and shallow water equations (Liska \\& Wendroff, 1998, 1999); real gas computations (Montarnal \\& Shu, 1999), wave propagation using Fey's method of transport (Noelle, 2000); etc.  In the context of cosmological applications, we have developed a hybrid N-body/hydrodynamical code that incorporates a Lagrangian particle-mesh algorithm to evolve the collisionless matter with the fifth order WENO scheme to solve the equations of gas dynamics. This paper is to detail this code and assess its accuracy using some numerical tests. We proceed as follows. In \\S 2, we present the basic cosmological hydrodynamic equation for the baryon-CDM coupling system. \\S 3 gives a brief discussion of the numerical scheme for solving the hydrodynamic equations, especially about the implementation of the finite difference fifth order WENO scheme and the TVD time discretization. In \\S 4, we validate the code using a few challenging numerical tests. Concluding remarks are drawn in \\S 5. ",
        "conclusions": "In this paper, we have described a newly developed hybrid cosmological hydrodynamic code based on weighted essentially non-oscillatory (WENO) schemes for the Euler system of conservation laws. We implement the fifth order finite difference WENO to solve the inviscid fluid dynamics on a uniform Eulerian grid combining with a third order low storage Runge-Kutta TVD scheme for advancing in time. In order to solve the cosmological problem involving both collisional baryonic matter and collisionless dark matter, we incorporate the particle-mesh method for computing the self-gravity into our cosmological code.  The code has been subjected to a number of tests for its accuracy and convergence. As expected, the WENO scheme demonstrates its capacity of capturing shocks and producing sharp and non-oscillatory discontinuity transition without generating oscillations. In comparison with other existing hydrodynamic codes such as the TVD or PPM schemes, one striking feature of the WENO code is that it retains higher order accuracy in smooth regions including at smooth extrema even in multidimensions, and yet it is still highly stable and robust for strong shocks. In performance, the WENO scheme needs more floating point operations per cell than those of the PPM and TVD schemes. However, in compensating for twice or more loss of the computational speed, the WENO scheme achieves both higher order accuracy and convergence rate than PPM and TVD codes according to our numerical tests.  In the presence of gravity, the hydrodynamics become more challenging than that without gravity due to the highly non-linearity of gravitational clustering. One serious problem encountered in many cosmological applications is the so called high Mach number problem. To address this problem, we have incorporated an extra technique into our cosmological WENO/PM code, which is actually a combination of the dual energy algorithm (Bryan et al. 1995) and the energy-entropy algorithm (Ryu et al. 1993). Namely, instead of solving the internal energy equation in regions free of shocks as was done in the dual energy algorithm (Bryan et al. 1995), we solve the modified entropy equation (Ryu et al. 1993), which takes a conservative form and can be easily solved using the standard WENO scheme. This improvement over our hydrodynamic WENO code ensures an accurate tracking of the temperature field in regions free of shocks.  It is pointed out that the high order WENO discretization, e.g., the fifth order WENO scheme adopted in this paper, introduces a quite small numerical viscosity, which does not lead to a significant violation of energy conservation in the presence of gravitational fields. While for a second-order TVD scheme, the numerical diffusion is no longer negligible. In order to have a better conservation of the total energy, it is usually corrected by adding a compensation term in the gravitational force term (Ryu et al, 1993).  The Euler hydrodynamics on fixed meshes have several distinct advantages which includes simplicity for implementation, easy data parallelization, relatively low floating point cost, large dynamic range in mass and high resolution of shock capturing. In particular, the WENO scheme can also achieve a higher accuracy on coarse meshes and a better resolution on the largest meshes allowed by available memory. To suit for large simulations of the cosmological problem, further improvement of the hydrodynamic WENO code is needed in its implementation on distributed memory computers. The parallel version of the code based on the Message-Passing Interface (MPI) has been under development."
    },
    "0403/astro-ph0403005_arXiv.txt": {
        "abstract": "This is the written version of an invited review talk for the 13 Feb 2004 AAAS Meeting in Seattle. The talk's goal is to present a philosophical view of extragalactic astronomy as it applies to the sub-field of galaxy evolution.  The talk is divided into three parts: 1) How we got to where we are (technology drivers to our science goals), 2) What's new and special (how that technology has achieved our recent science results) and 3) How an improved worldview will help us in the near future.  The intended audience for this talk is a generally knowledgeable scientist, but not an astronomer by training.  This talk is also *not* intended to be a complete review of the field of galaxy evolution and only includes a few recent results extracted from the astro-ph archives to present the current state of our field. ",
        "introduction": "Historians will no doubt look back onto the last part of the 20th century as a true \"Golden Age\" for observational cosmology and galaxy evolution due to the enormous advancements driven primarily by technological leaps in the telescopes, detectors and computers.  Like most modern sciences, the study of galaxies is beyond the basic human senses (although the Magellenic Clouds are visible to the naked eye in the southern hemisphere and Andromeda is at the limit of the human eye for very dark sites).  Thus, the study of galaxies did not even begin until the development of 1m class telescopes at the end of the 19th century.  With the discovery of an expanding Universe (Hubble 1936) and the confirmation of a Creation point (Penzias \\& Wilson 1965), it was quickly realized that astronomers have a unique capability, that no other science has. This is the capability to actually follow the evolution of distant objects while the size of the Universe is much greater than the speed of light (i.e. lookback time).  To imagine the impact of lookback time on a field of science, imagine the response of a palaeontologist if offered a device that would allow them to observe animal behavior during the Jurassic era, or a musicologist who is allowed to listen in on Mozart's practice sessions.  Lookback time means that the more distant an object, the farther into our past it is. Thus began in extragalactic astronomy the great hunt for high redshift objects, redshift being a measure of distance in an expanding Universe.  Of course, more distant also means fainter in apparent luminosity, which drives the need for larger and larger telescopes to collect the photons from high redshift, or very distant, galaxies.  Our first quantum jump in the study of galaxy evolution was during the 1950's and 60's with the development of the Palomar 5m and NOAO's 4m telescopes.  With this technology, spectroscopy out to several billion light-years was doable.  Solid state technology in the 80's gave us highly efficient and highly reliable (for photometry) detectors.  The 90's gave us 10m class telescopes and the Hubble Space Telescope, arguably the most important instrument for the study of galaxy evolution.  Parallel to observational achievements were breakthroughs in many theoretical arenas, primarily the coming of age of the computer simulation in the 80's as the main tool for astrophysical research (observers call themselves astronomers, while theorists call themselves astrophysicists for some strange reason).  During the 90's, decreasing access to telescopes (due to an increasing population of observers) and decreasing computer hardware costs produced a glut of theoretical Ph.D. theses.  Within a decade, our field went from the point where theorists numbered one for every four observers to near parity.  When combined with uncertain values for the basic cosmological constants (i.e. $H_o$ and $q_o$), this lead to an explosion of ideas, but with very few constraints to narrow our focus on what was really happening after galaxy formation.  Unfortunately, the same technology that allows us to observe the past also confines our window of knowledge.  For example, galaxies present themselves over a range of luminosities, which corresponds directly with the mass of the galaxy (i.e. the number of stars).  While the brightest galaxies (up to 10$^{13}$ $M_{\\sun}$) are the most dominant members of clusters of galaxies, and the easiest to obtain photometric and spectroscopic information, they are not the dominant members of the Universe by number.  That honor goes to the small, dwarf galaxies (masses around 10$^8$ $M_{\\sun}$) who number over a factor of 10 to 100 more numerous than the brightest galaxies in clusters.  Thus, the limiting magnitude of your telescope (the depth in luminosity that can be achieved for a decent S/N) will allow you to study bright galaxies to distant redshifts (far into the past) or both bright and faint galaxies at lower redshifts. This bias forced our conclusions about galaxy evolution in the 80's and 90's to be limited to only the very brightest galaxies in rich environments.  The newest technology allows us to address the full galaxy population at cosmologically interesting distances of 0.5 to 0.7 (see Figure 1.1).  \\begin{figure} \\centering \\includegraphics[width=12cm]{lum_morph.epsf} \\caption{ This figure displays the effects of limiting magnitude to the ability of galaxy evolution studies to resolve their science goals.  Early technology (shown by the 5m limit) shows that only the brightest galaxies could be measured at cosmologically interesting distances.  The Hubble Deep Field (HDF) extended the depth to more normal galaxies and the Hubble Ultra Deep Field (HUDF) will allow a comparison of the evolution of dwarf and giant galaxies, a key test to the hierarchical merger scenario of galaxy formation. } \\end{figure}  The greatest progress with respect to the study of galaxy evolution was the determination of the cosmic distance scale parameters, $H_o$ and $q_o$ by the Hubble Cepheid Key Project and the High Redshift Supernova Cosmology Project.  While these two parameters are not directly related to galaxy evolution parameters, their determination has finally released researchers from being forced to interpret evolutionary effects through a myriad of combinations of $H_o$ and $q_o$.  Or, worse, attempting to deduce these cosmological parameters from poorly known evolutionary effects.  Galaxy evolutionists are now free to focus solely on the star formation history of galaxies with lookback time determined solely by the redshift of the galaxy. ",
        "conclusions": ""
    },
    "0403/astro-ph0403696_arXiv.txt": {
        "abstract": "We revisit the notion that galaxy motions can efficiently heat intergalactic gas in the central regions of clusters through dynamical friction.  For plausible values of the galaxy mass-to-light ratio, the heating rate is comparable to the cooling rate due to X-ray emission.  Heating occurs only for supersonic galaxy motions, so the mechanism is self-regulating: it becomes efficient only when the gas sound speed is smaller than the galaxy velocity dispersion. We illustrate with the Perseus cluster, assuming a stellar mass-to-light ratio for galaxies in the very central region with the dark-matter contribution becoming comparable to this at some radius $r_s$.  For $r_s \\la 400~{\\rm kpc} \\sim 3 r_{\\rm cool}$---corresponding to an average mass-to-light ratio of $\\sim10$ inside that radius---the dynamical-friction coupling is strong enough to provide the required rate of gas heating.  The measured sound speed is smaller than the galaxy velocity dispersion, as required by this mechanism.  With this smaller gas temperature and the observed distribution of galaxies and gas, the energy reservoir in galactic motions is sufficient to sustain the required heating rate for the lifetime of the cluster.  The galaxies also lose a smaller amount of energy through dynamical friction to the dark matter implying that non--cooling-flow clusters should have flat-cored dark-matter density distributions. ",
        "introduction": "\\label{sec:intro}  Galaxy cluster gas loses thermal energy copiously through X-ray emission. In the absence of energy input, radiative cooling in cores of clusters should result in  substantial gas inflow  (see Fabian 1994 for a review). These ``cooling flows'' would have associated mass deposition rates of several hundred ${\\rm M_\\odot yr^{-1}}$ in some clusters (Peres \\etal 1998). Nevertheless, recent high-resolution X-ray observations (e.g., Peterson \\etal 2001, 2003) have revealed that there is little evidence for the expected multi-phase gaseous structures, strongly suggesting that mass dropout is being prevented by some source that is heating the gas, thereby balancing radiative energy loss in  the central region of clusters. Recent work  has focused on two prospective heating mechanisms: (1) diffusive heat transport, via thermal conduction (\\pcite{tuc83,bre88,nar01,voi02,fab02,zak03,kim03a}) and/or turbulent mixing (\\pcite{cho03,kim03b,voi04}), from hotter gas in the outer region to that in the core; (2) energy input from jets, outflows, and radiation from a central active galactic nucleus (\\pcite{cio01,chu02,bru02,kai03}).   There is however at least one additional mechanism that appears to have been overlooked. It involves the energy lost by  concentrated clumps of matter (galaxies) as they move through the cluster. Part of this energy may go into re-arranging the dark-matter mass distribution  (El-Zant \\etal 2003), but a  significant fraction  should end up deposited in the gas. That dynamical-friction (DF) coupling can  transform the dynamical energy of  galaxies into thermal motion in the gas has been known for at least four decades (e.g., Dokuchaev 1964). Relatively recent work involving this notion includes the analysis by Miller (1986) of the Perseus cluster and that of Just \\etal (1990) concerning the Coma cluster.  Both studies confirm that, provided that the mass-to-light ratio of galaxies is $\\sim 20$, energy lost by galaxies to the gas should be sufficient in counteracting the radiative cooling of the gas in the central region of these clusters.    Several developments on the empirical side suggest renewed relevancy for this mechanism. One obvious one involves  recent X-ray data confirming that a heating mechanism {\\em is} required; whereas the consensus in the 1980's was against this conclusion, it now seems inescapable. The second involves revisions to the inferred gas electron densities in the central region of clusters; best values referred to by Miller and Just \\etal are of the order of $10^{-3} {\\rm cm^{-3}}$, while current best estimates are rather in the range of $10^{-2}-10^{-1} {\\rm cm^{-3}}$ (Kaastra \\etal 2004).  This leads to  an  order-of-magnitude increase in the dynamical friction coupling between the galaxies and gas. There has also been progress in determining the mass-to-light ratio of galaxies. On the theoretical side, work by Just \\etal (1990) and Ostriker (1999) has since shed some light on the behavior of the dynamical-friction coupling in a gaseous medium in the transonic and subsonic regimes.   There appears to be {\\em a priori} no reason why the rate of energy loss from galaxies to gas via dynamical friction should be of the same order of that radiated by the gas.  However, {\\em this coupling is active only when the sound speed of the gas is smaller than the typical velocity of galaxies}.  The mechanism is therefore self-regulating; the gas cools until the dynamical-friction heating rate is always equal to the cooling rate.  We start by pointing out, in the next Section, why this is expected to be so, before moving on to develop a Monte-Carlo model to estimate the  total energy expected to flow into the cooling region of the Perseus cluster, using recent data for galaxy luminosities and projected positions, as well as for the gas parameters of that cluster, and averaging over a set of different realizations where three-dimensional positions, velocities, and mass-to-light ratios are treated as stochastic quantities.  The final Section discusses briefly the central dark-matter distribution and energetics, as well as some remaining questions. ",
        "conclusions": "\\label{sec:disc}   That dynamical-friction coupling between cluster galaxies and gas can significantly heat the latter component is not a new concept. Indeed Miller (1986) has pointed out that a few luminous galaxies  in the Perseus cluster's central  region may alone  deposit sufficient energy to keep gas in that region from cooling, provided that they had a mass-to-light ratio of about 20. We revisited this cluster using recently compiled galaxy and gas data and  determinations of  the mass-to-light ratios in galaxies. Our Monte Carlo model  shows that the rate at which energy is deposited by galaxies into the gas within the cooling radius can completely compensate for radiative loss from within that radius, if galaxies at radii  $\\la 400$ kpc from the center of the cluster have dark mass that is comparable to their stellar mass (with galaxies at smaller radii having progressively smaller dark-matter content, vanishing for those near the center of the cluster). The   associated  average mass-to-light ratio within this radius is about 10.    A robust and potentially important prediction of the model is that the drop in gas temperature invariably  observed in the cooling regions of clusters  is necessary for the DF coupling to be sufficiently strong. There is a natural interpretation for this phenomenon: the gas cools until dynamical-friction heating becomes sufficiently efficient to prevent a further drop in temperature.  The gas random motion is then coupled to that of the galaxies,  and so remains in approximate equilibrium with it. All this  simply follows from the fact that dynamical friction drops sharply for subsonic motion. In the case of the Perseus cluster, the transition from weak to strong coupling  occurs precisely in the range of sound speeds found in the cooling core of the cluster.   The action of the mechanism discussed in this paper also has consequences for the dark-matter distribution. As shown by El-Zant \\etal (2003), dynamical friction from the galaxies will heat a   density cusp, creating a core  with radius corresponding to roughly a fifth of the original Navarro, Frenk \\& White (1997) scale length (cf., Fig.~1 of El-Zant et al.), where dark matter has been spread out to larger radii (Fig.~3 of El-Zant et al.). Approximating the initial density distribution in this region such that  $\\rho_i = \\rho_0 (r_0/r)$ and the final one with $\\rho_f = \\rho_0$, the energy required for this transformation is $\\Delta \\Phi = (8 \\pi G)^{-1}  \\int_{0}^{r_0} (|d \\phi/dr|_i^2 - |d\\phi/dr|_f^2) 4 \\pi r^2 dr = 22 \\pi^2 \\rho_0^2 r_0^5/45$ [with $d\\phi/dr = (4 \\pi G/r^2 )\\int \\rho r^2 dr$]. Taking $r_0 = r_{\\rm cool}$ and $\\rho_0$ to correspond to four times the value defined by the electron density $\\rho_e (r_0) = 0.0053$ (where we have used the gas to  gravitational mass ratio in Table~5 of Peres \\etal 1998 and equation (4) of Churazov \\etal 2003 for the electron density), one finds $\\Delta \\Phi = - 6 \\times 10^{60}~{\\rm erg}$, which is comparable to the binding energy of the CD galaxy $\\Phi_{\\rm CD} \\approx - (G M_{\\rm CD}^2/R_{\\rm CD})$. It is however significantly smaller than the energy radiated from gas inside the cooling radius for the age of the cluster (e.g., for five Gyr this amounts to  $8 \\times 10^{61}\\, {\\rm erg}$). Nevertheless, this energy is  easily available from fast moving galaxies from beyond the very inner region---for example its value coincides with the  kinetic energy in a mass similar to that of the CD galaxy and moving at $2000~\\kms$, and that material can be supplied solely by the stellar mass of galaxies within $2 r_{\\rm cool}$---the region from which, according to our model, the bulk of energy input to the gas comes from.   It is worth noting here that, under the circumstances just described, the mass of the gas within $r_{\\rm cool}$ [which using equation (4) of Churazov \\etal 2003, adjusted for $h=0.7$ with $r_{\\rm cool} = 130\\,{\\rm kpc}$ amounts to $2.2 \\times 10^{12} {\\rm M_\\odot}$] is comparable to that of the galaxies in that region. Due to the temperature drop in the gas, thermal motions can be significantly smaller than that of the galaxies. The energy of the gas can therefore be smaller than that of the galaxies by up to an order of magnitude. There is therefore sufficient energy in galaxies to heat the gas many times over,  with the crucial parameter, which was the focus of this paper, being the rate at which this is transferred. Furthermore, once the dark matter in the center has been heated, it absorbs little additional energy from the galaxies, since  dynamical friction does not act on fast particles, and the coupling with the dark matter at larger radii decreases rapidly with radius (see El-Zant \\etal 2003 for further discussion).  The energy lost to the dark matter is therefore less than that going into  keeping the gas at constant temperature. The mass of the CD galaxy, as well as the spatial and velocity distribution of galaxies in the central region of the cluster will reflect the history of energy loss to both components.    Several questions remain open. Prominent among these is the issue of thermal stability. From the condition that thermal stability requires a remarkably narrow range of heating rates, Bregman \\& David (1989) have argued against Miller's proposition that the Perseus cluster gas is heated by DF from galaxies.  However the functional form of the heating rate used by these authors does not agree with later calculations that have been borne out by numerical studies---e.g., their postulated form has a {\\em negative} energy transfer rate ($\\sim -1/V^3$) for highly subsonic velocities, resulting in gas {\\em cooling}, which is incompatible with the positive (even if small)  heating found by Just \\etal (1990) and Ostriker (1999) in that limit. It will be therefore necessary to repeat these calculations using adequate forms for the heat-loss function. These calculations would also address another crucial question, that concerning  the precise manner in which  the energy lost by the galaxies is distributed in the gas. We have adopted the simple approach where this energy is deposited isotropically and equally in logarithmic intervals. This should approximate how the energy is initially deposited, at least for highly supersonic galaxies (and since, in the cooling region,  the sound speed may be several times smaller than the  velocity dispersion, this may not be a very bad assumption). Even then, however, the deposited  energy  may still  be transported toward the center of the cluster by wave motions, as suggested, for example,  by Balbus \\& Soker (1990), and redistributed.  We note here that the claim of these authors that DF from galaxies is insufficient to heat the cluster core is based upon outdated values for the electron density and an outdated ($1/r$) gas distribution.  With updated values for the electron density and gas distribution (e.g., the empirical formula of Churazov \\etal 2003), their conclusions are changed, as we have shown.  We thus conclude that it is far from obvious that DF-heating of cluster gas is irrelevant."
    },
    "0403/astro-ph0403375_arXiv.txt": {
        "abstract": "{ We present the \\XMM{} spectra of three low-redshift intermediate Seyferts (one Sy1.5, and two Sy1.8), from our survey of hard spectrum \\Rosat{} sources. The three AGN are well fitted by absorbed powerlaws, with intrinsic nuclear photoelectric absorption from column densities between 1.3 and 4.0 $\\times 10^{21}$~cm$^{-2}$.  In the brightest object the X--ray spectrum is good enough to show that the absorber is not significantly ionized.  For all three objects the powerlaw slopes appear to be somewhat flatter ($\\Gamma\\sim1.3-1.6$) than those found in typical unabsorbed Seyferts.  The constraints from optical and X--ray emission lines imply that all three objects are Compton-thin. For the two fainter objects, the reddening deduced from the optical broad emission lines in one of them, and the optical continuum in the other, are similar to those expected from the X--ray absorption, if we assume a Galactic gas-to-dust ratio and reddening curve.  The broad line region Balmer decrement of our brightest object is larger than expected from its X--ray absorption, which can be explained either by an intrinsic Balmer decrement with standard gas-to-dust ratio, or by a $>$Galactic gas-to-dust ratio. These $\\ge$~Galactic ratios of extinction to photoelectric absorption cannot extend to the high redshift, high luminosity, broad line AGN in our sample, because they have column densities $>10^{22}$~cm$^{-2}$, and so their broad line regions would be totally obscured. This means that some effect (e.g., luminosity dependence, or evolution) needs to be present in order to explain the whole population of absorbed AGN. ",
        "introduction": "According to the unified model for Active Galactic Nuclei (AGN)(Antonucci 1993), broad-line Seyferts (type 1) and narrow-line (type 2) Seyferts are intrinsically the same type of object but are viewed with different orientations to our line of sight. In this model the central engine of the AGN, and the high velocity clouds that produce the broad optical and UV emission lines, are surrounded by a thick torus of dust and cool, molecular gas. In type 1 objects we have a direct view of the central engine and the broadline clouds, whereas in type 2 objects the torus blocks our line of sight to these regions. Such a torus has a large photoelectric opacity at soft X--ray energies, which explains why type 1 Seyferts have steep X--ray spectra with little absorption (Nandra \\& Pounds 1994), and type 2 Seyferts have absorbed X--ray spectra (Smith \\& Done 1996). Narrow emission lines are produced in more distant gas clouds on scales larger than the dusty torus, and so can be seen in both types of object.  Objects in which the broad lines are attenuated, but not completely obscured, are called intermediate Seyferts, and are assigned classifications ranging from 1.5-1.9.  However, even in the type 1 objects, an extra absorption component from ionized gas is often seen in the X--ray band (George et al.  1998).  This ionized gas, the ``warm absorber'', appears to be distributed throughout the BLR and NLR of Seyfert galaxies, and in type 2 objects it can scatter nuclear radiation into our line of sight. The warm absorber has a much lower opacity to soft X--rays than cold gas, but if it contains a significant element of dust (e.g. Brandt, Fabian \\& Pounds 1996) it could produce considerable extinction at optical and ultraviolet wavelengths.  An understanding of absorption in AGN is extremely important.  For example, unified AGN models for the X--ray background (XRB) (see e.g., Setti \\& Woltjer 1989, for an early proposal, and Gilli, Salvati \\& Hasinger 2001, for a late development) explain the spectrum of the XRB by the superposition of spectra of AGN with various degrees of absorption. Under this scheme, soft spectrum X--ray sources should be mostly type 1 AGN, while hard spectrum X--ray sources would be predominantly type 2 AGN. In the 1990s this matched the observations quite well, because \\Rosat{} surveys in the soft band were dominated by type 1 AGN, (e.g. Mason et al. 2000, Lehman et al. 2001) while surveys selected in harder bands with {\\sl BeppoSAX} (Fiore et al. 1999) were much richer in type 2 AGN.  Despite these early successes, recent developments have put in jeopardy the identity between optical type 1 and X--ray unabsorbed objects, on one hand, and optical type 2 and X--ray absorbed objects, on the other. For example, identifications of our survey of \\Rosat{} sources with hard spectra (Page, Mittaz \\& Carrera 2000, 2001) produced mostly type 1 AGN, contrary to the expectations of the unified model. Other examples of X--ray absorbed type 1 objects have been found by Akiyama et al. (2000) using {\\sl ASCA} data, and Mainieri et al. (2002) and Page et al. (2003) using \\XMM{} data. In principle, high gas-to-dust ratios in the X--ray absorbing gas (perhaps due to dust sublimation close to the central X--ray source, Granato, Danese \\& Franceschini 1997), or large dust grains (Maiolino, Marconi \\& Oliva 2001), could give rise to high levels of X--ray absorption, without much optical obscuration.  In the opposite sense, Pappa et al. (2001) have found several examples of type 2 AGN with very little or no X--ray absorption. Panessa \\& Bassani (2002) estimate that 10-30\\% of Seyfert 2 galaxies have this property.  One striking example is H1320+551 (Barcons, Carrera \\& Ceballos 2003), a Seyfert 1.8/1.9 galaxy with strong optical (BLR and NLR) obscuration, but without any corresponding X--ray absorption from cold gas. The high quality of their \\XMM{} data allows these authors to rule out a warm absorber in this source, which leads them to the conclusion that the BLR is intrinsically reddened in this object: its Sy 1.8/1.9 appearance cannot arise from obscuration of a Seyfert 1 spectrum.  The situation is therefore complex.  Possible explanations include Compton thick obscuration which could suppress completely the nuclear emission below 10~keV. This spectral range could then be filled by X--rays scattered off the warm absorber, or by extranuclear emission, which would not have in principle a particularly hard or absorbed spectrum. This could result in optically obscured type 2 AGN which appear to be absorption free at X--ray energies. Such a model can in principle be tested, since Bassani et al. (1999) have developed a diagnostic diagram that permits identifying Compton thick sources, as those with high equivalent width Fe emission lines (originating in fluorescence in the torus material), and low 2-10~keV to [OIII] flux ratio. This is based on the observation that [OIII] originates in the NLR, outside the torus, and thus in principle [OIII] should be free of obscuration. Neither H1320+551, nor any of the sources discussed in Pappa et al. (2001), lie in the Compton thick region of this diagram. They represent therefore genuine mismatches between optical and X--ray classifications, at odds with the unified AGN model.   Here, we analyze optical and \\XMM{} spectroscopic data on three AGN (RXJ133152.51\\-+111643.5, RXJ163054.25\\-+781105.1, and RXJ213807.61\\--423614.3) from the sample of Page, Mittaz \\& Carrera (2001). Of the objects in this sample which show broad optical emission lines, these three had the highest X--ray fluxes.  All three show strong signs of absorption in their \\Rosat{} spectra. In section \\ref{OSp} we present their optical spectra, finding evidence for optical obscuration in at least two of them. We then analyze their \\XMM{} spectra in section \\ref{XSp}, and in particular we measure their intrinsic X--ray absorption. The differences between the levels of optical obscuration and X--ray absorption are discussed in section \\ref{discussion}, as well as a comparison of the X--ray spectral properties of our sources with respect to those of other samples at similar flux levels. Finally, in section \\ref{discussion} we summarize our results.  For brevity, we will refer to the three sources using truncated versions of their names (i.e. RXJ1331, RXJ1630 and RXJ2138) in the text.  We have used the currently fashionable values of $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{m}=0.3$, and $\\Omega_{\\Lambda}=0.7$, throughout this paper. ",
        "conclusions": "\\label{conclusions}  We have presented X--ray spectra from \\XMM{} of the three brightest AGN exhibiting broad optical emission lines (all three are intermediate Seyferts) from the sample of hard spectrum \\Rosat{} sources (Page, Mittaz \\& Carrera 2000, 2001). The X--ray spectra of all three sources are well fitted by powerlaws ($\\Gamma\\sim1.5$), absorbed by moderate amounts of intrinsic nuclear cold material ($N_{\\rm H}\\sim$a few $\\times10^{21}$ cm$^{-2}$). Similarly absorbed sources are an important part of the $S_{\\rm 2-10\\,keV}\\geq 10^{-13}$~cgs AGN population, although the three sources studied here appear to have harder intrinsic power law slopes than the majority of AGN at this flux level.  The equivalent width of the narrow emission line at about 6.4~keV found in RXJ1331 is typical of other radio-quiet Compton-thin Seyferts.  Detailed analysis of our optical spectroscopic data confirm the classification of these sources as intermediate-type Seyfert galaxies (Sy1.5-Sy1.9). For the two objects in which both H$\\alpha$ and H$\\beta$ are visible (RXJ1331 and RXJ2138), the NLR is shown to be almost free of reddening, while the BLR is significantly reddened. The column density necessary to produce this effect is about 1.5 times that inferred from the X--ray absorption in RXJ1331, and of the order of it in RXJ2138 and RXJ1630 (in this last case from its broad band optical spectrum), if standard Galactic gas-to-dust ratios are assumed.  None of the three sources are Compton-thick.  The X--ray data for RXJ1331 require that the absorber is cold, allowing us to rule out the presence of dust embedded in a warm absorber in this source.  These three low redshift broad line AGN from our sample of hard sources show a ratio of optical extinction to X--ray absorption which is similar to, or larger than, the interstellar medium of our own Galaxy. This cannot be the case for the high luminosity, high redshift, broad line AGN in our sample, because the broad lines would be completely obscured at the X--ray column densities ($>10^{22}$~cm$^{-2}$) observed in these sources.  To explain the whole population of absorbed AGN, their effective gas-to-dust ratio must show a large variety, perhaps depending on luminosity, evolving with redshift, or showing geometries different from those proposed by the unified AGN model."
    },
    "0403/astro-ph0403469_arXiv.txt": {
        "abstract": "Stars form by gravoturbulent fragmentation of interstellar gas clouds. The supersonic turbulence ubiquitously observed in Galactic  molecular gas generates strong density fluctuations with gravity taking over in the densest and most massive regions. Collapse sets in to build up stars and star clusters.  Turbulence plays a dual role. On global scales it provides support, while at the same time it can promote local collapse.  Stellar birth is thus intimately linked to the dynamical behavior of parental gas cloud, which determines when and where protostellar cores form, and how they contract and grow in mass via accretion from the surrounding cloud material to build up stars.  Slow, inefficient, isolated star formation is a hallmark of turbulent support, whereas fast, efficient, clustered star formation occurs in its absence.  The fact that Galactic molecular clouds are highly filamentary can be explained by a combination of compressional flows and shear.  The dynamical evolution of nascent star clusters is very complex. This strongly influences the stellar mass spectrum. The equation of state (EOS) plays a pivotal role in the fragmentation process. Under typical cloud conditions, massive stars form as part of dense clusters.  However, for gas with effective polytropic index greater than unity star formation becomes biased towards isolated massive stars, which may be of relevance for understanding Pop III stars. ",
        "introduction": "\\label{sec:intro} Star clusters form by gravoturbulent fragmentation in interstellar clouds. The supersonic turbulence ubiquitously observed in Galactic gas clouds generates strong density fluctuations with gravity taking over in the densest and most massive regions.  Once such cloud regions become gravitationally unstable, collapse sets in and leads to the formation of stars and star clusters. Yet the conditions for fragmentation and the physical processes that govern the early evolution of nascent star clusters are poorly understood.  Following up on analytical studies (starting with Jeans 1902; and later  Larson 1969; Shu 1977; Elmegreen 1993; Padoan 1995; Padoan \\& Nordlund 2002), most current investigations concentrate on a numerical approach to star cluster formation.  For example, the effects of interstellar turbulence have been studied extensively in a series of 3D simulations by Klessen, Burkert, \\& Bate (1998), Klessen, Heitsch, \\& Mac Low (2000), Klessen \\& Burkert (2000, 2001), Heitsch, Mac Low, \\& Klessen (2001a,b), Klessen (2001). See also Ballesteros-Paredes et al.\\ (1999ab, 2003), Padoan \\& Nordlund (1999), Padoan et al.\\ (2001), Bate, Bonnell, \\& Bromm (2003) or Bonnell, Bate, \\& Vine (2003).  A complete overview is given in the reviews by Larson (2003) and Mac Low \\& Klessen (2004).   In this proceedings paper we call your attention to the dynamical complexity arising from the interplay between supersonic turbulence and self-gravity, and introduce the concept of gravoturbulent fragmentation.  We argue that in typical star forming clouds turbulence generates the density structure in the first place and then gravity takes over in the densest and most massive regions to build up the star cluster.  In Section 2 we focus on spatial distribution and timescale of star formation, then in Section 3, we discuss a specific example of a star forming filament similar to those observed in Taurus, and in Section 4 we speculate about the mass spectra of clumps and stars in the context of the gravoturbulent fragmentation model. Finally, in Section 5 we demonstrate that the equation of state (EOS) of the interstellar gas plays a pivotal role in gravoturbulent fragmentation. The EOS determines whether molecular cloud regions build up clusters of low to intermediate-mass stars, or form isolated high-mass objects. ",
        "conclusions": ""
    },
    "0403/astro-ph0403143_arXiv.txt": {
        "abstract": "{ We present 2\\farcs4 resolution, high sensitivity radio continuum observations of the nearby spiral galaxy  NGC\\,4258 in total intensity and linear polarization obtained with the Very Large Array at $\\lambda$3.6~cm (8.44~GHz). \\\\ The radio emission along the northern jet and the center of the galaxy is polarized and allows investigation of the magnetic field. Assuming energy-equipartition between the magnetic field and the relativistic particles and distinguishing between (1) a relativistic electron-proton jet and (2) a relativistic electron-positron jet, we obtain average magnetic field strengths of about (1) 310~$\\mu$G and (2) 90~$\\mu$G. The rotation measure is determined to range from $-$400 to $-$800~rad/m$^{2}$ in the northern jet. Correcting the observed E-vectors of polarized intensity for Faraday rotation, the magnetic field along the jet turns out to be orientated mainly along the jet axis. An observed tilt with respect to the jet axis may indicate also a toroidal magnetic field component or a slightly helical magnetic field around the northern jet. ",
        "introduction": "The nearby galaxy NGC\\,4258 (M\\,106) is a bright SAB(s)bc spiral (de Vaucouleurs et al.\\ \\cite{vauc}) at a distance of 7.2~Mpc (Herrnstein et al.\\ \\cite{herr99}). It seems to possess a small bright nucleus with a highly excited emission line spectrum (Burbidge et al.\\ \\cite{bur}) and has been classified as a weakly active Seyfert~2-type galaxy.  Most striking are the two so-called `anomalous arms', which are not visible in the optical and were first detected in H$\\alpha$ by Court\\`{e}s \\& Cruvellier (\\cite{cour}) in the inner region of the galaxy. Van der Kruit et al. (\\cite{kruit}) detected these anomalous arms in the radio range where they extend out to the optical periphery. Spectral index studies indicate that their radio emission is of non-thermal origin (de Bruyn\\ \\cite{bruyn}; van Albada \\cite{alb}).  The anomalous arms of NGC\\,4258 have been extensively discussed in terms of ejection of matter from the nucleus. The detection of a water-maser (Claussen et al.\\ \\cite{claus}; Henkel et al.\\ \\cite{henk}) and an accretion disk around a supermassive central object (Miyoshi et al.\\ \\cite{miyo}; Herrnstein et al.\\ \\cite{herr}), and the fact that inner anomalous arms are orientated parallel to the rotation axis of the accretion disk and can be traced even on subparsec scale (Herrnstein et al.\\ \\cite{her}) indicate that the anomalous arms indeed are jets. Whereas in the inner regions the anomalous arms clearly reveal their jet character, many of the features at a greater distance from the nucleus (e.g. their bifurcation) remain unexplained.  The three-dimensional geometry of the galaxy and its jets has been discussed for a long time (cf. e.g. van Albada \\& van der Hulst\\ \\cite{albhulst}; Hummel et al.\\ \\cite{hum}). The detection of an accretion disk revealed directly for the first time that the central part of the galaxy has a significant tilt with respect to the galactic disk: the accretion disk itself has an inclination angle of $83\\degr$ and a position angle (p.a.) of $86\\degr$ (Miyoshi et al.\\ \\cite{miyo}), i.e. it is oriented nearly east-west. The p.a. of the galactic disk, however, is $150\\degr$, which is nearly north-south, and its inclination is $72\\degr$ (van Albada \\cite{albada}). Thus, the plane of the galactic disk and the plane of the accretion disk have a significant angle to each other. As the jets emerge perpendicular to the accretion disk, they have to pass the galactic disk at a rather small angle. Due to the projection of the whole system with respect to the Earth, the inner jet direction is almost parallel to the major axis of the galactic disk.  To further investigate the jet geometry and especially the magnetic field along the jets we obtained data with the Very Large Array (VLA)\\footnote{The VLA is a facility of the National Radio Astronomy Observatory. The NRAO is operated by Associated Universities, Inc., under contract with the National Science Foundation.} at $\\lambda$3.6~cm and compared these with already published but reprocessed VLA data by Hummel et al. (\\cite{hum}) at $\\lambda$6.2~cm (4.86~GHz) and $\\lambda$20~cm (1.49~GHz) and with observations at $\\lambda$2.8~cm (10.55~GHz) made with the Effelsberg 100-m telescope.  The observations and data reduction procedures are described in Sect.~2. In Sect.~3 we present the results and examine the measurements of the linearly polarized emission in terms of Faraday rotation, magnetic field strength and direction. The discussion of the results and the summary follow in Sect.~4 and Sect.~5, respectively. ",
        "conclusions": "We present interferometer data obtained with the VLA in its C-configuration at $\\lambda$3.6~cm (8.4399~GHz) in total power and linear polarization. For comparison and to obtain quantities like rotation measure and depolarization we also reprocessed VLA data at $\\lambda$6.2~cm (4.8851~GHz) and $\\lambda$20~cm (1.4899~GHz) that were previously published by Hummel et al. (\\cite{hum}) and observations at $\\lambda$2.8~cm (10.55~GHz) made with the Effelsberg 100-m telescope.  The high resolution maps ($2\\farcs4$ HPBW) at $\\lambda$3.6~cm as obtained with the C-configuration of the VLA are able for the first time to resolve the jets in the central region of the galaxy. Detailed inspection shows that they emerge from the galactic center along the projected spin axis of the accretion disk as determined by Miyoshi et al. (\\cite{miyo}). At a distance $24\\arcsec$ away from the center they change direction towards the previously seen northwest to southeast and bifurcate at $62\\arcsec$ from the nucleus in the northern jet and at $85\\arcsec$ in the southern jet. The multiple splitting of the northern jet is clearly visible at this resolution.  The polarized emission was detected exclusively along the jets and allowed the calculation of the magnetic field strength in the inner region of the northern jet. Energy equipartition considerations lead to magnetic field strengths of $310\\pm 15~\\mu$G assuming a relativistic electron-proton jet and $90\\pm 5~\\mu$G assuming an electron-positron jet. The rotation measure could be determined between $\\lambda$3.6~cm and $\\lambda$6.2~cm, at a linear resolution of $14\\arcsec$ HPBW. It is derived to be $-$400 to $-$800~rad/m$^2$ in the northern jet. Correcting the observed E-vectors of polarized emission for Faraday rotation, the magnetic field is mainly along the jet axis in the central region and tends to become somewhat tilted with respect to the jet direction in the outer part of the northern jet. This may be consistent with a slightly helical magnetic field around the northern jet or may indicate a superposition of a longitudinal magnetic field near the jet axis and a toroidal magnetic field away from the axis."
    },
    "0403/astro-ph0403625_arXiv.txt": {
        "abstract": "{\\it RXTE} ASM count rates from the X-ray pulsar Her X-1 began falling consistently during the late months of 2003.  The source is undergoing another state transition  similar   to  the  anomalous low state    of 1999. This new event  has triggered observations  from both space- and ground-based observatories.  In order  to aid data interpretation  and telescope scheduling, and to facilitate the phase-connection of cycles before  and  after the  state   transition, we have  re-calculated the precession ephemeris using cycles over  the last 3.5 years.  We report that the source has displayed a different  precession period since the last anomalous event.  Additional archival data from \\cgro\\  suggests that each  low  state is accompanied  by a  change in precession period and that   the subsequent period is correlated  with accretion flux.  Consequently our analysis reveals long-term accretion disk  behaviour  which    is predicted   by     theoretical models  of radiation-driven warping. ",
        "introduction": "\\label{sec:introduction}  Her X-1 is an eclipsing X-ray pulsar that displays variability on spin (1.24s),   orbital  (1.7d)   and    super-orbital   (35d)   timescales \\citep{tan72}.  The companion star is   of early-F or late-A type  and fills its  Roche lobe, resulting  in  accretion mainly  by Roche  lobe overflow through the inner Lagrangian ($L_1$) point \\citep{lea98}.  The 35d  period  is  revealed in  X-rays   as an absorbing   column of variable depth passing in front of the accretion source \\citep{sco00}, In the UV and optical, an intrinsic flux variation  is observed on the same  period \\citep{ger76,lea99}.   The  long-standing  model is  of a warped accretion  disk,  viewed close to  edge-on ($i  \\simeq 90$ deg) precessing  in   a  retrograde  direction  around  the   neutron star, \\citep{kat73,rob74}.  The disk  atmosphere   provides the  column  and casts an  X-ray  shadow over the  inner  face  of the  companion star, causing optical variability on 35d timescales.  \\begin{figure*} \\begin{picture}(100,0)(10,20) \\put(0,0){\\special{psfile=\"f1_color.eps\" hoffset=20 voffset=-550 angle=0 vscale=80 hscale=80}} \\noindent \\end{picture} \\vspace{190mm} \\figcaption{Top panel: \\her\\ \\cgro\\ BATSE (20--50 \\kev) one-day average count  rates from Earth occultations  between 1991--1999. Middle panel: \\xte\\  ASM (1--12  \\kev) one-day average  count rates  from \\her\\ from 1996--2004.   Bottom panel: ASM Hardness  ratio  averaged over each 35 day cycle.  Band  A = 1.3--3 \\kev,  band B =   3--5\\kev\\ and band C  = 5--12\\kev.  Horizontal  lines represent count rate  thresholds adopted in the $0-C$ analysis. Dwells  over the two  thresholds are plotted on the $O-C$ diagram of Fig. \\ref{fig:o-c}\\label{fig:asm}} \\end{figure*}  In X-rays,  the 35d cycle   is  characterized by  a  rapid turn-on,  a ``main-on''  state  which varies  in intensity greatly   from cycle to cycle, and then  decays back to an ``off''  state at $\\sim 1$\\% of the peak main-on flux. This is followed by a rise to a weaker ``short-on'' state and then another low  \\citep{tan72,gor82}. Superimposed on  this cycle are eclipses of the neutron star by the companion once per orbit and dips  that  occur  almost on the  beat   period between orbit  and precession \\citep{gia73}, which  are believed to  be due to  localized structure in the  disk, generated by the impact  of the  ballistic gas stream  from the $L_1$  point \\citep{cro80,sch96}.  The turn-on occurs at only two orbital phases, $\\phio  \\simeq 0.2$ or $\\phio \\simeq 0.7$, apparently  at random.   Consequently  it has  often been claimed that precession is not strictly   periodic (Ogelman 1987),   and precession cycles generally  have durations  of either 20,   20.5  or 21  orbital cycles,  leading  some authors  to  suggest  that  the precession  and orbital cycles are related physically \\citep{sco99}.  To date, it has not been possible  to determine the engine behind disk precession   observationally.  From   a  theoretical  perspective, the companion star can force a tilted disk to precess in the outer regions closest to  the donor star \\citep{lar98}. However  tidal forces do not account for the  mechanism that tilts  the disk.   Mechanisms that can simultaneously    warp   structure  and   drive   precession  are wind \\citep{sch94} and radiation \\citep{pri96,wij99} pressure, pushing down on  the  disk  plane.  The dominant  engine  behind  this pressure  is accretion   flux from  the neutron  star    which irradiates the  disk atmosphere  externally.  While  the  simplifying  assumptions used  in these theoretical  calculations   prevent a detailed  comparison  with observation, the success of the models  in producing warped disks with the observed precession rate indicates that some combination of tides, wind and radiation pressure are capable of driving the \\her\\ clock.  Her  X-1   has occasionally  displayed  consecutive main-on  states of significantly  lower flux  than   average,  e.g.   1993 \\citep{vrt94}. Also,   between  1999--2000,  the source  missed   a  large number  of consecutive cycles  altogether  \\citep{par99,cob00,vrt01,sti01}.   See also \\citet{par85}  for the description  of a 1983  event prior to the epoch of all-sky monitoring surveys.  During  these times, the optical flux  does   not dip appreciably  \\citep{vrt01}.  In   addition to the decrease in main-on amplitude,  the neutron star experiences an abrupt transition from the nominal state  of pulse period spin-up to episodes of spin-down \\citep{par99}.  Although these events are not homogeneous in  terms of  either duration   or  magnitude, they have  been  coined collectively the `Anomalous Low  States (ALS)'.  Quite clearly the ALS are significant disk events; they point to the accretion disk changing state  for a limited period of  time before returning seemingly to the same preferred period and warp shape, albeit with an offset of several orbital cycles in the   turn-on clock, according to  \\citet{oos01} and \\citet{man03}. In this paper, we show this picture to be incorrect.  {\\it Rossi X-ray Timing Explorer  (RXTE)} All-Sky Monitor (ASM)  count rates from  Her X-1  in the  main-on state  began falling consistently during  the late  months of  2003.   \\her\\  appears  to be  undergoing another state  transition similar to the ALS  of 1999--2000.   Since a glitch   in  the precession clock  was   reported  after the  last ALS \\citep{oos01}, it is now    prudent  to recalculate  the    precession ephemeris using the last 3.5 years of  data.  Rather than a glitch, we report  that the source  has displayed a  different precession period, \\pprec, since  the  last ALS.   Additionally, archival   data from the BATSE  experiment that  was   onboard  the  {\\it  Compton  Gamma   Ray Observatory (CGRO)} suggests that each ALS  is accompanied by a change in \\pprec.  Contrary to  previous reports \\citep{oos01}, these results indicate  that the accretion  disk does not  return to  the same state after each ALS. ",
        "conclusions": "\\label{sec:conclusions}  We  have  redefined the accretion  disk  precession ephemeris  for the X-ray pulsar \\her\\ between the end of the 1999--2000 ALS and the start of the 2003--2004 ALS event.   Relative to the most recently published work  on  cycle  timing  \\citep{man03}, the  new ephemeris  provides a correction to the  turn-on time of the  precession cycle of $\\sim 14$d at the time of  writing, critical  to  research teams  and observatory schedulers during the  current  ALS. A period  analysis  of the \\cgro\\ BATSE  and \\xte\\  ASM time series   since 1991 reveals  four epochs of stable precession period. The  precession epochs are each separated by an ALS and the change of period appears to be instantaneous within the resolution  of the  $O-C$  diagram,  i.e.   a  few precession  cycles. Testable predictions are that  \\her\\  will return from  the 2003--2004 ALS with a new precession period and epoch-averaged peak main-on flux."
    },
    "0403/astro-ph0403413_arXiv.txt": {
        "abstract": "We present a systematic spectral analysis of six ultraluminous X-ray sources (NGC1313 X-1/X-2, IC342 X-1, HoIX X-1, NGC5408  X-1 and NGC3628  X-1) observed with {\\sl XMM-Newton} Observatory. These extra-nuclear X-ray sources in nearby late-type galaxies have been considered as intermediate-mass black hole candidates. We have performed Monte-Carlo simulations of Comptonized multi-color black-body accretion disks. This unified and self-consistent spectral model assumes a spherically symmetric, thermal corona around each disk and accounts for the radiation transfer in the Comptonization. We find that the model provides satisfactory fits to the {\\sl XMM-Newton} spectra of the sources. The characteristic temperatures of the accretion disks ($T_{in}$), for example, are in the range of $\\sim 0.05-0.3$ keV, consistent with the intermediate-mass black hole interpretation. We find that the black hole mass is typically about a few times $10^3~\\rm{M_\\odot}$ and has an accretion rate $\\sim 10^{-6} - 10^{-5} ~\\rm{M_\\odot~yr^{-1}}$. For the spectra considered here, we find that the commonly used multi-color black-body accretion disk model with an additive power law component, though not physical, provides a good mathematical approximation to the Monte-Carlo simulated model. However, the latter model provides additional constraints on the properties of the accretion systems, such as the disk inclination angles and corona optical depths. ",
        "introduction": "Probably the most exciting recent development in the field of black hole (BH) study is the discovery of numerous candidates for intermediate-mass BHs (IMBHs) with masses in the range of $\\sim 10^2-10^5 M_\\odot$ (see Miller \\& Colbert 2003 for a recent review).  The presence of such BHs was first proposed to explain ultraluminous X-ray sources (ULXs), defined as point-like extra-nuclear X-ray sources observed in nearby galaxies and with  inferred isotropic X-ray luminosities in excess of $10^{39} {\\rm~erg~s^{-1}}$, about an order of magnitude greater than the Eddington limit of a solar mass object (e.g., Fabbiano 1989; Colbert \\& Mushotzky 1999; Miller et al. 2003a,b; Strohmayer \\& Mushotzky 2003). While unambiguous detections of individual IMBHs currently do not exist, there are observational hints from studies of microlensing events, globular clusters, and centers of nearby galaxies (van der Marel 2003 and references therein). Although ULXs actually represent a heterogeneous population, a majority of them are likely to be accreting BHs. The controversy is centered on the X-ray emission mechanisms and on the masses of the BHs (e.g., Makishima et al. 2000; King et al. 2001; Begelman 2002).  The IMBH interpretation, though probably the most straightforward and exciting, has serious difficulties (e.g., King et al. 2001; Kubota et al. 2002). In addition to their high X-ray luminosities, many ULXs show convex-shaped spectra, especially in the energy band $\\lesssim $ a few keV. Such spectra are characteristic of the ``soft state'' of accreting BH binaries and are often approximated by black-body-like models such as the multi-color disk model (MCD; {\\em diskbb} in the {\\em XSPEC} spectral analysis software package (Arnaud 1996; e.g., Makishima et al. 1986). In the MCD model, each annulus of the axis-symmetric optically-thick accretion is assumed to radiate as a blackbody with a radius-dependent temperature. The characteristic temperature  $T_{in}$ of the innermost portion of the disk is $\\propto (\\dot{M}/M)^{1/4}$, where $M$ and $\\dot{M}$ are the BH mass and the accretion mass rate. However, $T_{in}$ inferred from the model fit is almost always too high for the required high mass, assuming the Eddington limit on $\\dot{M}$. Equivalently, the inner disk radius $R_{in}$ is much smaller than the last stable orbit for a non-spin BH. Even more disturbing is that the inferred value of $R_{in}$ is sometimes found to be time-variable, in contrast to the soft-state of confirmed BH X-ray binaries, where $R_{in}$ is approximately constant for each source.  Furthermore, it has been shown recently that many X-ray spectra of ULXs cannot be satisfactorily described by a single MCD model, especially in observations with good counting statistics and with a broad energy coverage (e.g., Miller et al. 2003a,b). The usual practice is then to fit such a spectrum with an additive combination of an MCD and a power law (PL; e.g, Miller 2003a,b; Cropper \\etal 2003). The requirement of this latter component, which often becomes important at energies $\\gtrsim$ a few keV and may extend up to 200 keV as indicated by stellar mass BH systems, suggests that a high temperature electron cloud (corona) exists around an accretion disk, producing inverse-Compton scattering of the disk photons (e.g., Kubota et al. 2002; Page et al. 2003). A spectral fit with this additive, phenomenological model combination typically leads to an acceptable fit and a much lower (apparently more reasonable) disk temperature. Nevertheless, the model is over-simplified in the following two aspects:  First, the extension of the power-law component straight to the low-energy limit of the spectrum is nonphysical. Because the power-law component is assumed to mimic the effect of the inverse Compton scattering, a low-energy cutoff as in the MCD component must be present in the Comptonized component (see also Page et al. 2003). Neglecting this low-energy cutoff could mis-characterize the spectral shape of the MCD component and could lead to an artificially high absorption in spectral fitting.  Second, the additive combination of the MCD and the PL components does not account for the radiative transfer process or the removal of photons from the MCD component to the Comptonized component. Furthermore, this process depends on both disk photon and corona electron energies. Therefore, one may not, in general, directly take the MCD normalization derived from the spectral model fitting to infer the inner disk radius or the BH mass, as realized by some authors (e.g., Kubota, Makishima \\& Ebisawa 2001). Otherwise, the inferred (but probably nonphysical) inner disk radius, for example, may appear to vary significantly when an accretion system changes from its hard state to its soft state or vice versa, as in XTE J2012+381 (Campana et al. 2002).  These oversimplifications in the MCD+PL model certainly obscure the physical dependence of the Comptonization on the corona and disk properties, and could also seriously affect the inference of accretion disk parameters.  Yao \\etal (2003, hereafter Paper I) have recently developed a Comptonized multi-color disk (CMCD) model. They use Monte-Carlo simulations  to directly generate the Comptonized X-ray spectra, removing the above mentioned over-simplifications and avoiding the complications of using the two~(unrelated)-component model and then trying to correct for various radiative transfer effects. While the MCD model is still used to describe the accretion disk emission, the Comptonized radiation is no longer an independent component. This self-consistent treatment thus provides a new tool to constrain the physical properties of the corona and to study its relationship to the accretion disk. But most importantly, the  model enables us to recover the same original disk flux in a spectral fit, which is essential to a reliable mass estimate of the putative BH.  In the present work, we apply the CMCD model to the analysis of {\\sl XMM-Newton}  spectra of six ULXs which have been suggested as IMBH candidates. Our main objectives are to check whether or not the CMCD model provides an adequate spectral description of these sources and to see what potential new insights we may gain from such an application. The sources and data are described in \\S 2, whereas the implementation of the model for this application is discussed in \\S 3, which also includes a summary of various corrections required to infer BH masses from the present model fits.  We also test the PL, MCD, and MCD+PL models and compare them with the CMCD model. We present the results of our spectral fits in \\S 4 and discuss the implications and conclusions in \\S 5. ",
        "conclusions": "The satisfactory fits of the CMCD model to the {\\sl XMM-Newton} spectra of our selected ULXs suggest that they are consistent with the IMBH interpretation. In particular, the model does not have the high $T_{in}$ problem as is faced by the MCD model. The problem is apparently caused by the neglect of Comptonization in the model. Although this neglect is statistically allowed when both the counting statistics and the energy band coverage of an observed spectrum are poor, the fitted spectral parameters are far from being physical. We conclude that the MCD model alone should not be used in the interpretation of ULXs as IMBHs.  We confirm that both $T_{in}$ and $N_H$ inferred from the MCD+PL model are reasonably accurate for the sources considered here. This apparent agreement between the CMCD and MCD+PL models suggests that the latter model as a whole is mathematically a good representation of the former model, at least for the IMBH candidates considered here. This is rather surprising when one considers the over-simplifications in the MCD+PL model, as discussed in \\S 1. It appears that the nonphysical extrapolation of the PL to the low energy nearly compensates the failure to include the radiation transfer loss of soft disk photons. However, this does not mean that the MCD+PL model could be used to describe the Comptonized disk emission in general. The various nonphysical effects can cause problems for other sources, especially those with higher $T_{in}$ values ($\\sim$ 1 keV; see the discussion in \\S 3; e.g., LMC~X-1 and LMC~X-3; Paper II).  In comparison, the CMCD model provides more reliable measurements of the disk parameters as well as unique constraints on the physical properties of the coronae. In the following, we briefly discuss both the function of these new parameters and the physical reason for their different degrees of constraints:  \\begin{itemize} \\item The opacity $\\tau$ is relatively well constrained, which is the key parameter that determines the total number of Comptonized photons (e.g., Fig.~\\ref{fig:compare}). For example, IC 342 X-1 with the largest best-fit $\\tau$ value appears to have the disk emission nearly completely Comptonized, explaining why the spectrum of the source can be characterized by a PL alone. Also for this source, because the saturated Comptonization dominates over the thermal emission, the constraints on $T_{in}$, $R_{in}$, and eventually on BH mass are very weak.  \\item The corona electron temperature $T_c$ is chiefly responsible for the overall energy extent of the Comptonized spectral component. Because $T_c$ is $\\gtrsim 30$ keV for all the sources, the high-energy turning-off of the component is well beyond the \\xmm\\ band limit. Therefore, the data do not constrain the upper limit of $T_c$.  The lower limit is determined because a minimum electron energy is needed to up-scatter soft disk photons to the high energies covered by the spectra.  \\item Whereas the nearly isotropic Comptonized flux is barely affected by the disk inclination angle $\\theta$, the observed strength of the soft disk component is proportional to cos($\\theta$). In a  spectral fit, however, this difference in the disk inclination dependence may be partially compensated by a change in the $\\tau$ value. But if $\\theta$ is large (for a nearly edge-on disk),  its geometric effect cannot be canceled by adjusting other parameters, which would also effectively alter the spectral shapes of both the disk and Comptonized components in a spectral fit (e.g., Fig.~\\ref{fig:compare}). Consequently, we may constrain the upper limit to $\\theta$. This constraint, though not very tight, is important for the estimation of the BH masses (\\S 3).  \\item $R_c$ determines the effective corona radius, within which the disk emission is most affected by the Comptonization. Photons from larger radii have relatively little chance to be scattered and may contribute to the un-Comptonized disk component even if $\\tau$ is large ($\\geq 1$). But the amount of soft X-ray radiation from the disk also decreases with the increasing radius. The combination of these two effects may thus place a constraint on $R_c$. \\end{itemize}  Apparently, these parameters are correlated in a spectral fit. This, together with the limited counting statistics and bandwidth of the data, explains why the parameters are not tightly constrained. Nevertheless, the results presented above demonstrate the potential of the CMCD model to shed new insights into the physical properties of the accretion disk coronae, in addition to a more reliable mass estimate of  the putative BHs.  Table~\\ref{results:flux} includes our estimated BH masses, assuming no spin and the best-fit $\\theta$ values. The typical BH mass is in the range of $\\sim 10^3-10^4$ M$_\\odot$, although the upper limit for IC~342~X-1 is slightly higher. If a BH spins rapidly, the inferred BH mass could be several times higher than the value quoted in the table (depending on $\\theta$; Fig. 2).  Assuming that the bolometric luminosity (estimated in the 0.05--100 keV range) $L_{bol} = 0.1  \\dot{M}c^2$, we further estimate the accretion rate $\\dot{M}$ for each source (Table 5), which is in the range of 1--10 $\\times 10^{-6}$ M${\\rm_{\\odot}yr^{-1}}$.  The present work represents, at most, an incremental step in developing a fully self-consistent model for accreting BH systems. The CMCD model used here deals only with the Comptonization by static disk coronae. To study the dynamics, one needs to understand the formation and evolution of the coronae as well as the physics of the accretion disks. We also have not considered other proposed scenarios that may explain some of the ULXs; e.g., the anisotropic emission of the radiation \\citep{king01}, the relativistic motion of the X-ray-emitting plasmas \\citep{fab01}, and the possible super-Eddington emission (Begelman 2002). Spectral models for such scenarios, yet to be developed, need to account for the apparent presence of the soft thermal component, in addition to the power law, for these sources except in the case of IC 342~X-1. The bottom line here is that the \\xmm\\ spectra are consistent with the IMBH interpretation of the sources."
    },
    "0403/astro-ph0403139_arXiv.txt": {
        "abstract": "We present a deep JHK$_s$-band imaging survey of the W3 Main star forming region, using the near-infrared camera, SIRIUS (Simultaneous three-color InfraRed Imager for Unbiased Surveys), mounted on the University of Hawaii 2.2 m telescope. The near-infrared survey covers an area of $\\sim$ 24 arcmin$^2$ with 10 $\\sigma$ limiting magnitudes of $\\sim$ 19.0, 18.1, and 17.3 in  J, H, and K$_s$-band, respectively. We construct JHK color-color and J/J-H and K/H-K color-magnitude diagrams to identify young stellar objects and estimate their masses. Based on these color-color and color-magnitude diagrams, a rich population of YSOs is identified which is associated with the W3 Main region. A large number of previously unreported red sources (H-K $>$ 2) have also been detected around W3 Main. We argue that these red stars are most probably pre-main sequence stars with intrinsic color excesses. We find that the slope of the K$_s$-band luminosity function of W3 Main is lower than the typical values reported for the young embedded clusters. The derived slope of the KLF is the same as that found by Megeath et al. (1996), from which analysis by Megeath et al. indicates that the W3 Main region has an age in the range of 0.3--1 Myr. Based on the comparison between models of pre-main sequence stars with the observed color-magnitude diagram we find that the stellar population in W3 Main is primarily composed of low mass pre-main sequence stars. We also report the detection of isolated young stars with large infrared excesses which are most probably in their earliest evolutionary phases. ",
        "introduction": "The W3 giant molecular cloud (GMC) complex is located in the Perseus spiral arm at a distance of 1.83$\\pm$0.14 kpc from the Sun (Imai et al. 2000). W3 GMC hosts two massive and active star forming regions, W3 Main in the north, and W3 (OH) in the south. The W3 Main star forming region contains objects such as H II regions, embedded infrared sources (including the extremely luminous cluster of sources W3 IRS 5), OH and water masers (Wynn-Williams et al. 1974, Forster et al. 1977, Gaume \\& Mutel 1987), which are in different stages of evolution.  The millimeter continuum observations have shown the existence of two dense clumps of about 2000 M$_{\\odot}$ associated with the luminous infrared sources IRS 4 and IRS 5 in the W3 core (Richardson et al. 1989). VLA observations have resolved the IRS 5 region into a cluster of seven distinct centimeter radio sources (Tieftrunk et al. 1997, hereafter TGC97; Claussen et al. 1994, hereafter CG94). TGC97 proposed that the spatial and kinematic relation of the compact, ultracompact, and hypercompact radio continuum regions toward W3 Main is indicative of sequentially triggered star formation caused by the pressure of the expanding H II regions and the subsequent compression of the molecular gas. Recently, high resolution continuum imaging at 1.3 and 0.7 cm of four hypercompact H II regions in W3 IRS 5, suggested that these sources contain B or O stars (Wilson et al. 2003).  From a recent near-infrared (NIR) survey of a $\\sim$ 1\\arcmin.5$\\times$1\\arcmin.5 region towards W3 Main, Megeath et al. (1996) found a dense concentration of stars in the molecular clump surrounding W3 IRS 5. These data showed a large, embedded population of intermediate to low mass stars co-existing with recently formed OB stars. They also argued that the formation of high mass stars is associated with the formation of dense clusters of low mass stars in the W3 Main star forming region. Several X-ray sources were detected in the W3 core by Chandra X-ray Observatory (Hofner et al. 2002). Most of these sources are located at the peak radio positions of the W3 H II regions. Hofner et al. (2002) postulated that the X-ray sources are the young massive stars that are also responsible for the ionization of the compact and ultracompact H II regions in the W3 core.  In this paper we present deep J, H, and K$_s$-bands NIR observations of the W3 Main star forming region. In comparison with the previous NIR survey (Megeath et al. 1996), our survey covers a larger area ($\\sim$ 24 arcmin$^2$) surrounding W3 IRS 5, including compact H II regions W3 A, W3 B and W3 D, diffuse H II regions W3 H, W3 J and W3 K, and the cometary ultracompact (UC) H II regions W3 C, W3 E and W3 F. These individual H II regions have been labeled following the scheme introduced by Wynn-Williams (1971) and Harris \\& Wynn-Williams (1976). Our motivation is to look for new young stellar objects (YSOs) associated with the W3 Main region, to determine their evolutionary stages, and to discuss their nature. Tieftrunk et al. (1998) have presented the three 10\\arcmin$\\times$10\\arcmin ~mosaics ($\\sim$ 300 arcmin$^2$) in K' filter of W3 region. This mosaic has similar spatial resolution and depth to the SIRIUS K$_s$ image and covers the region stretching from W3 Main to W3(OH). However, due to the crowded nature of the sources and the large pixel size, the K'-band mosaic was not suitable for photometry of the embedded stellar clusters. In Sects. 2 and 3 we present the details of observations and data reduction procedures, Sect. 4 deals with the results and discussion and we summarize our conclusions in Sect. 5. ",
        "conclusions": "A deep JHK$_s$-band NIR imaging survey of YSOs associated with the W3 Main star forming region is presented. The survey covers a 4\\arcmin.9$\\times$4\\arcmin.9 area down to a limiting magnitude (10 $\\sigma$) of J = 19.0, H = 18.1, and K$_s$ = 17.3. From the analysis of these images we derive the following conclusions :  1) A cluster of YSOs (Class II and Class I sources derived from their NIR colors) has been detected in the W3 Main core and near the compact, ultracompact, and diffuse H II regions.  2) A large number of red stars (H-K $>$ 2) are detected in the molecular cloud region, most of them clustered around the molecular clumps associated with IRS 5 and IRS 4. Some of them are also associated with the diffuse emission near the dense molecular clumps. We argue that most of the reddest stars are YSOs with circumstellar materials.  3) The KLF of the W3 Main region shows the power-law slope : $\\alpha$ = 0.26$\\pm$0.02, which is lower than the typical values reported for the embedded young clusters. Our finding also confirms the previous results of Megeath et al. (1996) for a smaller region around W3 IRS 5.  4) The observed density of the cluster region around W3 IRS 5 is $\\sim$ 2000 stars pc$^{-3}$ for K $<$ 17.5, which is larger than the typical values ($\\sim$ 1000 stars pc$^{-3}$) reported for other embedded clusters.  5) Using the age of W3 Main in the range of 0.3--1 Myr determined by Megeath et al. (1996), we find that about 80\\% of the YSO candidates have an upper mass limit of 4 M$_{\\odot}$. We estimate that the lowest mass limit of Class II \\& Class I candidates in our observations is 0.1 M$_{\\odot}$. Therefore, the stellar population in W3 Main is primarily composed of low mass PMS stars."
    },
    "0403/astro-ph0403280_arXiv.txt": {
        "abstract": "After a brief introduction to the necessary of {\\em quark stars} in modelling pulsars, I present a qualitative analysis of the {\\em solidification} of quark matter with low-temperature but high-density. The reason, that a solid neutron star could {\\em not} be possible, is also given. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403555_arXiv.txt": {
        "abstract": "{ We analyse the evolution of the fractional ionisation in a steady-state protoplanetary disc over $10^6$~yr. We consider a disc model with a vertical temperature gradient and with gas-grain chemistry including surface reactions. The ionisation due to stellar X-rays, stellar and interstellar UV radiation, cosmic rays and radionuclide decay is taken into account. Using our reduction schemes as a tool for the analysis, we isolate small sets of chemical reactions that reproduce the evolution of the ionisation degree at  representative disc locations with an accuracy of 50\\%--100\\%. On the basis of fractional ionisation, the disc can be divided into three distinct layers. In the dark dense midplane the ionisation degree is sustained by cosmic rays and radionuclides only and is very low, $\\la 10^{-12}$. This region corresponds to the so-called ``dead zone'' in terms of the angular momentum transport driven by MHD turbulence. The ionisation degree can be accurately reproduced by chemical networks with about 10 species and a similar number of reactions. In the intermediate layer the chemistry of the fractional ionisation is driven mainly by the attenuated stellar X-rays and is far more complicated. For the first time, we argue that surface hydrogenation of long carbon chains can be of crucial importance for the evolution of the ionisation degree in protoplanetary discs. In the intermediate layer reduced networks contain more than a 100 species and hundreds of reactions. Finally, in the unshielded low-density surface layer of the disc the chemical life cycle of the ionisation degree comprises a restricted set of photoionisation-recombination processes. It is sufficient to keep about 20 species and reactions in reduced networks. Furthermore, column densities of key molecules are calculated and compared to the results of other recent studies and observational data. The relevance of our results to the MHD modelling of protoplanetary discs is discussed. ",
        "introduction": "Nowadays a paradigm of the evolution of a protoplanetary disc is widely accepted in which most of the disc matter is assumed to move steadily toward a protostar due to redistribution of the angular momentum. Ionisation in such objects is an important factor that enables the angular momentum transport to occur via magnetohydrodynamic (MHD) turbulence driven by the magnetorotational instability (MRI; Balbus \\& Hawley~\\cite{MRI}). From this point of view, a disc is conventionally divided into the ``active'' layer, adjacent to the disc surface, and the ``dead'' zone, centred on the midplane. The active part of the disc is irradiated by high energy stellar/interstellar photons. Thus, the fractional ionisation there is relatively high, which implies that the magnetic field is well coupled to the gas. Due to this coupling, the active layer is unstable to the MRI, and the developing turbulence allows the accretion to occur (e.g., Gammie~\\cite{gammie}; Fleming \\& Stone~\\cite{fs2003}). The shielded midplane region is almost neutral, decoupled from the magnetic field, and, thus, quiescent.  The location of the boundary between these two regions may prove to be very sensitive to the disc physical properties and chemical composition (e.g., Fromang, Terquem, \\& Balbus~\\cite{from2002}). However, the MHD modelling with non-ideal effects included is a very demanding computational task. This is why in the MHD modelling of protoplanetary discs (and protostellar clouds) a very simple chemical scheme is usually assumed with a few ions and a network that includes only ionisation and recombination reactions, often neglecting the presence of dust grains (Sano et al.~\\cite{sano}; Fromang et al.~\\cite{from2002}; Fleming \\& Stone~\\cite{fs2003}). The medium is believed to be in chemical equilibrium so that the fractional ionisation $x_\\mathrm{e}$ can be expressed as (e.g., Gammie~\\cite{gammie}) \\begin{equation} x_\\mathrm{e}(\\mathrm{eq})=\\sqrt{\\zeta\\over\\beta n_\\mathrm{H}}, \\label{xe} \\end{equation} where $\\zeta$ is the ionisation rate, $\\beta$ is a typical recombination coefficient, and $n_\\mathrm{H}$ is the hydrogen number density. This approach may indeed be valid if only the cosmic ray ionisation is taken into account and only gas-phase chemical processes are considered. However, dust plays an important role in the evolution of the fractional ionisation being an efficient electron donor for recombining ions and a sink for neutrals in cold parts of a disc. Moreover, newly born stars possess a relatively high X-ray flux with photon energies from about 1 to 5 keV (e.g. Igea \\& Glassgold~\\cite{IG99}). These factors give rise to a more complicated chemistry relevant for the ionisation degree. In particular, different ions dominate the fractional ionisation at different times and in different parts of the disc. Therefore, the straightforward application of Eq.~(\\ref{xe}) for evaluating the ionisation degree can be fraught with errors in some parts of a disc.  To check the validity of Eq.~(\\ref{xe}), we analyse the detailed ionisation structure of a protoplanetary disc computed with the full UMIST\\,95 chemical network (Millar et al.~\\cite{umist95}) with the surface chemistry included. We adopt a reference disc model to serve as a guide to the range of possible physical conditions that may be encountered in real protoplanetary objects. Within this model, we choose several representative disc locations and investigate in detail which chemical processes control the time-dependent ionisation degree there. Our intention is to show that the oversimplified treatment of the ionisation in a protoplanetary disc as an equilibrated ionisation-recombination cycle can lead to $x_\\mathrm{e}$ values that differ by more than an order of magnitude from values computed with all the available information on the disc chemical evolution.  We describe the chemical processes that are responsible for ionisation in a disc in terms of reduced networks that are subsets of the full UMIST\\,95 network, supplied in some cases by a few surface reactions. These networks contain only those species and reactions that are needed to reproduce $x_{\\rm e}$ with up to a factor of 2 uncertainty. The utilised reduction methods are described in Wiebe, Semenov \\& Henning (\\cite{papi}, hereafter Paper~I). The species-based reduction rests upon the Ruffle et al.~(\\cite{Rea02}) technique and consists of choosing species that are important in a particular context, and then selecting from the entire network only those species that are necessary to compute abundances of important species with a reasonable accuracy. In the reaction-based method, analysis starts from reactions that govern the abundance of important species. All reactions in the entire network are assigned weights according to the influence they have on abundances of important species. Then, only those reactions are selected that have weights above a cut-off parameter that is selected on the basis of the requested accuracy. Only those species are included in the reduced network that participate in selected reactions.  The organisation of the paper is the following. In Section~\\ref{mod} we describe the disc model and its physical structure as well as updates to the chemical model of Paper~I. In Section~\\ref{icon} we discuss the initial conditions for the disc chemistry. The processes responsible for the fractional ionisation in various parts of the disc are outlined in Section~\\ref{res}. In Section~\\ref{colden} column densities are tabulated and compared to other studies and observational data. Results of the analysis and their relevance to the MHD modelling are discussed in Section~\\ref{diss}. Final conclusions are drawn in Section~\\ref{concl}. ",
        "conclusions": "\\label{concl}  Chemical processes, responsible for the ionisation structure of a protoplanetary disc with a central star, are analysed by means of reduced networks that reproduce the ionisation fraction within a factor of 2. These networks are available from the authors upon request. Because of the wide range of physical conditions met in a typical disc, there is a corresponding diversity in the chemical reactions that control the fractional ionisation in different parts of the disc. Generally, it can be divided into three layers. In the midplane the ionisation is provided by cosmic rays and radioactive elements only. Above the midplane, the intermediate layer is located where the ionisation is dominated by X-rays. In the surface layer UV photons are the main ionising factor. In each of these layers we analyse several representative points and construct reduced chemical networks that are needed to reproduce the fractional ionisation as a function of time during $10^6$ years of evolution within a factor of 2~uncertainty. In the midplane the chemistry, which determines the fractional ionisation, is very simple. Reduced networks typically include no more than 10 species and a similar number of reactions. The same holds true for the surface layer. The intermediate layer is the most complicated region from the chemical point of view. To compute the fractional ionisation with the targeted accuracy, in some regions one has to take into account carbon chains containing up to 6~carbon atoms, which leads to reduced networks with over a hundred species and reactions."
    },
    "0403/astro-ph0403024_arXiv.txt": {
        "abstract": "The quantum model of the homogeneous and isotropic universe predicts logarithmic-law dependence of coordinate distance to source on redshift $z$ which is in good agreement with type Ia supernovae and radio galaxies observations for the redshift range $z = 0.01 - 1.8$. A comparison with phenomenological models with dark energy in the form of cosmological constant and without dark energy component is made. Fluctuations of the cosmological scale factor about its average value which can arise in the early universe produce accelerating or decelerating expansions of space subdomains containing separate sources with high redshift whereas the universe as a whole expands at a steady rate. ",
        "introduction": "\\label{Intro}  The high-redshift type Ia supernova (SN Ia) observations \\cite{Rie,Per,Ton} can be explained by an accelerating expansion of the present-day universe \\cite{Rie,Per,Tu1,Pee,Liv}. Such conclusion assumes that observed dimming of the SNe Ia is hardly caused by physical phenomena non-related to overall expansion of the universe as a whole (see e.g. Refs. \\cite{Ton,Vish} for discussion and review). Furthermore it is supposed that matter component of energy density in the universe $\\rho_{M}$ varies with the expansion of the universe as $a^{-3}$ (i.e. it has practically vanishing pressure, $p_{M} \\approx 0$), where $a$ is the cosmological scale factor, while mysterious cosmic fluid (so-called dark energy \\cite{Tu2,Ost}) is describes by the equation of state $p_{X} = w_{X} \\rho_{X}$, where $-1 \\leq w_{X} \\leq - \\frac{1}{3}$ \\cite{Tu1,Pee}. In the models with the cosmological constant ($\\Lambda$CDM) one has $w_{X} = - 1$ \\cite{Ton,Tu1,Pee}, while in general case $w_{X}$ may vary with time (models with quintessence) \\cite{Pee,Liv}. Even if regarding baryon component one can assume that it decreases as $a^{-3}$ (pressure of baryons may be neglected due to their relative small amount in the universe), for dark matter (whose nature and properties can be extracted only from its gravitational action on ordinary matter) such a dependence on the scale factor may not hold in the universe taken as a whole (in contrast to local manifestations e.g. in large-scale structure formation, where dependence $a^{-3}$ may survive). Since the contribution from all baryons into the total energy density does not exceed 4 \\% \\cite{Hag}, the evolution of the universe as a whole is determined mainly by the properties of dark matter and dark energy. The models of dark energy \\cite{Tu1,Pee,Tu2} show explicitly unusual behaviour of this component during the expansion of the universe.  In the present paper we notice that the quantum model of the homogeneous and isotropic universe filled with primordial matter in the form of the uniform scalar field proposed in Refs. \\cite{K,KK} allows to explain the observed coordinate distances to SNe Ia and radio galaxies (RGs) in wide redshift range. ",
        "conclusions": "\\label{Con}  In addition to the prediction about the steady-speed expansion of the universe as a whole (at the same time the accelerating or decelerating motions of its subdomains remain possible on a cosmological scale as it is shown in Sect. \\ref{Quan}) the quantum model allows an increase of quantity of matter/energy in matter dominated universe according to (\\ref{3}). If the mass $m$ of elementary excitations of the scaler field remains unchanged during the expansion of the universe, then the increase of $M$ can occur due to increase in number $s$ of these excitations. But the increase in $s$ does not mean that a quantity of observed matter in some chosen volume of the universe increases. According to the model proposed in Refs. \\cite{KK2,KK3} the observed ``real'' matter (both luminous and dark) is created as a result of the decay of excitations of the scalar field (under the action of gravitational forces) into baryons, leptons and dark matter. The undecayed part of them forms what can be called a dark energy. Such a decay scheme leads to realistic estimates of the percentage of baryons, dark matter and energy in the universe with $\\langle a \\rangle \\gg 1$ and $M \\gg 1$. Despite the fact that the quantity of matter/energy can increase, the mean total energy density decreases and during the expansion of the universe mainly the number of elementary excitations of the scalar field increases. Their decay probability is very small, so that basically only the dark energy is created. These circumstances can explain the absence of observed events of creation of a new baryonic matter on a cosmologically significant scale.  The proposed approach to the explanation of observed dimming of some SNe Ia may provoke objections in connection with the problem of large-scale structure formation in the universe, since the energy density $\\langle \\rho \\rangle$ in the form (\\ref{4}) cannot ensure an existence of a growing mode of the density contrast $\\delta \\langle \\rho \\rangle/\\langle \\rho \\rangle$ (see e.g. Refs. \\cite{Ol,We,Pee2}). As we have already mentioned above in Sect. \\ref{Coor} the density $\\langle \\rho \\rangle$ (\\ref{4}) describes only homogenized properties of the universe as a whole. It cannot be used in calculations of fluctuations of energy density about the mean value $\\langle \\rho \\rangle$. Under the study of large-scale structure formation one should proceed from the more general expression for the energy density (\\ref{2}). Defining concretely the contents of matter/energy $M$, as for instance in the model of creation of matter mentioned above, one can make calculations of density contrast as a function of redshift. The problem of large-scale structure formation is one of the main problems of cosmology (see e.g. Refs. \\cite{Ol,Dol}). It goes beyond the tasks of this paper and requires a special investigation. The ways of its solution in the quantum model are roughly outlined in Ref. \\cite{KK}."
    },
    "0403/astro-ph0403212_arXiv.txt": {
        "abstract": "Dense high-precision photometry of microlensed stars during a fold-caustic passage can be used to reveal their brightness profiles from which the temperature of the stellar atmosphere as function of fractional radius can be derived. While the capabilities of current microlensing follow-up campaigns such as PLANET allowed for several precise measurements of linear limb-darkening coefficients, all attempts to reveal a second limb-darkening coefficient from such events have failed. It is shown that the residual signal of a second coefficient characterizing square-root limb darkening is $\\sim\\,$25 times smaller which prevents a proper determination except for unlikely cases of very high caustic-peak-to-outside magnification ratios with no adequate event being observed so far or for source stars passing over a cusp singularity. Although the presence of limb darkening can be well established from the data, a reliable measurement of the index of an underlying power-law cannot be obtained. ",
        "introduction": "Since \\citet{SW1987} have proposed to use caustics of gravitational lenses for measuring the brightness profile of closely-aligned background sources, dense high-precision photometry on microlensing events involving fold-caustic passages has provided measurements of linear limb-darkening coefficients for several stars \\citep{PLANET:M41, PLANET:O23, joint, PLANET:EB5}.  For stellar brightness profiles that involve linear, square, or square-root terms in $\\cos \\vartheta$, where $\\vartheta$ is the emergent angle, properties of light curves during fold-caustic passages have been studied by \\citet{Rhie:LD}, while prospects and strategies for determining the linear limb-darkening coefficient have recently been discussed in detail by \\citet{Do:FoldLD}, showing the ability of obtaining precise measurements even with moderate use of the current capabilities.  This letter will however show that the measurement of additional coefficients, e.g.\\ corresponding to a square-root law term, is not possible for typical events that involve a fold-caustic passage. Attempts of such measurements \\citep{PLANET:EB5} have resulted in the meaningful measurement of a single linear combination of the involved limb-darkening coefficients only, while a similar result has been obtained for an event where the source passes over a cusp \\citep{Abe}, where however the exhibited differential magnification is modest due to the source size being relatively large compared to the cusp. In contrast, the cusp-passage event discussed by \\citet{PLANET:M28} provides a better size ratio between source and cusp and to date presents the only measurement of both linear and square-root limb-darkening coefficients by microlensing. ",
        "conclusions": "If the stellar brightness profile $\\xi(\\rho)$ is modelled by a linear superposition of three base profile functions, as for the discussed examples of constant, linear, and square-root terms in $\\cos \\vartheta$, the normalization forces one of these profiles to mediate between the other two, so that this base profile can be approximated by a superposition of the others. With several fractional radii $\\rho$ being probed by the fold caustic for any passage phase $\\eta$ according to ${\\bmath{\\mathcal T}} (\\eta,\\rho)$, the residuals of this approximation are smaller for the caustic profile function $G_\\rmn{f}^\\star(\\eta; \\xi)$ than for the stellar brightness profile $\\xi(\\rho)$, which limits the power of photometric microlensing observations during fold-caustic passages for revealing the stellar brightness profile $\\xi(\\rho)$.  Compared to the determination of a linear limb-darkening coefficient, the residual signals of an additional square-root limb-darkening coefficient are $\\sim\\,$25 times smaller, making its measurement with current microlensing follow-up campaigns such as PLANET impossible, unless the caustic-peak-to-exit magnification ratio becomes $F_\\rmn{peak}/F_\\rmn{f}^\\star \\ga 40$ for a (typical) photometric accuracy  $\\sigma_\\rmn{f}^\\star/F_\\rmn{f}^\\star = 1.5\\,$\\% at the caustic exit where none of such events have been observed so far, or the source passes over a cusp singularity rather than a fold.  If more than one limb-darkening base profile is assumed to contribute, only a single characteristic coefficient can be accurately measured which corresponds to a specific superposition of the corresponding base profiles as found for worked examples \\citep{PLANET:EB5,Abe}."
    },
    "0403/astro-ph0403448_arXiv.txt": {
        "abstract": "We report here the first evidence for planetesimal infall onto the very young Herbig Be star LkH$_\\alpha$234. These results are based on observations acquired over 31 days using spectroscopy of the sodium D lines, the He I 5876\\AA, and hydrogen H$_\\alpha$ lines. We find Redshifted Absorption Components (RAC) with velocities up to 200 km/s and very mild Blueshifted Absorption Components (BEC) up to 100 km/s in the Na I lines. No correlation is observed between the appearance of the Na I RAC \\& BEC and the H$_\\alpha$ and He I line variability, which suggests that these (Na I RAC \\& BEC) are formed in a process unrelated to the circumstellar gas accretion. We interpret the Na I RAC as evidence for an infalling evaporating body, greater than 100 km in diameter, which is able to survive at distances between 2.0 to 0.1 AU from the star. The dramatic appearance of the sodium RAC and mild BEC is readily explained by the dynamics of this infalling body making LkH$_\\alpha$234 the youngest (age $\\sim$ 0.1 Myr) system with evidence for the presence of planetesimals. ",
        "introduction": "LkH$_\\alpha$234 (m$_v$=11.9mag) is a star embedded in a nebula associated with the NGC 7129 star forming region and is at a distance of about 1250 parsecs. This is one of the youngest known Herbig Be objects and has an age $\\sim$ 0.1 million years and spectral type of B5 (Fuente et al. 2001). Herbig Ae/Be stars are intermediate mass (2 to 10M$_\\odot$) pre-main-sequence stars (Herbig 1960). A number of Herbig Ae stars have been found to possess circumstellar disks and show both photometric and spectroscopic variability (Herbst et al. 1999). LkH$_\\alpha$234 though much hotter and more massive also exhibits these properties (Fuente et al. 2001, Herbst et al. 1999, Polomski et al. 2002). The star shows photometric variability of 1.2 magnitude in the V band and is known posses a circumstellar disk with the disk mass to the stellar mass ratio (M$_D$/M$_*$) $\\leq$ 0.02. A detailed description of LkH$_\\alpha$234 and its circumstellar region can be found in Fuente et al. (2001, and references therein).  Spectroscopic variations in metallic lines (like Na I, Ca II, etc.) in young A type stars like $\\beta$ Pictoris (20 to 100 million years) are often interpreted as the signature of star grazing comets or falling evaporating bodies with sizes of tens of kilometers (Grinin et al. 1996, Smith \\& Terrile 1984, Thebault \\& Beust 2001). The only Herbig Be star where grazing cometary transient events have been inferred so far is HD 100546 (Viera et al. 1999). However, HD100546 is much older ($\\geq$10 Myrs) and cooler (spectral type $\\sim$B9) compared to LkH$_\\alpha$234. High resolution spectroscopic monitoring of very young (0.1 to 1 Myrs) Herbig Be stars are scanty and uncharted.  Here we report dramatic variations in the spectroscopic lines of Na I 5890\\AA (D2), and 5896\\AA (D1), He I 5876\\AA, and H$_\\alpha$ of LkH$_\\alpha$234 over a period of 31 days during Oct.-Nov. 2003. The observations consists of high resolution spectra (R=30,000) from the Hobby-Eberly Telescope (HET) using the High Resolution Spectrograph (HRS). Section 2 describes observations and in section 3 we discuss the results and section 4 gives the conclusion. ",
        "conclusions": "We have presented here high resolution spectroscopic monitoring of a Herbig Be star LkH$_\\alpha$234 in the Na I (D2\\&D1), He I (5876\\AA) and H$_\\alpha$ lines. We found no correlations between the Na I and the He I and H$_\\alpha$ line variability in the five data sets over a period of 31 days (7th Oct. 2003 to 8th Nov. 2003) though we do find a correlation between He I and H$_\\alpha$. Thus the origin of variation in the Na I line (RAC up to 200 km/s and BEC up to 100 km/s seen only on 13th Oct. 2003) is different from those of H$_\\alpha$ and He I. While the H$_\\alpha$ and He I line variations are due to episodic gaseous accretion, the Na I RAC indicate a dramatic transient event of solid body infall on to the star and the BEC as the blown away parts of the solid body not projected against the stellar surface. Considering the Keplerian dynamics and the harsh environment of a B5 star, we estimate that a solid body of size $\\geq$100 km broke up and disintegrated at a distance between 0.1 to 2.0 AU from the star. This makes LkH$_\\alpha$234 the youngest system ($\\sim$0.1 Myrs) with evidence for protoplanetary bodies of asteroidal size. We plan to pursue further spectroscopic monitoring of the star with the HET to determine the frequency of such events and further understand the complicated dynamics of this very young circumstellar environment."
    },
    "0403/physics0403019_arXiv.txt": {
        "abstract": "In MicroPattern Gas Detectors (MPGD) when the pixel size is below 100 $\\mu$m and the number of pixels is large (above 1000) it is virtually impossible to use the conventional PCB read-out approach to bring the signal charge from the individual pixel to the external electronics chain. For this reason a custom CMOS array of 2101 active pixels with 80 $\\mu$m pitch, directly used as the charge collecting anode of a GEM amplifying structure, has been developed and built. Each charge collecting pad, hexagonally shaped, realized using the top metal layer of a deep submicron VLSI technology is individually connected to a full electronics chain (pre-amplifier, shaping-amplifier, sample \\& hold, multiplexer) which is built immediately below it by using the remaining five active layers. The GEM and the drift electrode window are assembled directly over the chip so the ASIC itself becomes the pixelized anode of a MicroPattern Gas Detector. With this approach, for the first time, gas detectors have reached the level of integration and resolution typical of solid state pixel detectors. Results from the first tests of this new read-out concept are presented. An Astronomical  X-Ray Polarimetry application is also discussed. ",
        "introduction": "The most interesting feature of the Gas Electron Multiplier (GEM) is the possibility of full decoupling of the charge amplification structure from the read-out structure. In this way both can be independently optimized. Indeed, by organizing the read-out plane in a multi-pixel pattern it is possible to get a true 2D imaging capability. At the same time a high granularity of the read-out plane would also allow to preserve the intrinsic resolving power of the device and its high rate capability that otherwise would be unavoidably lost by using a conventional projective read-out approach. However, when the pixel size is small (below 100 $\\mu$m) and the number of pixels is large (above 1000) it is virtually impossible to bring the signal charge from the individual pixel to a chain of external read-out electronics even by using an advanced, fine-line, multi-layer, PCB technology. The fan-out which connects the segmented anodes collecting the charge to the front-end electronics is the real bottleneck. Technological constraints limit the maximum number of independent electronics channels that can be brought to the peripheral electronics. Furthermore,  the crosstalk between adjacent channels and the noise due to the high input capacitance to the preamplifiers become not negligible. In this case, it is the electronics chain that has to be brought to the individual pixel. We have implemented this concept by developing and building a CMOS VLSI array of 2101 pixels with 80 $\\mu$m pitch which is used directly as the charge collecting anode of the GEM. A description of the read-out ASIC for a MPGD and of its advantages is given in the next section. Section 3 describes the coupling of the chip die to the amplifying electrode, the assembly of the full detector and the results of laboratory tests obtained with a 5.9 keV X-ray source. The use of this new detection concept for Astronomical X-Ray Polarimetry and other applications are discussed in the last section. ",
        "conclusions": "A system in which the GEM foil, the absorption gap and the entrance window are assembled directly over a custom CMOS chip die has been developed. The transfer of  charge from the amplifying region to the collection and read-out region occurs via electric fields. The ASIC itself becomes at the same time, the charge collecting anode and the pixelized read-out of a MicroPattern Gas Detector. For the first time the full electronics chain and the detector are completely integrated without the need of complicated bump-bonding. At a gain of 1000 a high sensitivity to single primary electron detection is reached. No problems have been found up to now in operating the system under HV and in a gas environment. An astronomical X-ray Polarimeter application has been presented. Final design will have 16$\\div$32 k channels and 60$\\div$70 microns pixel size ($\\simeq 1 cm^{2}$ active area). Depending on pixel and die size, electronics shaping time, analog vs. digital read-out, counting vs. integrating mode, gas filling, many others applications can be envisaged. This would open new directions in gas detector read-out, bringing the field to the same level of integration of solid state detectors."
    },
    "0403/astro-ph0403162.txt": {
        "abstract": "We present a detailed analysis of a 3.5 s long burst from SGR1900+14 which occurred on 2001 July 2.  The 2-150 keV time-integrated energy spectrum is well described by the sum of two blackbodies whose temperatures are approximately 4.3 and 9.8 keV.  The time-resolved energy spectra are similarly well fit by the sum of two blackbodies. The higher temperature blackbody evolves with time in a manner consistent with a shrinking emitting surface. The interpretation of these results in the context of the magnetar model suggests that the two blackbody fit is an approximation of an absorbed, multi-temperature spectrum expected on theoretical grounds rather than a physical description of the emission. If this is indeed the case, our data provide further evidence for a strong magnetic field, and indicate that the entire neutron was radiating during most of the burst duration. ",
        "introduction": "\\label{intro}  % KEVIN CAN YOU CHECK THE REFERRE's COMMENTS ON THE INTRO, AND ANSWER THEM...  The soft gamma repeater \\sgr\\ was discovered in 1979 when it emitted 3 short bursts of soft gamma-rays in 3 days (Mazets et al. 1979).  Its next recorded appearance occurred some 13 years later (Kouveliotou et al. 1993).  Attempts were made over the years to obtain a precise position for the source and identify its counterpart (Hurley et al. 1994; Vasisht et al.  1994; Hurley et al. 1996), but this remained elusive until 1998, when the source entered a new period of activity, allowing it to be localized accurately by the interplanetary network (IPN: Hurley et al. 1999a).  The precise source location was found to be consistent with that of a previously identified ROSAT quiescent X-ray source (Vasisht et al. 1994; Hurley et al. 1996).  Observations with ASCA further revealed that the source had a 5.16 s period (Hurley et al. 1999b), and observations with RXTE demonstrated that the period was increasing rapidly ($\\rm 1.1 \\times 10^{-10}$ s s$^{-1}$, Kouveliotou et al. 1999). The counterpart to \\sgr~has not been found yet.  If it is associated with the Galactic supernova remnant G42.8+0.6, it could be as close as 5 kpc (Hurley et al. 1999b).  However, the source position also appears to be very close to a cluster of high mass stars, and it has been proposed that this may be the birthplace of the neutron star (Vrba et al. 2000).  If so, its distance could be roughly 12-15 kpc. In all this paper we adopt a distance of 10 kpc.  The 1998 activity of \\sgr~culminated in the giant flare of August 27 (Frail et al. 1999; Hurley et al. 1999c; Feroci et al. 1999; Mazets et al. 1999), which was followed by numerous, smaller bursts (e.g. Ibrahim et al. 2001). The next major period of activity of \\sgr~came in 2001 (Guidorzi et al. 2001; Hurley et al.  2001a, b; Feroci et al. 2001; Ricker et al. 2001a, b, c; Montanari et al. 2001, Feroci et al. 2004). During this episode, it became apparent that this source emits not only the common, short SGR bursts, with durations of about 200 ms and fluences  $\\gtrsim 10^{-6} \\rm erg \\, cm^{-2}$, and the much rarer giant flares, lasting for minutes and having fluences $\\gtrsim 10^{-3} \\rm erg \\, cm^{-2}$, but also, a class of high-fluence \\it intermediate \\rm bursts, whose durations and fluences fall somewhere in between, and may in fact form a continuum (Kouveliotou et al. 2001; Woods et al. 2003; Feroci et al. 2003).  Such events had been observed in the aftermath of the 1998 August 27 giant flare, but were thought to be related to it.  In retrospect, it seems likely that such bursts were also emitted by SGR0525-66, in the aftermath of the famous 1979 March 5 giant flare (see Golenetskii et al. 1984). The properties of these intermediate bursts are interesting for numerous reasons, not the least of which are that some display X-ray afterglows, similar to the one observed during the 1998 August 27 event (Feroci et al. 2001; Thompson and Duncan et al. 2001), and that their longer duration permits studies of their spectral evolution. %OLD possess unique spectral properties.  Duncan and Thompson (1992), Paczy\\'nski (1992), and Thompson and Duncan (1995, 1996) have proposed that the soft gamma repeaters are \\it magnetars \\rm, i.e., neutron stars with magnetic fields B $\\rm \\approx 10^{15} G$.  In this model, magnetic dissipation causes the neutron star crust to fracture, and Alfv\\'en waves accelerate electrons, resulting in short (200 ms) bursts of soft gamma-radiation. Much more rarely, magnetic reconnection provides the energy for a longer, extremely energetic giant flare involving the entire neutron star magnetosphere.  The periodicity observed in both the quiescent soft X-ray emission from \\sgr~and in the giant flare, the high spin-down rate, and the energetics of the giant flare, are all consistent with the main features of the magnetar model.  In this paper, we analyze an intermediate burst from \\sgr , which occurred on 2001 July 2, emphasizing broad band and time-resolved X-ray spectral modeling obtained using the data of the FREGATE (FREnch GAmma-ray TElescope) and WXM (Wide Field X-ray Monitor) experiments aboard the HETE (High Energy Transient Explorer) spacecraft. ",
        "conclusions": "We have presented HETE-FREGATE and WXM observations of an intermediate burst from SGR1900+14.  This is the first time that time-resolved spectral analysis of such a burst has been possible over a wide energy range, with good spectral and temporal resolution. We have found that the spectrum is well described by the sum of two blackbodies. One possible interpretation is that we are observing two distinct emitting volumes.  Another is that we are observing emission which is affected by radiation transfer effects in a superstrong magnetic field. Both interpretations are consistent with the basic features of the magnetar model of soft gamma repeaters.  However, it is not possible to choose between the two on the basis of a single observation.  Fortunately, HETE-2 has observed numerous SGR bursts from both SGR1900+14 and SGR1806-20. Analysis of these observations is underway, and may provide less ambiguous support for one of these hypotheses. % Thus, this intermediate burst may % be viewed as a scaled-down giant flare. % JLA:SHOULD WE KEEP THAT ??? % JFO:Je crois pas."
    },
    "0403/astro-ph0403118_arXiv.txt": {
        "abstract": "ESPaDOnS is a new-generation cross-dispersed \\'echelle spectropolarimeter, the commissioning phase of which is scheduled at the Canada-France-Hawaii Telescope (CFHT) for spring 2004.  This instrument will provide full coverage of the optical domain (370~nm to 1,000~nm) in all polarization states (circular and linear) at a resolving power of about 70,000, with a peak efficiency of 20\\% (telescope and detector included).  It includes a bench-mounted spectrograph, fiber-fed from a Cassegrain-mounted module including all polarimetric and calibration facilities.  ESPaDOnS should be the most powerful tool dedicated to stellar spectropolarimetry, therefore opening unprecedented perspectives for major issues of stellar physics, from studies of stellar interiors to investigations of stellar atmospheres, stellar surfaces, stellar magnetic fields, and to observations of circumstellar environments and extra-solar planets. ",
        "introduction": "The combination of a polarimeter and a high-resolution \\'echelle spectrograph constitutes a powerful tool for investigating the physics of stellar atmospheres and environments. In particular, stellar magnetic fields are mostly detected by their polarizing effect on spectral lines. The analysis of this so-called ``Zeeman effect'' requires polarized, high resolution spectra, so that polarization across spectral line profiles can be investigated in detail, in order to measure both field strength and orientation. A spectral resolution in excess of 50,000 (to obtain at least two resolved elements in the thermal width of a typical spectral line) and the possibility to record linearly and circularly polarized spectra are therefore requested for optimal investigations of stellar magnetism. In addition to high resolution, a wide spectral coverage (preferentially covering the whole optical domain) is needed by spectroscopic studies (polarimetric or not) searching for tiny spectral signatures, the detection of which can greatly benefit from the rise in S/N produced by a simultaneous extraction of the signal in many spectral lines.  To fulfill this range of ambitious technical performances, the ESPaDOnS concept was proposed by Donati et al. (1998). The instrument, developed by a collaboration including French, Canadian and Dutch experts, is now in its final assembling and testing phase at Observatoire Midi-Pyr\\'en\\'ees (France) before its first light scheduled at the CFHT during the second half of the year.  We first describe the main characteristics of this instrument, which includes a polarimeter mounted on the Cassegrain bonnette of the telescope, fiber linked to a bench mounted spectrograph. We then evoke some of the most promising scientific issues that such a unique facility will make it possible to address in a near future. ",
        "conclusions": ""
    },
    "0403/astro-ph0403432_arXiv.txt": {
        "abstract": "In February 1997 the Japanese radio astronomy satellite Halca was launched to provide the space-bourne element for the VSOP mission. Approximately twenty-five percent of the mission time has been dedicated to the VSOP Survey, a 5~GHz survey of bright, compact AGN. We present the results from the ongoing analysis. Both the final, calibrated, high resolution images and plots of visibility amplitude versus {\\em uv} distance for the first 102 of the sources have been prepared and has been submitted. Papers on the methods and the models from fitting the cumulative {\\em uv} amplitudes will also be submitted. The analysis of the second half is well underway. ",
        "introduction": "The radio astronomy satellite HALCA was launched by the former Institute of Space and Astronautical Science (now part of Japanese Aerospace eXpolaration Agency (JAXA))in February 1997 to participate in Very Long Baseline Interferometry (VLBI) observations with arrays of ground radio telescopes. It was was placed in an orbit with an apogee height above the Earth's surface of 21,400\\,km, a perigee height of 560\\,km, and an orbital period of 6.3~hours. HALCA provides the longest baselines of the VLBI Space Observatory Programme (VSOP), an international endeavor that has involved over 25 ground radio telescopes, five tracking stations and three correlators (Hirabayashi et al. 1998; 2000a).   HALCA has now passed the end of the Guaranteed Observing Time period, and with the completion of the Memorandum of Understanding in February 2002 the NASA tracking and orbital calculation support ceased and the observation program has turned to completing the Survey. The orbital determination and spacecraft tracking are now completely indigenous to Japan and ISAS. ",
        "conclusions": "The satellite continues to make survey observations, and will do so while the satellite is functioning. These observations are being analysed cumulatively and individually and providing interesting results.  These results are not only important for the understanding the target sources, mainly AGN's, but also for the planning of future space-VLBI missions such as the VSOP-2 mission, which will have a resolution of nearly a magnitude better."
    },
    "0403/astro-ph0403604_arXiv.txt": {
        "abstract": "We evaluate the logarithmic derivative of the depth of the solar convective zone with respect to the logarithm of the radiative opacity, $\\partial \\ln R_{\\rm CZ}/\\partial \\ln \\kappa$. We use this expression to show that the radiative opacity near the base of the solar convective zone (CZ) must be known to an accuracy of $\\pm 1$\\% in order to calculate the CZ depth to the accuracy of the helioseismological measurement, $R_{\\rm CZ} = (0.713 \\pm 0.001)R_\\odot$. The radiative opacity near the base of the CZ that is obtained from OPAL tables must be increased by $\\sim 21$\\% in the Bahcall-Pinsonneault (2004) solar model if one wants to invoke opacity errors in order to reconcile recent solar heavy abundance determinations with the helioseismological measurement of $R_{\\rm CZ}$. We show that the radiative opacity near the base of the convective zone depends sensitively upon the assumed heavy element mass fraction, $Z$. The uncertainty in the measured value of $Z$ is currently the limiting factor in our ability to calculate the depth of the CZ. Different state-of-the-art interpolation schemes using the existing OPAL tables yield opacity values that differ by $\\sim 4$\\% . We describe the finer grid spacings that are necessary to interpolate the radiative opacity to $\\pm 1$\\%. Uncertainties due to the equation of state do not significantly affect the calculated depth of the convective zone. ",
        "introduction": "\\label{sec:intro} The depth of the solar convective zone has been measured by helioseismological techniques to high accuracy.  In the most comprehensive study to date, Basu \\& Antia (1997)\\nocite{basucz} have investigated the influence of  observational and theoretical systematic uncertainties as well as measurement errors. Basu and Antia concluded that the base of the solar convective zone currently lies at a depth of \\begin{equation} R_{\\rm CZ} ~=~ (0.713 \\pm 0.001)R_\\odot \\, . \\label{eq:rczmeasured} \\end{equation} The result of Basu and Antia is consistent with the earlier measurements of Kosovichev \\& Fedorova (1991),\\nocite{kosovichevcz} who obtained $R_{\\rm CZ} ~=~ (0.713 \\pm 0.001)R_\\odot$, and Christensen-Dalsgaard, Gough, \\& Thompson (1991),\\nocite{jcdCZ} who also obtained $R_{\\rm CZ} ~=~ (0.713 \\pm 0.003R_\\odot$, as well as with the determination of Guzik \\& Cox (1993),\\nocite{guzikcz} who found $R_{\\rm CZ} ~=~ (0.712 \\pm 0.001)R_\\odot$.  Basu (1998)\\nocite{basu98} also studied the effect of the assumed value of the solar radius on the inferred depth of the convective zone and found $R_{\\rm CZ} ~=~ (0.7135 \\pm 0.0005)R_\\odot$. The analyses in these different studies span a wide range of reference solar models and analysis techniques.  On the basis of the analyses cited above, the measurement of the depth of the solar convective zone appears robust and precise.  Recently, new precision measurements have been made of the C, N, O, Ne, and Ar abundances on the surface of the Sun (Allende Prieto, Lambert, \\& Asplund 2001;\\nocite{allende01} Allende Prieto, Lambert, \\& Asplund 2002;\\nocite{allende02} Asplund et al. 2004;\\nocite{asplund04} Asplund et al. 2000;\\nocite{asplundetal00} Asplund 2000)\\nocite{asplund00}.  These new abundance determinations use three-dimensional rather than one-dimensional atmospheric models, include hydrodynamical effects, and pay particular attention to uncertainties in atomic data and the observational spectra. The new abundance estimates, together with the previous best-estimates for other solar surface abundances~(Grevesse \\& Sauval 1998)\\nocite{oldcomp}, imply $Z/X = 0.0176$, much less than the previous value of $Z/X = 0.0229$ (Grevesse \\& Sauval 1998)\\nocite{oldcomp}.  For a solar model with the recently-determined heavy element to hydrogen ratio, the calculated depth of the convective zone is~(Bahcall \\& Pinsonneault 2004)\\nocite{BP04} \\begin{equation} R_{\\rm CZ}(Z/X = 0.0176) ~=~ 0.726 R_\\odot \\, , \\label{eq:rcznewzoverx} \\end{equation} which is very different from the measured depth of the CZ (see equation~[\\ref{eq:rczmeasured}]). On the other hand, Basu and Antia (2004)~\\nocite{basu04} have shown that the helioseismological determination of $R_{\\rm CZ}$, equation \\ref{eq:rczmeasured},  is not affected if one assumes the correctness of the lower heavy element abundances ($Z/X = 0.0176$).  Something is wrong. We have a new solar problem: ``the convective zone (CZ) problem.''  The radiative opacity is a key ingredient in calculating the depth of the convective zone. Moreover, about 95\\% of the total radiative opacity near the base of the convective zone involves bound electrons, either bound-free or bound-bound opacity (Iglesias 2004).\\nocite{iglesias04} Thus opacity calculations in this region involve details of the ionization balance and other delicate atomic physics properties.  In this paper, we focus on determining the accuracy with which the opacity near the base of the convective zone must be known in order to calculate precisely the depth of the CZ with a stellar evolution code.  We also evaluate the accuracy with which the opacity near the base of the CZ can be interpolated from OPAL tables. For a related comparison of the Los Alamos LEDCOP opacities and the OPAL opacities, see Neuforge-Verheecke et al. (2001)~\\nocite{losalamos}. For comprehensive discussions of stellar radiative opacities, the reader is referred to the important reviews by Rogers and Iglesias (1998)\\nocite{OPALreview} and by Seaton et al. (1995)\\nocite{opbook}.  We investigate in a paper in preparation (Bahcall, Basu, Pinsonneault, and Serenelli 2004)~\\nocite{inpreparation} the helioseismological implications of the changes in opacity that are discussed in the present paper. The viability of any proposed change in the opacity discussed in the present paper must be tested by comparing a solar model that is evolved with the changed opacity with a complete set of precise helioseismological data. There is no compelling reason to believe that the illustrative change in opacity considered here, which is highly peaked in radius,  will be either reproduced exactly by new opacity calculations or will be precisely consistent with helioseismological constraints. In the future, once new opacity calculations are available that satisfy the requirements described in this paper, it will be possible to test simultaneously the new opacities, the solar model evolution, and the helioseismological implications.  We derive in \\S~\\ref{sec:dependence} the dependence, $\\partial \\ln R_{\\rm CZ}/\\partial \\ln \\kappa$, of the calculated depth of the solar convective zone upon the assumed radiative opacity. We apply this result to determine the accuracy with which the opacity must be known in order to calculate the depth of the CZ to the accuracy with which it is measured helioseismologically.  We also determine the change in the standard OPAL opacity that is required in order to reconcile the helioseismological measurement with the recent determinations of heavy element abundances. We evaluate in \\S~\\ref{sec:composition} the dependence of the radiative opacity near the base of the convective zone upon the stellar composition. We find that the opacity depends sensitively upon the assumed heavy element abundance. We compare in \\S~\\ref{sec:comparison} the opacities obtained from two different interpolation schemes that are both applied to the same published OPAL opacity tables. Throughout this paper, we adopt the OPAL opacities (Iglesias \\& Rogers 1991a,b; Rogers \\& Iglesias 1992; Iglesias \\& Rogers 1996)\\nocite{iglesias91a,iglesias91b,rogers92,opalopacity96} as standard, when supplemented at low temperatures by values from Alexander \\& Fergusson (1994)\\nocite{alexanderopac}. We use simulated opacity tables in \\S~\\ref{sec:shifted} to estimate the likely uncertainties that result from interpolations within the existing OPAL opacity tables and to determine the grid sizes to obtain small interpolation errors. For completeness and for contrast, we use four different equations of state  to show in Appendix~A that uncertainties due to the choice of EOS are not important, at the present level of accuracy, for the calculation of the depth of the solar convective zone. We also demonstrate in Appendix~B that uncertainties in the nuclear reaction rates affect the depth of the solar convective zone only at the level of 0.1\\% . In Appendix~C, we verify that the conversion of carbon and oxygen in CNO burning, which cannot be accurately modeled with existing opacity codes, causes a 0.1\\%  uncertainty in the calculated depth of the convective zone. Basu and Antia (2004) (see also Asplund et al. 2004)  have shown that errors in the calculation of the diffusion coefficients are unlikely to be the correct explanation of the discrepancy between measured and calculated depth of the solar convective zone. Other solar model ingredients, including the element diffusion coefficients, can affect the calculated depth of the convective zone. A complete investigation of all the possible effects on the convective zone is beyond the scope of the present paper and would distract the reader from our main concern, the effect of the radiative opacity. Moreover, we believe that the radiative opacity and the heavy element abundance provide the single largest contributions to the error budget for the calculation of the solar convective zone. The effect of the heavy element abundance on the calculated depth of the convective zone has been evaluated in Bahcall and Pinsonneault (2004). We summarize and discuss our main results \\hbox{in \\S~\\ref{sec:summary}}. ",
        "conclusions": "\\label{sec:summary}  The primary goal of this paper is to determine how accurately the radiative opacity near the base of the convective zone must be known in order to use measurements of the CZ depth to draw conclusions about other solar parameters. There are two separate but related issues with respect to the accuracy of the radiative opacity, namely, the accuracy with which the tabulated values in opacity tables are calculated and the accuracy with which the opacity can be interpolated within tables of a specified grid size. We first summarize our conclusions regarding the accuracy of tabulated opacity values and then we summarize our results with respect to the accuracy of interpolations within the standard OPAL opacity tables. The helioseismological implications of the opacity changes considered in this paper will be discussed in Bahcall, Basu, Pinsonneault, and Serenelli, (2004, in preparation).  We show in \\S~\\ref{sec:dependence} that the logarithmic derivative of the convective zone depth with respect to the logarithm of the opacity satisfies $\\partial \\ln R_{\\rm CZ}/\\partial \\ln \\kappa \\approx -0.095$. We conclude from this relation that the radiative opacity must be known to an accuracy of 1\\% in order to calculate in a solar model the depth of the CZ to the accuracy, 0.14\\%, with which the depth is measured by helioseismology. On the other hand, if one accepts the recent measurements of heavy element abundances, then the OPAL opacities must be increased by about 21\\% in order to reconcile the calculated solar model depth of the CZ and the measured depth of the CZ. This change of 21\\% could conceivably arise from a combination of errors in the tabulated values of the opacity and interpolation errors, which are discussed below. However, as we shall see, the total change of 21\\% is too large to be ascribed solely to errors in interpolation.  It would be very instructive to have a comprehensive study of the absolute accuracy of state-of-the-art radiative opacity calculations. A detailed comparison of the calculated opacity near the base of the convective zone obtained by the Opacity Project (Seaton, Yan, Mihalis, \\& Pradhan 1994)~\\nocite{opacityproject} with the results of the OPAL project (Iglesias \\& Rogers 1996) would be very informative. The interested reader is referred to the informative and insightful comparison by Neuforge-Verheecke et al. (2001) of the Los Alamos LEDCOP opacities and the OPAL opacities. The largest differences are found near the base of the convective zone, with the OPAL opacities being as much as 6\\% larger than the LEDCOP opacities in this region. As part of a comprehensive discussion of factors that affect the accuracy of solar models, Boothroyd \\& Sackmann (2003)~\\nocite{sackmann03} have investigated ways that the opacities can affect helioseismological parameters.  We show in \\S~\\ref{sec:composition} that the radiative opacity near the base of the convective zone depends sensitively upon the assumed chemical composition (see especially equation~\\ref{eq:kappaoncomposition} and equation~\\ref{eq:compositionderivatives}). If one wanted to calculate the depth to an accuracy of $0.6$\\%, then one would need to know the heavy element mass fraction, $Z$, to an accuracy of 1\\%. This precision is far beyond the current state-of-the-art accuracy in the determination of the heavy element abundance.  The entire difference between the measured depth of the solar convective zone (equation~\\ref{eq:rczmeasured}) and the value calculated  using a solar model with the recent low determinations of the heavy element abundances (equation~\\ref{eq:rcznewzoverx}) could be explained by the present uncertainty, $\\sim 15$\\%, in the ratio of Z/X (see Bahcall \\& Pinsonneault 2004)\\nocite{BP04}.  Of course, the changes in opacity caused by changing $Z/X$ are not limited to any particular region. Changing the assumed surface value of $Z/X$  affects the composition and hence the opacity throughout the solar model.  We have approximated in this paper the dependence of the opacity upon composition by the dependence upon just two variables,  the mass fractions $X$ and $Z$.  In reality, the situation is more complex. Different chemical elements contribute differently to the stellar opacity. For example, Bahcall, Pinsonneault, and Basu (2001) found that the depth of the convective zone was most sensitive to the abundances of the lighter metals, which are significant opacity sources at $2 \\times 10^6$K, while the heavier metals were much more important for the core structure and the estimated initial solar helium abundance.  However, we are not yet at a level of precision that we can specify well the opacity-weighted uncertainties of the different heavy elements. This is a refinement that will have to await further progress in determining the different heavy element abundances and more extensive opacity calculations.   We compare in \\S~\\ref{sec:comparison} the radiative opacity values obtained with two different interpolation routines from the standard OPAL opacity tables. We find that the difference in interpolated values of the radiative opacity can be as large as 4\\% near the base of the convective zone. We also tested in \\S~\\ref{sec:shifted} the accuracy with which interpolations can be performed within simulated opacity tables of different grid sizes. We find that errors of the order of 3\\% may be expected from tables with the grid spacings of the existing OPAL tables. However, we show that the interpolation uncertainties could be reduced to the level of 1\\% or below by using a denser grid with $\\Delta \\log T = 0.025$, $\\Delta \\log r = 0.125$, and with $Z$ ranging from $Z = 0.0100$ to $Z = 0.0225$ with $\\Delta Z = 0.0025$.  For completeness, we report in the Appendix on the calculated depth of the CZ that was found using  four different equations of state. In agreement with other authors, we find that the choice of equation of state affects the calculated depth of the CZ by only about $\\pm 0.1$ \\%. We also show in the Appendix that current uncertainties in nuclear reaction rates also affect the calculated depth of the convective zone at the level of 0.1\\%."
    },
    "0403/astro-ph0403097_arXiv.txt": {
        "abstract": "Milagro is a water Cherenkov extensive air shower array that continuously monitors the entire overhead sky in the TeV energy band.  The results from an analysis of $\\sim$3 years of data (December 2000 through November 2003) are presented. The data has been searched for steady point sources of TeV gamma rays between declinations of 1.1 degrees and 80 degrees. Two sources are detected, the Crab Nebula and the active galaxy Mrk 421. For the remainder of the Northern hemisphere we set 95\\% C.L. upper limits between 275 and 600 mCrab (4.8-10.5$\\times 10^{-12}$ cm$^{-2}$ s$^{-1}$) above 1 TeV for source declinations between 5 degrees and 70 degrees. Since the sensitivity of Milagro depends upon the spectrum of the source at the top of the atmosphere, the dependence of the limits on the spectrum of a candidate source is presented.   Because high-energy gamma rays from extragalactic sources are absorbed by interactions with the extragalactic background light the dependence of the flux limits on the redshift of a candidate source are given. The upper limits presented here are over an order of magnitude more stringent than previously published limits from TeV gamma-ray all-sky surveys. ",
        "introduction": "\\label{sec:Introduction}  Sources of very-high-energy (VHE, $>$100 GeV) gamma rays are observed to be non-thermal in nature and are typically the sites of particle acceleration.  This acceleration is thought to occur in astrophysical shocks such as those believed to exist in plerions \\citep{deJager1992}, supernova remnants \\citep{Volk2000}, active galactic nuclei \\citep{Blandford1978}, and galaxy clusters \\citep{Loeb2000} (among other sources).  These shocks may accelerate protons or electrons, both of which lead to the emission of gamma rays.  Since gamma rays are unaffected by the magnetic fields that pervade the Galaxy and the Universe, they can be used to pinpoint the sites of particle acceleration. In addition to these ``classical'' astrophysical sources of VHE gamma rays, other more exotic objects such as primordial black holes, topological defects, and the decay of relic particles from the big bang may also emit VHE gamma rays.  Perhaps of most interest is the possible existence of a new type of source that has yet to be postulated.  A comprehensive survey of the sky sensitive to emission at all time scales is necessary to detect the many possible sources. The analysis presented in this paper is part of an ongoing effort by the Milagro collaboration to search the entire northern hemisphere for such objects.  The search for short bursts of TeV gamma rays has been addressed in a previous paper \\citep{Atkins2004}.  The analysis presented here deals specifically with steady point sources of TeV gamma rays.  It has been over a decade since the discovery of the first source of VHE gamma rays, the Crab Nebula \\citep{Weekes89}.  Since that time there have been seven other confirmed sources of TeV gamma rays \\citep{Horan2003}, six of which lie in the northern hemisphere.  With the exception of the Crab Nebula and the supernova remnant PKS 1706-44, these objects are all active galactic nuclei of the blazar class \\citep{Horan2003}.   All of these objects have been discovered by atmospheric Cherenkov telescopes searching for counterparts to sources discovered at lower energies. In contrast, the EGRET instrument detected over 270 objects emitting high-energy gamma rays above 100 MeV \\citep{Hartman99}. One hundred and seventy of these objects are not identified at other wavelengths.  The VHE regime is a natural energy band to search for counterparts to these objects.  The small field-of-view and low duty factor of the atmospheric Cherenkov telescopes (ACTs) make comprehensive sky surveys difficult to perform. As a result, very few comprehensive surveys of the VHE sky have been performed to date. The first VHE survey was performed by a non-imaging ACT \\citep{Helmken79, Weekes79} in 1979.  The Milagrito instrument (a prototype to Milagro, with no background rejection capability and a higher energy threshold) also performed a survey of the northern hemisphere \\citep{Atkins2001} and set limits of $\\le$3 Crab from any point source in the northern hemisphere.  More recently flux limits of 4-9 times that of the Crab Nebula (above 15 TeV) have been published by the AIROBICC collaboration \\citep{Aharonian2002}.  (1 Crab is equivalent to an integral flux above 1 TeV of $F(>$1 TeV)$=1.75 \\times$10$^{-11}$ cm$^{-2}$ s$^{-1}$.) The limits presented here are over an order of magnitude more stringent than these previous surveys.  Since many of the confirmed sources of VHE gamma rays are extragalactic the limits must account for the absorption of TeV gamma rays by interactions with the extragalactic background light (EBL) \\citep{Primack2000, Stecker2002, Kneiske2002}. The EBL is comprised of visible radiation emitted by stars and infrared radiation emitted by dust due to reprocessed starlight. Since direct measurements of the intensity and spectrum of the EBL are problematic due to the foreground light from our galaxy, a model is used to determine the effect of the EBL on the observed spectrum at earth from a distant source.  Before employing the results presented here the limitations of this survey need to be understood.  First, the limits presented only apply to point sources, not to extended objects, such as the galactic plane.  Second, they only apply to the average VHE emission during the time period over which the data was obtained. ",
        "conclusions": "\\label{sec:Conclusions} A complete survey of the northern hemisphere (declination $1.1^{\\circ}$ to $80^{\\circ}$) for point sources of TeV gamma rays has been performed.  These limits apply to the average flux level during the roughly 3 year period from December 15, 2000 through November 25, 2003.  The average 95\\% C.L. upper limits range from 275 mCrab to 600 mCrab depending upon the declination of the source and are over an order of magnitude more restrictive than previous limits. A prescription has been given to calculate the corresponding upper limits for sources with different spectra and for extragalactic sources.  For sources with differential spectral indices of -2.0 the upper limits are 57\\% lower.  For a source at a redshift of 0.03 the flux limits are a factor of 2.4 larger.  While these limits are the best available to date, Milagro has recently been completed with the construction of an array of 175 water tanks surrounding the central reservoir.  A comparable dataset, with this now complete Milagro detector, would improve these limits by a factor $\\sim$2."
    },
    "0403/astro-ph0403574_arXiv.txt": {
        "abstract": " ",
        "introduction": "The current paradigm for the formation of giant gaseous planets is based on the so-called core accretion model in which a growing solid core reaches a critical mass and accretes rapidly a massive atmosphere (Pollack et al. \\cite{P96}, hereafter referred to as P96). While this model has many appealing features, it suffers at least from two shortcomings which, as we shall show later, are actually coupled.  First, the timescale (close to 10 Myr) found by P96 to form Jupiter at its present location is uncomfortably close to the typical lifetime of protoplanetary discs which is believed to be of the order of 1-10 Myr (Haisch et al. \\cite{Haisch}). This timescale problem has led others to look for more rapid formation mechanisms based on direct gravitational collapse (Boss \\cite{boss02}). Second, P96 assumed that the giant planets of our solar system have been formed where they are observed today. However, the discovery over the last decade of extrasolar planets at very short distances to their parent star, has open the possibility that planets may actually migrate over large distances (Lin et al.  \\cite{linetal96}; Trilling et al. \\cite{Trilling}; Papaloizou \\& Terquem  \\cite{PT99}, hereafter referred to as PT99). The timescale of migration is still very uncertain, but conservative estimates give values between  $0.1$ and $10$ Myr. This timescale is therefore comparable to the planet formation timescale and the disc lifetime and thus migration cannot be neglected in a self-consistent picture.  In this Letter, we extend the core accretion model of P96 to include migration, disc evolution and gap formation and show that the formation timescales of giant planets can be reduced by factors of ten or more. ",
        "conclusions": "We have calculated the formation of giant planets up to runaway gas accretion and studied the effect of migration and gap formation on the resulting formation timescale. Our main result is that phase 2 described in P96 is suppressed, and this leads to a formation time around $1$ Myr for the initial conditions chosen here. This is roughly an order of magnitude faster that the favored model of P96 even though our initial disc had less mass than the one used by P96.  We have performed many tests and convinced ourselves that this speed-up is robust against changes in various parameters. For example, in a calculation in which the reduction of type I migration ($f_I$) is set to 0.1, an embryo starting at $15$AU undergoes runaway accretion in less than 3 Myr. The assumed size of the planetesimals plays a critical role as already noted by P96. For example, assuming planetesimals of 10km instead of 100km leads to runaway accretion after only 0.3 Myr! As already mentioned in section \\ref{infalling}, since the effect of ablation is negligible in our calculations, mass loss of planetesimals occurs very deep in the envelope. The switch from sinking to no-sinking approximation (see P96) has then a small effect. In the migrating case (without gap formation) we obtain a formation timescale of $\\sim 1.2$ Myr using the sinking approximation, compared to $\\sim 1$ Myr in the no-sinking case.  Thus, migration, besides explaining the presence of giant planets at short distances to their stars, also plays an important role in the formation process itself. By ensuring that the feeding zone is never depleted, migration suppresses almost entirely the protracted phase 2 therefore shortening enormously the formation time. This of course,  does not preclude other effects such as reduction of opacity or formation of vortices prior to planet formation to further reduce this timescale. The formation of giant planets through the core-accretion scenario can therefore proceed over timescales in good agreement with disk lifetimes, and {\\it without having to consider disks significantly more massive than the minimum mass solar nebula}.  Finally, we point out that the formation of giant planets appears only possible if the currently available type I migration rates are reduced by at least a factor of 10. This suggests that their might still be a serious problem in our understanding of this type of migration.  The authors would like to acknowledge the help of C. Winisdoerffer and the ENS-Lyon team of the C.R.A.L. We thank P. Bodenheimer and G. Meynet for useful discussions. This work was supported in part by the Swiss National Science Foundation."
    },
    "0403/astro-ph0403268_arXiv.txt": {
        "abstract": "Current data exclude cosmic strings as the primary source of primordial density fluctuations. However, in a wide class of inflationary models, strings can form at later stages of inflation and have potentially detectable observational signatures. We study the constraints from WMAP and SDSS data on the fraction of primordial fluctuations sourced by local cosmic strings. The Bayesian analysis presented in this brief report is restricted to the minimal number of parameters. Yet it is useful for two reasons. It confirms the results of \\cite{PTWW03} using an alternative statistical method. Secondly, it justifies the more costly multi-parameter analysis. Already, varying only three parameters -- the spectral index and the amplitudes of the adiabatic and string contributions -- we find that the upper bound on the cosmic string contribution is of order $10\\%$. We expect that the full multi-parameter study, currently underway \\cite{PWW04}, will likely loosen this bound. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403560_arXiv.txt": {
        "abstract": "We construct eigenspectra from the DR1 quasars in the SDSS using the Karhunen-Lo\\`eve (KL) transform (or Principal Component Analysis, PCA) in different redshift and luminosity bins. We find that the quasar spectra can be classified, by the first two eigenspectra, into a continuous sequence in the variation of the spectral slope. We also find a dependence on redshift and luminosity in the eigencoefficients. The dominant redshift effect is the evolution of the blended Fe~II emission (optical) and the Balmer continuum (the ``small-bump'', $\\lambda_{rest} \\approx 2000-4000$\\AA), while the luminosity effect is related to the Baldwin effect. Correlations among several major broad emission lines are found, including the well-known ``Eigenvector-1''. ",
        "introduction": "Spectroscopic observations from the Sloan Digital Sky Survey (SDSS) have the advantage of large numbers of quasars (QSOs) and a large redshift range that provides an unique opportunity to study their intrinsic properties, sample variation, and spectral classification in great detail. A powerful method to address these questions is the KL transform. The idea is to derive from the observed spectra a lower dimensional set of eigenspectra (Connolly et al.\\ 1995), in which the essential spectral properties are represented. In the following we outline some results from applying this method to the SDSS quasars. ",
        "conclusions": ""
    },
    "0403/astro-ph0403083_arXiv.txt": {
        "abstract": "Several short-period Jupiter-mass planets have been discovered around nearby solar-type stars.  During the circularization of their orbits, the dissipation of tidal disturbance by their host stars heats the interior and inflates the sizes of these planets. Based on a series of internal structure calculations for giant planets, we examine the physical processes which determine their luminosity-radius relation. In the gaseous envelope of these planets, efficient convection enforces a nearly adiabatic stratification. During their gravitational contraction, the planets' radii are determined, through the condition of a quasi-hydrostatic equilibrium, by their central pressure.  In interiors of mature, compact, distant planets, such as Jupiter, degeneracy pressure and  the non-ideal equation of state determine their structure. But, in order for young or intensely heated gas giant planets to attain quasi-hydrostatic equilibria, with sizes comparable to or larger than two Jupiter radii, their interiors must have sufficiently high temperature and low density such that degeneracy effects are relatively weak. Consequently, the effective polytropic index monotonically increases whereas the central temperature increases and then decreases with the planets' size.  These effects, along with a temperature-sensitive opacity for the radiative surface layers of giant planets, cause the power index of the luminosity's dependence on radius  to decrease with increasing radius.  For planets larger than twice  Jupiter's radius, this index is sufficiently small that they become unstable to tidal inflation.  We make comparisons between cases of uniform heating and cases in which the heating is concentrated in various locations within the giant planet.  Based on these results we suggest that accurate measurement of the sizes of close-in young Jupiters can be used to probe their internal structure under the influence of tidal heating. ",
        "introduction": "One of the surprising findings in the search for planetary systems around other stars is the discovery of extrasolar planets with periods down to 3 days (\\cite{mq95}).  Nearly all planets with period less than 7 days have nearly circular orbits.  In contrast, known extrasolar planets with periods longer than 2-3 weeks, have nearly a uniform eccentricity distribution.   The shortest-period planets and their host stars induce tidal perturbations  on each other.  When these disturbances are dissipated, angular momentum is exchanged between the planets and their host stars, leading toward both a spin synchronization and orbital circularization (\\cite{ra96}).  \\cite{BLM} (hereafter Paper I), considered the effect of tidal dissipation of energy during the synchronization of these planets' spin and the circularization of their orbits.  In that analysis, they compute a series of numerical models for the interior structures of weakly eccentric Jovian planets at constant orbital distances under the influence of interior tidal heating and stellar irradiation. In these previous calculations, the interior heating rate per unit mass was imposed to be constant in time and uniformly distributed within the planet.  Under these assumptions, they showed that Jovian planets can be inflated to equilibrium sizes considerably larger than those deduced for gravitationally contracting and externally heated planets. For the transiting planet around HD 209458, they suggested that provided its dimensionless dissipation $Q$-value is comparable to that inferred for Jupiter (\\cite{yp81}), a small eccentricity ($e \\simeq 0.03$) would provide adequate tidal heating to inflate it to its observed size (\\cite{brown01}).  Since the orbital circulation time scale is expected to be shorter than the life span of the planet, they also suggested that this eccentricity may be excited by another planet with a longer orbital period.  The prediction of a small eccentricity and the existence of another planet are consistent with existing data (\\cite{bll03}).  There are at least two other scenarios for the unexpectedly large size of HD 209458b.  Heating by stellar irradiation reduces the temperature gradient and the radiative flux in the outer layers of short-period planets.  This process could significantly slow down the Kelvin-Helmholtz contraction of the planet and explain the large size (\\cite{bur00}). However, even though the stellar flux onto the planet's surface is 5 orders of magnitude larger than that released by the gravitational contraction and cooling of its envelope, this heating effect alone increases the radius of the planet by about 10\\%, not by 40\\% as observed (\\cite{gs02}).  An alternative source is the kinetic heating induced by the dissipation of the gas flow in the atmosphere which occurs because of the pressure gradient between the day and night sides (\\cite{gs02}). In order to account for the observed size of the planet, conversion of only 1\\% of the incident radiative flux may be needed, provided that the dissipation of induced kinetic energy into heat occurs at sufficiently deep layers (tens to 100 bars).  \\cite{sg02} suggest that the Coriolis force associated with a synchronously spinning planet may induce the circulation to penetrate that far into the planet's interior, and that dissipation could occur through, for example, Kelvin-Helmholtz instability.  A follow-up analysis suggests that this effect may be limited (\\cite{blb03}; \\cite{jl03}).  In order to distinguish between these three scenarios, the effect of tidal heating for planets with modest to large eccentricities was considered.  In a follow-up paper (Gu {\\it et al.} 2003, hereafter Paper II), we showed that the size ($R_p$) of compact Jupiter-mass planets slowly increases with the tidal dissipation rate.  For computational simplicity, we adopted a conventional equilibrium tidal model which describes the planets' continuous structural adjustment in order for them to maintain a state of quasi-hydrostatic equilibrium in the varying gravitational potential of their orbital companion.  In this prescription, a phase lag into the response is introduced to represent the lag being proportional to the tidal forcing frequency and attributable to the viscosity of the body.  The phase lag gives rise to a net tidal torque and dissipation of energy and the efficiency of the tidal dissipation can be parametrized, whatever its origin, by a specific dissipation function or $Q$-value (quality factor) (\\cite{gs66}).  External perturbation can also induce dynamical tidal responses through the excitation of g modes (Ioannou \\& Lindzen 1993a,b) or inertial waves (\\cite{ol04}) which can be damped by viscous dissipation in the interior (\\cite{gn77}) or radiative and nonlinear dissipation in the atmosphere of the planet (\\cite{ltl97}).  In both the prescription for equilibrium and dynamical responses, the tidal dissipation rate is a rapidly increasing function of the planets' radius.  But, their surface luminosity increases even faster with $R_p$ such that, planets with relatively small eccentricities and modest to long periods attain a state of thermal equilibrium in which the radiative loss on their surface is balanced by the tidal dissipation in their interior.  For planets with short periods and modest to large eccentricities, the rate of interior heating is sufficiently large that their $R_p$ may inflate to more than two Jupiter radii.  In this limit, the surface luminosity of the planet becomes a less sensitive function of its $R_p$ and the eccentricity damping rate is smaller than the expansion rate of the planet so that the increases in their surface cooling rate cannot keep pace with the enhanced dissipation rate due to their inflated sizes.  These planets are expected to undergo runaway inflation and mass loss.  We suggested that the absence of ultra-short-period Jupiter-mass planets with $P<3$ days, which corresponds to an orbital semi-major axis $a$ of 0.04 AU, may be due to mass loss through Roche-lobe overflow resulting from such a tidal inflation instability. Other scenarios have been proposed to explain the lack of Jupiter-mass planets with $P<3$ days: truncation of inner part of a disk (\\cite{lin96}; \\cite{kl02}), orbital migration due to the spin-orbit tidal interaction between the close-in planet and the parent star (\\cite{ra96}; \\cite{ws02}; \\cite{pr02}; \\cite{jiang03}), and Roche-overflowing planets with the help of disk-planet interaction but without tidal inflation (\\cite{tri98}).  In this contribution, we continue our investigation on the internal structure of tidally heated short-period planets.  The main issues to be examined here are: 1) how does tidal energy dissipation actually lead to the expansion of the envelope? 2) how does the internal structure of the planet depend on the distribution of their internal tidal dissipation rate? and 3) what are the important physical effects which determine the tidal inflation stability of the planets?  These issues are important in determining the mass-radius relation of short-period planets which is directly observable.  Structural adjustments may also modify the efficiency of dissipation and the planets' Q-value for both equilibrium and dynamical tides.  In the extended convective envelopes of gaseous giant planets and low-mass stars, turbulence can lead to dissipation of the motion that results from the continual adjustment of the equilibrium tide. However, the turbulent viscosity estimated from the mixing-length theory ought to be reduced by a frequency-dependent factor owing to the fact that the convective turnover time scale is usually much longer than the period of the tidal forcing. Based on the present-day structure of Jupiter, Goldreich \\& Nicholson (1977) estimated $Q\\approx5\\times10^{13}$. However, within the intensely heated (by tidal dissipation) short-period extra solar planets, convection is expected to be more rigorous with higher frequencies whereas their tidal forcing frequencies are smaller than that of Jupiter.  The Q-value for the equilibrium tide within extra solar planets is likely to much smaller than that within Jupiter.  For the dynamical response of short-period extra solar planets planet, the forcing frequencies are typically comparable to their spin frequencies but are small compared to their dynamical frequencies. Convective regions of the planets support inertial waves, which possess a dense or continuous frequency spectrum in the absence of viscosity, while any radiative regions support generalized Hough waves (Ioannou \\& Lindzen 1993a, b).  Inertial waves provide a natural avenue for efficient tidal dissipation in most cases of interest.  The resulting value of $Q$ depends, in principle, in a highly erratic way on the forcing frequency. Since the planets' spin frequency adjusts with their $R_p$, which is a time-delayed function of the tidal dissipation rate within them, the efficiency of tidal dissipation may fluctuate while the overall evolution is determined by a frequency-average Q-value (\\cite{ol04}).   In \\S2, we briefly recapitulate the basic equations which determine the quasi-static evolution of the planets' structure.  We show the simulation results for inflated giant planets in the case of constant internal heating per unit mass in \\S3 and analyze the results, in the Appendix, in terms of polytropic models which allow us to conclude that the onset of the tidal runaway inflation instability is regulated by a transition in the equation of state for the interior gas from a partially degenerate/non-ideal state toward a more ideal-gas state. We examine the dependence of planetary adjustment on different locations of the energy dissipation in \\S4.  Finally, we summarize the results and discuss their observational implications in \\S5. ",
        "conclusions": "In this paper, we continue our investigation on the adjustment of a planetary interior as a consequence of intense tidal heating.  As a giant planet's interior is heated and inflated, we showed in Paper II that its interior remains mostly convective.  Efficient energy transport leads to an adiabatic stratification.  With a constant heating rate per unit mass, we deduced an unique luminosity-radius relation regardless of how intense the heating rate is.  The planet's luminosity increases with its radius.  But the growth rate of $\\cal L$ is a decreasing function of $R_p$.  At the same time, the tidal dissipation heats the interior of the planet at a rate which  increases rapidly with $R_p$. At around $2~R_J$, $\\cal L$ can no longer sustain sufficient growth to maintain a thermal equilibrium with the tidal dissipation rate. Thereafter, the planet's inflation become unstable and it overflows its Roche radius and become tidally disrupted.  Here we show that the change of luminosity during the planet's expansion is directly linked to the evolution of its interior, in particular, the equation of state.  We employ the polytrope approach to investigate the interior structure of an inflated giant planet. According to the simulation, interior profiles deviate away from $P\\propto \\rho^2$ as the planet expands. The central temperature $T_c$ increases and then decreases with the size of the planet. Also $P_c$ and $\\rho_c$ are less steep functions of $R_p$ than the polytrope theory with a constant polytropic  index $n$  indicates. All of these effects suggest that the planet interior does not evolve in a self-similar manner,  but $n$ gets larger  as $R_p$ increases.  In conjunction with the numerical results that the degeneracy $D$ decreases, and that $T_c$ rises and then drops during the course of inflation, the result of a positive value of $dn/dR_p$ can be interpreted as a manifestation of a reduction in degeneracy during the expansion.  We reason that the coefficient of thermal expansion $\\delta$ increases in response to a decrease in degeneracy and non-ideal effects, leading to an increase in $n$ through equation (\\ref{eq:n_delta}). Consequently the planet compressibility at constant mass $C_M$ increases with $R_p$ (see equation \\ref{eq:C_M}). This pattern can be translated into the phenomenon of a decrease in luminosity growth during the inflation, as a consequence of an one-to-one relation between $K$ and $\\cal L$ in the case of uniform heating in mass under the conditions of the polytropic interior and hydrostatic equilibrium. We also compare the results between a planet with a core and without a core.  To be inflated to the same size, a planet with a core, therefore possessing a larger gravitational binding energy, needs a larger intrinsic luminosity $\\cal L$ than a planet with no core and the same mass.  We also show that the opacity in the radiative envelope has a drastic effect on the final equilibrium size of an inflated planet: the size would be much smaller if grains are depleted in the radiative envelope.  We also consider the possibility of localized tidal dissipation.  Such a process may occur in differentially rotating planets or near the interface between the convective and radiative zones where the wavelength associated with dynamical tidal response is comparable to the density scale height. Localized dissipation may also occur through the dissipation of resonant inertial waves or radiative damping in the atmosphere.  In the strong shell-heating models, the one-to-one relation between $K$ and $\\cal L$ disappears because of the existence of a radiative region caused by temperature inversion beneath the shell-heating zone. The unheated planet's interior in such cases might still be inflated due to a significant expansion of the gas above the heating zone, although the overall expansion rate is less efficient than that in the uniform heating cases as a result of a greater amount of radiative loss from the planet's photosphere. Without gaining entropy, the expanding interior cannot lift its degeneracy and therefore cannot increase its elasticity. However, the gas above the shell-heating zone can lift its non-ideal properties and hence enhance its elasticity, leading to a decrease in $\\Gamma$ in that region and thereby diminishing the growth rate of $\\cal L$ as the planet expands.  Finally, we consider the self consistent response, taking into account the modification of heating rate due to a planet's expansion.  In this paper, we adopt a constant-$Q$ prescription for equilibrium tides in which the tidal dissipation rate is assumed to be proportional to $R_p^5$.  The results for the uniform heating model suggest that a young gaseous planet of $1~M_J$ without a solid core can be thermally unstable and inflated from $2~R_J$ to a size beyond $4~R_J$ if $e\\sim 0.2$ at $a=10~\\rsun$. If the dissipation rate  is proportional to $R_p^5$,  and if most of the tidal perturbation is deposited at $m/M_p=0.9$, a core-less young planet of $1 M_J$ would be thermally unstable and inflated from $2 R_J$ to a size beyond $3 R_J$ for $e>0.07$ at $a=0.03$ AU or $e>0.26$ at $a=10 \\rsun$.  With the same heating concentration $m/M_p=0.9$ and $R_p$-dependence in the dissipation rate, a young planet with a core at $a=0.03$ AU with an initial eccentricity $e>0.294$ can be inflated from $2~R_J$ to a size beyond its Roche radius.  We have assumed that the convective flow still behaves adiabatically even though the heating shell causes a narrow radiative zone. However, the condition away from adiabaticity implies that the internal heat is not transported away as efficiently as in the case of  adiabatic convection, leading to a more severe reduction in non-ideal properties of  the gas and therefore an even faster decrease of $\\gamma$ as the planet expands.  Note that the Eddington approximation for the surface boundary condition  is used in these models rather than more detailed frequency-dependent model atmospheres. This approximation is not necessarily valid for the strongly irradiated atmospheres studied here (\\cite{gs02}). However it is unlikely to make much difference for the main results discussed here, namely the behavior of the planet's radius as a function of tidal dissipation energy. It could, however, lead to errors in other kinds of predictions, such as the radius as a function of opacity.  The equilibrium tidal dissipation formula is based on an {\\it ad hoc} assumption of a constant lag angle.  In reality,  the dynamical tidal response of a planet through both gravity and inertial waves near the planet's surface and convective envelope may be much more intense, especially through global normal modes.  Their dissipation may provide the dominant angular momentum transfer mechanism for the orbital evolution and heating sources for the internal structure of close-in extrasolar planets.  In the limit of small viscosity, the intensity of tidal dissipation is highly frequency dependent (Ogilvie \\& Lin 2004).  When the forcing and response frequencies are in resonance, the energy dissipation rate is intense whereas between the resonances it is negligible.  As the planets undergo structure adjustments, their spin frequency, Brunt--V\\\"ais\\\"al\\\"a frequency distribution, the adiabatic index, and equation of state also evolve.  Since all of these physical effects contribute to the planets' dynamical response to the tidal perturbation from their host stars, their response and resonant frequencies are continually modified.  The results in this paper indicate that the structure of the planet adjusts on a radiation transfer time scale which generally differs from the time scale for a planet to evolve through the non resonant region.  In addition, the tidal forcing frequency also changes as the planets evolve toward a state of synchronous spins and circular orbits.  Therefore, it is more appropriate to consider a frequency averaged tidal dissipation rate. In the limit of small viscosity, the frequency averaged dissipation rate converges (Ogilvie \\& Lin 2004) such that the equilibrium tidal dissipation formula may be a reasonable approximation.  Nevertheless, we cannot yet rule out the possibility that some close-in planets may attain some non resonant configuration and stall their orbital evolution.  Therefore, accurate measurement of the sizes of close-in young Jupiters via planet transit surveys can be used to constrain the theories of tidal dissipation and hence internal structure for these objects."
    },
    "0403/astro-ph0403610_arXiv.txt": {
        "abstract": "The mass discrepancy in disk galaxies is shown to be well correlated with acceleration, increasing systematically with decreasing acceleration below a critical scale $a_0 \\approx 3700\\kmskpc = 1.2 \\times 10^{-10}\\;\\mathrm{m}\\,\\mathrm{s}^{-2}$.  For each galaxy, there is an optimal choice of stellar mass-to-light ratio which minimizes the scatter in this mass discrepancy-acceleration relation. The same mass-to-light ratios also minimize the scatter in the baryonic Tully-Fisher relation and are in excellent agreement with the expectations of stellar population synthesis. Once the disk mass is determined in this fashion, the dark matter distribution is specified. The circular velocity attributable to the dark matter can be expressed as a simple equation which depends only on the observed distribution of baryonic mass. It is a challenge to understand how this very fine-tuned coupling between mass and light comes about. ",
        "introduction": "The masses of stellar disks and the distribution of mass in dark matter halos pose a coupled problem. Rotation curves provide good measures of the mass enclosed by disks. But it has been difficult to disentangle how much of this mass is in the stellar disk, and how much is in the dark matter halo. Consequently, both the mass of the stellar disk and the distribution of the dark matter, $\\rho(r)$, have been unclear.  There have long been suggestions of a close connection between mass and light in spiral galaxies.  Perhaps the most obvious is the Tully-Fisher relation (Tully \\& Fisher 1977). Beyond this global scaling relation, there are indications of a local coupling between mass and light (e.g., Rubin \\etal\\ 1985; Bahcall \\& Casertano 1985; Persic \\& Salucci 1991).  One manifestation of this is in the efficacy of maximum disk (e.g., van Albada \\& Sancisi 1986) in describing the inner parts of rotation curves.  If one scales up the stellar contribution to the rotation curves of high surface brightness (HSB) spirals to the maximum allowed by the data, one often finds a good match in the details (the ``bumps and wiggles'') between the shape of the rotation curve and that predicted by the observed stellar mass (e.g., Kalnajs 1983; Sellwood 1999; Palunas \\& Williams 2000).  This only works out to some radius where dark matter must be invoked, but does suggest that the preponderance of the mass at small radii is stellar.  Beyond that, it is often possible to scale up the gas component to explain the remainder of the rotation curve (Hoekstra, van Albada, \\& Sancisi 2001).  Moreover, while some dark matter may be needed to stabilize disks, detailed analyses of disk stability frequently require rather heavy disks in order to drive the observed bars and spiral features (e.g., Athanassoula, Bosma, \\& Papaioannou 1987; Debattista \\& Sellwood 1998, 2000; Weiner, Sellwood, \\& Williams 2001; Fuchs 2003a; Bissantz, Englmaier, \\& Gerhard 2003; Kranz, Slyz, \\& Rix 2003).  While these lines of evidence favor nearly maximal disks in HSB spirals, there are contradictory indications as well.  The most significant of these is the lack of surface brightness (or scale length) residuals in the Tully-Fisher relation (Sprayberry \\etal\\ 1995; Zwaan \\etal\\ 1995; Hoffman \\etal\\ 1996; Tully \\& Verheijen 1997). The apparent lack of influence of the distribution of disk mass on the Tully-Fisher relation suggests that disks are submaximal (McGaugh \\& de Blok 1998a; Courteau \\& Rix 1999). However, galaxies which occupy the same location in the Tully-Fisher plane can have very different rotation curve shapes (de Blok \\& McGaugh 1996; Tully \\& Verheijen 1997).  This excludes the simple hypothesis that galaxies of equal luminosity reside in identical halos with no significant influence from the disk.  Recent data for low surface brightness (LSB) disks complicate matters further. These objects show large mass discrepancies down to small radii, implying that they are dark matter dominated with very submaximal disks (de Blok \\& McGaugh 1997).  However, it is often formally possible to obtain a fit with something like a traditional maximum disk (where the peak velocity of the disk component is comparable to $V_{flat}$), albeit at the cost of absurdly high ($> 10\\; \\Msun/\\Lsun$) mass-to-light ratios (de Blok \\& McGaugh 1997; Swaters, Madore, \\& Trewhella 2000; McGaugh, Rubin, \\& de Blok 2001). As anticipated by McGaugh \\& de Blok (1998b), density wave analyses imply nearly maximal disks for LSB galaxies (Fuchs 2002, 2003b). These high mass-to-light ratios are unlikely for stellar populations, so one might consider a disk component of dark matter in addition to the usual halo.  This seems contrived, and also causes problems with the baryonic Tully-Fisher relation (McGaugh \\etal\\ 2000; Bell \\& de Jong 2001). This relation between mass and rotation velocity works best, in the sense of having minimal scatter, for disk masses which are consistent with stellar population mass-to-light ratios (\\S 3). Maximal disks in LSB galaxies increase the scatter in the baryonic Tully-Fisher relation.  Among these apparently contradictory lines of evidence, there is nevertheless a clear theme.  The luminous and dark components are intimately linked.  This is true not only in a global sense, but also in a local one. This might be paraphrased as Renzo's Rule: {\\it ``For any feature in the luminosity profile there is a corresponding feature in the rotation curve''} (Sancisi 2003).  The distribution of baryonic mass is completely predictive of the distribution of dark matter, even in dark matter dominated LSB galaxies.  Renzo's rule is an empirical statement which is mathematically encapsulated by MOND (Milgrom 1983).  MOND is a modified force law hypothesized as an alternative to dark matter, and remains a viable possibility (Sanders \\& McGaugh 2002).  Even if dark matter is correct (as is widely presumed), MOND is still useful as a compact description of the mass discrepancy in spirals (Sanders \\& Begeman 1994).  In this paper, I show that the stellar mass-to-light ratios determined from MOND fits to rotation curves are optimal in a purely Newtonian sense. I derive a simple expression for the corresponding dark matter distribution, and generalize this to apply for {\\it any\\/} choice of stellar mass-to-light ratio. This expression provides a simple yet stringent test for dark matter theories which seek to explain rotation curves. ",
        "conclusions": "There is a strong relation between the distribution of baryonic and total mass in disk galaxies.  The shape of the rotation curve predicted by the observed distribution of baryons is homologous to the observed rotation.  The mass discrepancy, defined as the ratio of the gradients of the total to baryonic gravitational potential, can be described by a simple function of centripetal acceleration: \\begin{displaymath} \\Phi'/\\Phi_b' = {\\cal D}(x), \\end{displaymath} where $x = a/a_0$ and ${\\cal D}(x)$ is the inverse of equation~6. For $a > a_0 = 1.2 \\times 10^{-10}\\;{\\rm m}\\,{\\rm s}^{-2}$, there is no apparent need for dark matter.  For $a < a_0$, the amount of dark matter increases systematically as acceleration declines.  The empirical organization apparent in the data provides constraints on mass models which have not been considered in traditional disk-halo modeling. Making full use of the information present in the data provides a method for estimating the stellar mass-to-light ratio of spiral disks.  These mass-to-light ratios are optimal in that they \\begin{itemize} \\item minimize the scatter in the mass discrepancy-acceleration diagram (Fig.~5), \\item minimize the scatter in the baryonic Tully-Fisher relation (Fig.~6), and \\item are consistent with the expectations of stellar population models (Fig.~7). \\end{itemize} Indeed, the mass-to-light ratios are consistent not only with the mean value anticipated for stellar populations, but they also reflect the expected trends with color and the band-pass dependent scatter about the mean color-\\ML\\ relation.  It is hard to imagine that all this could be the case unless these mass-to-light ratios are essentially correct.  This appears to solve the long standing problem of the absolute stellar mass of galactic disks.  Once the stellar mass of a disk is fixed, the distribution of dark matter follows.  This can be expressed as a simple function of observable quantities: \\begin{displaymath} \\Vh^2(r) = \\left[\\Q^{-1} {\\cal D}(x) -1\\right] \\Vst^2(r) + \\left[{\\cal D}(x) - 1\\right] \\Vg^2(r) \\end{displaymath} (equation~22).  The dark matter distribution is completely specified by the observed baryonic matter distribution.  This presents a serious fine-tuning problem for any theory of galaxy formation.  The only freedom available in the detailed distribution of dark matter inferred from observation is encapsulated in a single parameter \\Q.  This parameter allows for the unlikely possibility that the actual mass-to-light ratios of stars differs from the optimum value described here.  With this limited freedom, the relation derived between baryonic and total mass specifies the dark matter distribution with complete generality.  Explaining how this strong coupling between mass and light arises is a major challenge for galaxy formation theory. All we have done is restate, in a generalized Newtonian fashion, the well-established result that MOND fits rotation curves. The question, of course, is what this means. There are two independent issues here which are often mistakenly conflated. One is the unconventional theory known as MOND. The second is the empirical regularity of the data. If the effective force law apparent in the data is not in fact MOND, then we need to understand how it comes about in the context of dark matter.  The possibilities are limited.  Either \\begin{itemize} \\item MOND is essentially correct, or \\item Dark matter results in MOND-like behavior in disk galaxies. \\end{itemize} This is the same conclusion reached by Sanders \\& Begeman (1994) and McGaugh \\& de Blok (1998b).  The improvement of the data since that time only makes this result more clear.  If dark matter is correct, the tightness of the observed coupling between baryons and dark matter implies a very strong regulatory mechanism. Within the context of dark matter, there seem to be two basic options. Either \\begin{itemize} \\item the processes of galaxy formation lead to the observed coupling, or \\item there is a direct interaction between dark matter and baryons. \\end{itemize}  In the first case, some combination of mundane astrophysical effects (e.g., adiabatic contraction, mergers, feedback) are presumed to be the root cause of the coupling.  However, no clear mechanism is known which has the required effect. Indeed, it seems extremely unlikely that the chaotic processes of galaxy formation can give such a highly ordered result, much less the finely-tuned coupling between mass and light.  Alternatively, the strong coupling between dynamical and baryonic mass might provide a hint about the nature of the dark matter itself. Rather than interacting with baryons only through gravity, there may be some direct interaction which results in the observed coupling. This amounts to a modification of the nature of the dark matter itself. One obviously wishes to retain the successes of CDM on large scales, but modify it to have the appropriate effect in individual galaxies.  Many modifications of dark matter have been discussed in recent years (e.g., warm dark matter, self-interacting dark matter), motivated at least in part by the difficulties posed by rotation curves.  Those proposed modifications which ignore the baryons would not seem to help.  It is not adequate to change the dark matter properties in order to insert a soft core in dark matter halos; one must explain equation~22.  A mechanism which provides for the direct connection observed between dark matter and baryons is unknown.  Ideas along this line are only now being discussed (e.g., Piazza \\& Marinoni 2003), so it is too soon to judge whether they are viable.  Irrespective of which possibility might seem to hold the most promise, there is clearly some important physics at work which has yet to be understood."
    },
    "0403/astro-ph0403426_arXiv.txt": {
        "abstract": "{The eclipsing supersoft X-ray binary CAL 87 has been observed with Chandra on August 13/14, 2001 for nearly 100 ksec, covering two full orbital cycles and three eclipses. The shape of the eclipse light curve derived from the zeroth-order photons indicates that the size of the X-ray emission region is about 1.5 \\rsun. The ACIS/LETG spectrum is completely dominated by emission lines without any noticeable continuum. The brightest emission lines are significantly redshifted and double-peaked, suggestive of emanating in a 2000 km/s wind. We model the X-ray spectrum by a mixture of recombination and resonant scattering. This allows us to deduce the temperature and luminosity of the ionizing source to be $kT \\sim 50-100$ eV and $L_X \\sim 5 \\times 10^{37}$ erg/s. }    \\addkeyword{binaries: close} \\addkeyword{Stars: individual: CAL 87} \\addkeyword{X-ray: stars}  \\begin{document} ",
        "introduction": "\\label{sec:intro}   CAL 87 was detected with the IPC onboard the Einstein X-ray observatory during the Columbia Astrophysical Laboratory survey of LMC (Long \\etal\\ 1981), and it was optically identified with an eclipsing binary in the Large Magellanic Cloud with an orbital period of 10.6 hrs (Pakull \\etal\\ 1988, Callanan \\etal\\ 1989, Cowley \\etal\\ 1990). A shallow X-ray eclipse was discovered with ROSAT (Schmidtke \\etal\\ 1993).  Early attempts to model the optical light curve of CAL 87 by Callanan \\etal\\ (1989) already indicated an elevated accretion disk rim. This model has been extended by Schandl \\etal\\ (1997), and successfully describes the optical light curve of the CAL 87 binary system, composed of a primary (WD) with a mass of M$_1$ = 0.75 M$_{\\odot}$, placed at a distance of 2.2$\\times$10$^{11}$ cm from the mass-donating secondary star with a mass of M$_2$= 1.5 M$_{\\odot}$. The model of Schandl \\etal\\ (1997) includes (1) optical emission from the secondary star which is irradiated by the emission of the WD and (2) emission of an accretion disk with a thick rim and an optically thick, cold, clumpy spray produced by the high mass-flow rate of the accretion stream impinging on the disk (hot spot). This spray, moving around the disk, nicely reproduced the asymmetry in the optical light curve and the depth of the secondary dip.  Based on the pre-Chandra, low-resolution X-ray spectra and the detection of emission up to 1 keV, CAL 87 has long been considered as one of the two hottest known SSS, thus spurring observations with all X-ray satellites since ROSAT (Parmar \\etal\\ 1997, Asai \\etal\\ 1998, Dotani \\etal\\ 2000, Ebisawa \\etal\\ 2001).  We have exploited the unique spectral capability of the Chandra low-energy transmission grating (LETG) to study in detail the X-ray emission of the canonical supersoft source CAL 87. ",
        "conclusions": "The shape of the eclipse light curve, the lack of any continuum emission and the dominance of emission lines due to resonant scattering and recombination suggests that CAL 87's X-ray emission comes from an extended wind (or outflow), similar to the accretion-disk-corona low-mass X-ray binary 4U 1822-37, as already suggested by Schmidtke (1993). However, the plasma in 4U 1822-37 (i) has a much higher ionization as evidenced by the H- and He-like ions of Ne, Mg, Si, S and Fe which are completely absent in CAL 87 and (ii) is much more dominated by recombination, indicative of larger column densities than those in CAL 87.   \\medskip \\outputfulladdresses  The wind is probably not spherical, but rather has a conical shape with an opening angle of 120\\degs, extending on both sides of the accretion disk. During eclipse, the donor  occults the wind on the southern disk side.   The primary X-ray emission region in CAL 87 certainly has a temperature substantially lower than $\\sim$1 keV  due to the lack of Fe L, Si XIII or Ne IX/Ne X lines. This rules out a neutron star or black hole (except for retrograde rotation) accretor, since they would exhibit a maximum temperature of the inner part of the accretion disk of 0.9--1.2 keV. Also, the outflow velocity of $\\sim$2000 km/s is consistent with a white dwarf primary.  A corona with an ionization parameter and electron density ranging from $>$3000/10$^{12}$ cm$^{-3}$ near the accretor to $\\sim$30/10$^{11}$ cm$^{-3}$ near the outer boundary would scatter about 1\\% of the irradiated X-ray emission. This suggests that the intrinsic X-ray luminosity of CAL 87 is of order a few times 10$^{37}$ erg/s.  Though the earlier SSS classification was based on a wrong interpretation of the X-ray spectrum, the new kT and L estimates from our LETG spectrum re-confirm the supersoft source nature of CAL 87."
    },
    "0403/astro-ph0403389_arXiv.txt": {
        "abstract": "We present a survey of \\ion N5 and \\ion O6 (and where available \\ion C4) in the Galactic halo, using data from the {\\it Far Ultraviolet Spectroscopic Explorer} (FUSE) and the {\\it Hubble Space Telescope} (HST) along 34 sightlines.  These ions are usually produced in nonequilibrium processes such as shocks, evaporative interfaces, or rapidly cooling gas, and thus trace the dynamics of the interstellar medium.  Searching for global trends in integrated and velocity-resolved column density ratios, we find large variations in most measures, with some evidence for a systematic trend of higher ionization (lower \\ion N5/\\ion O6 column density ratio) at larger positive line-of-sight velocities.  The slopes of log[N(\\N)/N(\\O)] per unit velocity range from $-0.015$ to $+0.005$, with a mean of $-0.0032 \\pm 0.0022$(r)$\\pm 0.0014$(sys) dex~(km~s\\ts{-1})\\ts{-1}. We compare this dataset with models of velocity-resolved high-ion signatures of several common physical structures. The dispersion of the ratios, \\ion O6/\\ion N5/\\ion C4, supports the growing belief that no single model can account for hot halo gas, and in fact some models predict much stronger trends than are observed. It is important to understand the signatures of different physical structures to interpret specific lines of sight and future global surveys. ",
        "introduction": "We are interested in understanding how the dynamics of the interstellar medium (ISM) of galaxies determines how the energy and matter released by stars are redistributed.  In a previous paper \\citep{remy04}, hereafter denoted Paper~I, we investigated theoretical models of the ionization ratios of Li-like absorbers, \\ion N5, \\ion O6, \\ion C4, in the Galactic halo. These ions are believed to form in nonequilibrium in shocks, evaporative interfaces, or rapidly cooling gas, all of which trace the dynamics of the interstellar medium.  As a useful new diagnostic, in that paper we focused on {\\it velocity-resolved} signatures of several common physical structures: (1) a cooling Galactic fountain flow that rises, cools, and recombines as it returns to the disk; (2) shocks moving toward the observer; (3) a conductive interface with the observer located in the hotter gas.  In this paper, we present a survey of 34 sightlines observed through the Galactic halo by the {\\it Far Ultraviolet Spectroscopic Explorer} (FUSE) and the {\\it Hubble Space Telescope} (HST) in the Li-like absorption lines of \\ion O6, \\ion N5, and sometimes \\ion C4.  We report the integrated column densities of these three ions and their ratios.  We also focus, as in Paper~I, on the new, and potentially useful diagnostic provided by the velocity dependence of these ion ratios.  The observable velocity-ionization trends are weak, because even strong trends are washed out by the large thermal width of the gas at different parts of the flow.  Additional confusion results since the long sightlines almost definitely pass through multiple structures.  We report an indication of a weak trend of decreasing N(\\N)/N(\\O) at more positive velocities.  The distribution of this slope is broad, with a mean value, $-0.0032\\pm0.0022$(r) $\\pm0.0014$(sys)~dex~(km~s\\ts{-1})\\ts{-1}, displaced somewhat from zero.  In this paper, \\S~\\ref{data} describes the data calibration and analysis. The basic observational results are described in \\S~\\ref{results}. Section \\ref{interpretation} describes models of the dynamical signatures of the Li-like ions, and interpretation. ",
        "conclusions": "We consider the diagnostic power of velocity-resolved column density ratios in understanding the Galactic halo.  Column density ratios of Li-like ions in the Galaxy are useful to diagnose the physical formation mechanism of the gas and to study the interstellar gas cycle, and a survey of these ions can reveal general trends.  We find, in a survey of sightlines observed with FUSE and HST, that the distribution of N(\\N) and N(\\O) in the halo does not appear to favor a dominant physical production mechanism.  We find a possible weak trend of decreasing N(\\N)/N(\\O) at more positive velocities ($-$0.0032$\\pm$0.0022(r)$\\pm$0.0014(sys)~dex~(km~s\\ts{-1})\\ts{-1}, Figure \\ref{histo-slope}).  The weakness of this trend also argues against halo sightlines dominated by a single structure, e.g. the Local Bubble interface.  In Paper~I we presented models of interfaces and cooling nonequilibrium gas, and the velocity-resolved N(\\N)/N(\\O) signatures of each.  Here, we consider the model of smoothly distributed gas corotating on cylinders with the Galactic disk, and implications for the scale heights of \\N\\ and \\O. Observable velocity-ionization trends are weak, because even very strong trends are washed out by the large thermal width of the gas at different parts of the flow.  Additional confusion results because the long sightlines almost definitely pass through multiple structures.  In fact, the dispersion of N(\\N)/N(\\O), both integrated and velocity-resolved, clearly indicates that no single production scenario known to date can completely explain the Galactic halo.  To truly understand the physical production of Li-like ions in the halo, one needs to analyze gas in localized areas of physical space, rather than velocity space.  Absorption spectroscopy towards many halo stars with close angular separation and different distances could help to isolate gas at a specific altitude. Similarly, the gas above known superbubble shells or chimneys could be isolated.  These observations have a greater chance of distinguishing between models of hot gas production than observations along long lines of sight."
    },
    "0403/astro-ph0403176_arXiv.txt": {
        "abstract": "We have studied published data from the {\\it Yohkoh} solar X--ray mission, with the purpose of searching for signals from radiative decays of new, as yet undiscovered massive neutral particles. This search is based on the prediction that solar axions of the Kaluza--Klein type should result in the emission of X--rays from the Sun direction beyond the limb with a characteristic radial distribution. These X--rays should be observed more easily during periods of quiet Sun. An additional signature is the observed emission of hard X--rays by SMM, NEAR and RHESSI. The recent observation made by RHESSI of a continuous emission from the non--flaring Sun of X--rays in the 3 to $\\sim 15$~keV range fits the generic axion scenario. This work also suggests new analyses of existing data, in order to exclude instrumental effects; it provides the rationale for targeted observations with present and upcoming (solar) X--ray telescopes, which can provide the final answer on the nature of the signals considered here. Such measurements become more promising during the forthcoming solar cycle minimum with an increased number of quiet Sun periods. ",
        "introduction": "In order to solve the strong CP problem (why the strong interaction contrary to weak interactions does not violate CP symmetry), a new neutral particle, the axion, with spin--parity $0^-$ has been invented \\citep[for recent review articles see, e.g.][]{raffeltg99,raffeltg00,bradley}. Axions, along with Weakly Interacting Massive Particles (WIMPs), are the two leading particle candidates for dark matter in the Universe. If axions exist, they should be abundantly produced inside the solar core. An axion can be seen as a very light neutral pion ($\\pi ^0$), with restmass  $\\sim 10^{-3\\pm 3}$ eV/c$^2$. However, while the  $\\pi ^0$ interacts strongly and decays to two photons with a mean lifetime of $\\sim 10^{-16}$ s, the axion interacts very feebly with matter and is expected to decay to two photons ($a\\rightarrow 2\\gamma$) with a lifetime much longer than the age of the Universe. Therefore, detection techniques utilise the axion interaction with the electric field of atoms of underground detectors or strong magnetic fields, in order to (coherently) convert them to real photons. In other words, electric or magnetic fields play the role of a catalyser, which can  transform axions into detectable photons. If very light axions ($m \\leq $ few eV/$c^2$) are produced inside the Sun, their thermal energy peaks at $\\sim 4.2$ keV and they are ultra-relativistic. They mostly stream out of the Sun and can have an impact only on the evolution of a Star through the additional escaping energy (as it happens with neutrinos).  However, in recent theories of extra--dimensions, proposed as extensions of the Standard Model, the `conventional', almost massless axions become as massive as the reaction energies involved. In the case of the solar axions, the expected mass spectrum of all the excited Kaluza--Klein (KK) states extends all the way to $\\sim 10$ keV/c$^2$ \\citep{kk,dl,dilella}. These high KK--masses imply a relatively shorter lifetime ($\\tau \\sim 10^{20}$s), because of the $\\tau \\sim m^{-3}$ dependence. The underlying axion--photon--photon coupling constant, $g_{a\\gamma \\gamma}$, which defines the interaction cross-section with ordinary matter, is the same for the `conventional' almost massless axion and for the massive KK--axions. In this work, the KK--axions are taken as a generic example of particles which can be created inside the hot solar core. A small fraction of them ($\\sim 10^{-7}$, as it has been estimated by \\citet[hereafter DZ03]{dilella}), are extremely nonrelativistic and they can be gravitationally trapped by the Sun itself in orbits where they accumulate over cosmic times. As shown in DZ03, their density increases enormously near the solar surface. The estimated mean distance of the KK--axion population from the Sun surface is $\\sim 6.2\\,R_{\\odot}$.  If axions of the KK--type are gravitationally trapped and decay to two photons,% \\footnote{Following the decay mode $a\\rightarrow \\gamma \\gamma$, the two photons have the same energy and are emitted back--to--back, because of momentum conservation, since they are highly non-relativistic.} then the observed X--ray surface brightness from the solar disk and limb can be a signature of the solar axion scenario. In this work, we discuss the expected surface density profile of the derived axion halo around the Sun (DZ03), and we compare it with published data  taken by the {\\it Yohkoh} soft X--ray telescope from the diffuse emission of two quiet Sun observations. In addition, the measured continuous emission of hard X--rays (below $\\sim15\\mbox{ keV}$) from the non--flaring Sun during the past $\\sim25$ years by different orbiting X--ray detectors is also considered. ",
        "conclusions": "Our work suggests that the importance of quiet--Sun solar X--rays should not be overlooked. This paper should provide motivation to the solar X--ray community to follow experimentally the arguments presented here and to reduce the uncertainties of relevant observations. The period during the (next) solar cycle minimum should provide still cleaner quiet Sun  conditions. For example, offpointing observations with the RHESSI solar X-ray telescope can result to interesting data above $\\sim 3$ keV.~Even the pointing tests of the spacecraft on the 8 and 23 May 2003 might well contain useful solar data, in spite of the announcement that `data in this period should not be used for solar studies'\\,% \\footnote{http://hesperia.gsfc.nasa.gov/hessi/ and http://hessi.ssl.berkeley.edu/$\\sim$dsmith/hessi/HME.html} $\\!\\!$. Similarly, we suggest that the announced offpointing to the Crab Nebula from 16 June 2003 should be evaluated in the light of this work. From Figure 2 and the energy threshold of the RHESSI detectors, it follows that the relevant angular extension from the surface of the Sun is approximately $\\theta \\leq 2^{\\circ}$.  Other non--solar X--ray telescopes can be of potential use for this work, provided they can reduce the elongation angle with the Sun. Then, instead of searching for solar axions only here at the site of the Earth, (solar) X--ray telescopes could provide new physics results associated with the solar X--ray halo. At present, these orbiting instruments are highly sensitive to X--rays below $\\sim 10$ keV, which is  within the region-of-interest for many astrophysical objects.  In order to explain the experimental approach of this work in a model independent way, an exotic effect like the axion scenario of this paper can be searched for in the following type of solar observations, provided conventional solar dynamical effects can be excluded: a) an off--pointing observation results to an increased radiation level while approaching the Sun. This should be the manifestation of a halo around the Sun {\\it whatever} its radial distribution, and/or b) the measured level of {\\it any} cosmic radiation from the blank sky does not diminish completely (according to the performance of the telescope) while pointing towards the solar disk, i.e.\\ solar shadow effect is missing. The search for such a residual radiation from the solar disk direction might be the most sensitive approach, especially at photon energies above a few keV, where practically no emission is expected from the quiet Sun.  {\\it In summary}, the radial distributions from the re--considered two {\\it Yohkoh} X--ray  observations of the quiet Sun (including the derived inward heat flux in the solar atmosphere of {\\it some} nonthermal energy deposition beyond $\\sim 1$ solar radius from the Sun surface) can be reconciled with a halo of decaying massive particles near the solar surface. New analyses of existing data may definitely clarify that instrumental scattering effects from bright points on the solar disk are  small. More extended radial X--ray distributions (preferentially during a solar cycle minimum) can provide important information for the quiet Sun. An additional and independent evidence in favour of the axion scenario is the observed {\\it continuous} emission of X-rays from the non-flaring Sun in the 3 to $\\sim 15$ keV range. All these observations, when considered together, might suggest a non--conventional mechanism of the type assumed here. New X--ray measurements could clarify the nature of these relatively high energy effects around the Sun, for which an alternative conventional explanation is still missing."
    },
    "0403/astro-ph0403495_arXiv.txt": {
        "abstract": "High resolution spectral studies were undertaken at orbital phases ($\\phi$) 0, 0.25 and 0.5 on the high-mass X-ray binary (HMXB) Vela X-1 using archival Chandra data. We present (a) the first detailed analysis of the multiple strong narrow emission lines present in $\\phi$ 0.5  (b) an analysis of the absorption of the continuum in  $\\phi$ 0.5, and (c) the first detection of narrow emission and absorption lines in $\\phi$ 0.25. Multiple fluorescent and H-and He-like emission lines in the band 1.6 - 20 Angstrom (\\AA) in eclipse are partially obscured at $\\phi$ 0.25 by the X-ray continuum. The $\\phi$ 0.25 spectrum displays 3 triplets, 2 with a blue-shifted resonance (r) line in absorption and the intercombination (i) and forbidden (f) lines in emission, and shows in absorption other blue-shifted lines seen in emission in eclipse. At $\\phi$ 0.5 the soft X-ray continuum diminishes revealing an \"eclipse-like\" spectrum, however line flux values are around 13-fold those in eclipse. We conclude the narrow emission lines in Vela X-1 become apparent when the continuum is blocked from line of sight, either by eclipse or by scattering and/or absorption from a wake or cloud. The H-and He-like lines arise in warm photoionised regions in the stellar wind, while the fluorescent lines (including a Ni K$\\alpha$ line) are produced in cooler clumps of gas outside these regions. Absorption of the 5-13 $\\mbox{\\AA} $ continuum at $\\phi$ 0.5 may be caused by an accretion wake comprised of dense stagnant photoionized plasma inside a Stromgren zone. Multiple fluorescent emission lines may be a common feature of the supergiant category of HMXBs. ",
        "introduction": "High-mass X-ray binaries (HMXBs) including Vela X-1 have served to illuminate the evolution of giant binary stars, to probe their radiation-driven winds, and to study accretion processes and the fundamental parameters of compact objects. The 130 known HMXBs fall into 2 categories, Be/X-ray and supergiant systems \\citep{liu00}. Vela X-1 is one of 11 supergiant systems (which include GX 301-2 and 4U 1700-37), while the vast majority (around 80\\%) of HMXBs are Be/X-ray systems detected as transients \\citep{kap01,whi95}. Supergiant systems have a relatively short life of 10,000 years before the secondary star explodes in a supernova leaving a neutron star binary or 2 single neutron stars \\citep{kap01}. Once a massive main-sequence star evolves to a supergiant the Roche lobe may become filled and accretion flow becomes sufficient to power a strong X-ray source.  Vela X-1 is an eclipsing system at an estimated distance of 1.9 kpc \\citep{sad85} containing an B0.5Ib supergiant (mass 23 $M_\\odot$ and radius 34 $R_\\odot$) and a wind-fed pulsar with a period of 283 s \\citep{nag89}. The neutron star (NS) orbits in a period of 8.96 days very close to the non-synchronously rotating surface of the supergiant with a semi-major axis 53.6 $R_\\odot$\\citep{bar01}, so it is fully embedded in the acceleration zone of the stellar wind.  The NS mass of 1.9 $M_\\odot$ is the highest known \\citep{bar01,van95} and the system has been studied at infrared \\citep{hut02}, UV \\citep{sad85, vanl01}, gamma \\citep{rau94} and at optical \\citep{van95, kap94} wavelengths. It has not been detected at radio frequency (detection limit at mean cm radio flux density (mJy) $<0.2$ \\citep{fen00}.  Orbital modulation of the spectra have been noted in the UV, X-ray, and optical wavelengths. \\citet{sad85} noted absorption of \\ion{Al}{3} and \\ion{Fe}{3} in the UV, limited to the orbital phase ($\\phi$) 0.5 and higher, and postulated the presence of a trailing wake. \\citet{bes75} reported an absorption component in the H$\\alpha$ spectrum around $\\phi$ 0.6.  \\citet{kap94} postulated that an absorption component in the optical spectrum resulted from a photoionization wake that trails the X-ray source, rather than the accretion wake which does not sufficiently obscure the supergiant. \\citet{smi01} studied an asymmetric dip in the UV lightcurve at $\\phi$ 0.46-0.7 and concluded spectral ratios (predip/dip) contain information about the mechanism responsible for the variability. Absorption of soft X-rays during the $\\phi$ 0.5 and higher was first reported in Tenma data \\citep{nag89}. \\citet{fel96} showed that a trailing photoionization wake can explain the asymmetry in the X-ray band during the eclipse ingress and egress.  \\citet{sak99} using ASCA data demonstrated that the spectrum in eclipse was comprised of recombination lines and radiative combination continua (RRC) produced by photoionization in the stellar wind; they suggested the fluorescent K-shell lines from near-neutral atoms indicated that the X-ray -irradiated portion of the wind in Vela X-1 consists of cool dense clumps in a hotter, more ionized gas.  \\citet{shu02a} presented the first high-resolution spectrum during eclipse, and using RRC (including \\ion{Ne}{10} RRC) deduced a temperature of 1.2 $\\times 10^5$ K for the plasma. The spectrum showed clear evidence of photoionization processes, and the resonance (r) lines in the He-like triplets were of roughly equal strength to the forbidden (f) lines, consistent with resonant scattering.  We unexpectedly found that in $\\phi$ 0.5 the  emission lines were stronger than in eclipse by an order of magnitude, associated with the near disappearance of the soft X-ray continuum.  This result, together with the unusual finding of multiple fluorescent emission lines, led us to further investigate emission line production in HMXBs by a study of Vela X-1 comparing eclipse, $\\phi$ 0.25 and $\\phi$ 0.5 using high-resolution Chandra datasets. We outline the data reduction proceedure in section 2, give our results in section 3 and present the discussion in section 4. ",
        "conclusions": "X-ray emission in Vela X-1 results from capture and accretion of gas in the stellar wind of the supergiant star onto the hot spot on the NS. Wind material departs uniformly from the surface of the star, and accretes onto the NS if it passes within a critical distance forming an accretion column. Adjacent material that is perturbed in its flow but is not accreted forms an accretion wake. The X-rays ionize and heat the surrounding gas, and are reprocessed in the stellar wind. Spectra at orbital phases 0, 0.25 and 0.5 reflect the processes of mass loss in the stellar wind of the supergiant and mass accretion onto the NS.  The Doppler shifts at $\\phi$ 0.5 indicate slowing of the supergiant's wind velocity in the vicinity of the NS, consistent with X-ray photoionization and destruction of ions with strong UV resonance transitions that otherwise drive the stellar outflow. The X-ray source creates a Stromgren zone in which the high photoionization conditions effectively quench the driving force and acceleration of the wind.  Spectra at $\\phi$ 0.25 reveal simultaneous evidence of emission lines produced by scattering, and absorption lines that are produced in different regions of the binary system.  Narrow emission lines in Vela X-1 become apparent when the continuum is blocked from line of sight, either by eclipse or by scattering and/or absorption from a wake or cloud. The X-ray source is very compact and is readily occluded by the accretion wake at $\\phi$ 0.5, accounting for the absorption of soft continuum. Absorption of the continuum at $\\phi$ 0.5 has allowed detection of scattered radiation and emission lines whose strength is an order of magnitude higher than those in the eclipse phase. The eclipse and 0.5 spectra have a similiar appearance, and the ratios of lines in the He-like triplets are consistent with photoionization and resonant scattering (involving photoexcitation) in a PI plasma.  Multiple fluorescent emission lines have been detected in 3 of 5 supergiant HMXBs for which high resolution X-ray spectra are available. This study suggests fluorescent emission from externally X-ray illuminated matter may occur from elements other than neutral iron if the direct radiation is blocked and only reflected radiation is visible, when neutral ions of Ne, Mg, Si, S, Ar, Ca, Cr and Ni may be detected.  Differences in Doppler shifts of various lines indicate different origins of lines, from direct radiation from the surface of the NS (lines seen in $\\phi$ 0.25 in absorption) and from adjacent regions in the stellar wind from scattered radiation."
    },
    "0403/astro-ph0403206_arXiv.txt": {
        "abstract": "We investigate in detail the mass distribution obtained by means of high resolution rotation curves of 25 galaxies of different morphological types. The dark matter contribution to the circular rotation velocity is well-described by resorting to a dark component whose density shows an inner core, i.e.  a central  constant density region. We find a very strong correlation between the core radius size $R_C$ and the stellar  exponential scale length $R_D$: $R_C \\simeq  13 \\  (\\frac {R_D} {5\\ {\\rm kpc}})^{1.05} \\ {\\rm kpc} $,  and between  $R_C$ and the galaxy dynamical mass at this distance, $M_{dyn}(R_C)$. These relationships would not be expected if the core radii  were the product of an incorrect decomposition procedure, or the biased result of wrong or misunderstood observational data. The very strong correlation between the dark and luminous scale lengths found here seems to hold also for different Hubble types and opens new scenarios for the nature of the dark matter in galaxies. ",
        "introduction": "\\label{sec:introduction} In spite of a large amount of knowledge gathered over the past 20 years on the phenomenon of Dark Matter (DM) in Galaxies, it is only recently that the attention has been focused on its radial density profile. Recall that, according to high-resolution N-body simulations in the $\\Lambda$-CDM framework, the dark halo density, characterized by the virial halo mass as free parameter, is described by an inner power--law cusp (but see Ricotti 2003): \\begin{equation} \\rho_{\\rm NFW}(r) = \\frac{\\rho_s}{(r/r_s)(1+r/r_s)^2} \\label{eq:rho_nfw} \\end{equation} where $r_s$ and $\\rho_s$ are related parameters (Navarro, Frenk and White, 1996b, hereafter NFW). The DM contribution to the circular velocity can be written as: \\begin{equation} V_{\\rm NFW}^2(r)= V_{vir}^2 \\frac{c}{A(c)} \\frac {A(x)}{x} \\end{equation} where $x \\equiv   c r/R_{vir}$, $c \\equiv R_{vir}/r_s$ is the concentration parameter, and $A(x)\\equiv \\ln (1+x) - x/(1+x)$. We also have that $V^2_{vir}= G M_{vir}/ R_{vir}$, $\\rho_s \\simeq 101/3 \\ \\rho_{crit} \\ c^3 \\ / (\\ln(1+c)-c/(1+c))$ and $M_{vir}  = \\frac{4 \\pi}{3} \\delta_{vir} \\rho_{mean} R^3_{vir}$. \\footnote {$\\rho_{crit}$ is the critical density of the universe, $\\delta_{vir} \\simeq 337$ is the virial overdensity and $\\rho_{mean}$ is the mean universal density at the galaxy's redshift. The values for the numerical constants are valid for a $\\Lambda$-CDM cosmogony with $\\Omega_0$=0.3 and $\\Omega_\\Lambda$=0.7.}  A number of reliable mass modellings, obtained for spirals and low surface brightness galaxies (LSBs) has supported the early claim (Moore, 1994) that dark halos around disk  galaxies have a central density distribution much shallower than the NFW one (e.g. Gentile et al., 2004 and references therein; Weldrake, de Blok \\& Walter, 2003; de Blok \\& Bosma, 2002). The first studies on elliptical galaxies seem to indicate that also  these objects share the same phenomenon (Gerhard et al., 2001;  Borriello et al., 2003 ).  This  ``galaxy by galaxy'' comparison between the predicted $\\Lambda$-CDM density distribution and those actually detected for the dark halos around disk galaxies has been the main goal of several published works. However, although the study of the discrepancy between data and theoretical predictions remains necessary, we believe that time has come that we investigate {\\it per se} the distribution of DM around galaxies, independently of the existing cosmological implications. In fact, the hot debate on possible falsification of $\\Lambda$-CDM on galactic scales has carried the research off from the original topic of understanding the {\\it DM phenomenon} in virialized systems. Let us stress that a {\\it direct} knowledge on the presence, nature and interaction with baryons of the  dark galactic  component  is still very rough and limited, unlike the complex and refined scenario that theory and simulations have put forward.   In the decomposition of rotation curves (RC's), the circular velocity  $V(r)$ is best-fit  modelled as the sum of a baryonic component that  includes a stellar and a gaseous disk, and a spherical dark halo: \\begin{equation} V^2(r)=V_D^2(r)+V^2_{gas}(r)+V^2_H(r) \\end{equation} where the labels D, gas and H refers to the corresponding components.  It is common to represent the dark halo  contribution to the circular velocity by means  of the   2-parameter function describing the circular velocity of a pseudo-isothermal (PI) halo: \\begin{equation} V^2_{H,PI}(r) =  4\\pi G \\rho_0 R_C^2 \\left( 1 - \\frac{R_{C}}{r} {\\rm arctan} \\frac{r}{R_{C}} \\right). \\label{eq:pi} \\end{equation} The central density $\\rho_0$ and the halo {\\it density core} radius $R_C$ are  free parameters to be tuned to fit $V^2(r)$.  Persic, Salucci \\& Stel (1996, PSS hereafter)  introduced, for the  dark halo velocity contribution, an equally  simple function, \\begin{equation} V^2_{H, URC}(r)=(1-\\beta) V_{opt}^2 r^2 (1+a^2)/(r^2+a^2 R_{opt}^2) \\label{eq:urc} \\end{equation} to fit, once   added to the  disk contribution,  the  Universal Rotation Curve (URC) of Spirals,  i.e. the  ensemble of  synthetic RC's, each one   derived from a large number  of individual RC's  belonging to  objects with a luminosity  falling   inside a  fixed  range (see Section 4 of PSS for details).  The two free parameters are: the {\\it velocity} core radius $a R_{opt}$, and $(1-\\beta) \\ V^2_{opt}$, the halo velocity amplitude at $R_{opt}$.\\footnote{$R_{opt}$ is  the radius   encompassing 83\\% of the total luminosity of the galaxy. In  the  case of a (stellar) exponential thin  disk $R_{opt}$ is 3.2  times the disk scale length $R_D$.} For the aim   of this work, in  the region of interest,  $1/3 R_{opt}  \\leq r \\leq  R_{opt}$,  the PI  circular velocity is indistinguishable  from the {\\it halo } URC  of Eq. 5. In fact, the transformation law is obtained from the known  quantities $R_{opt}$ and $V_{opt}$, and by setting the  PI free  parameters $R_C$ and $\\rho_0$ in Eq. 4 at the values  $R_C=  a\\, R_{opt}/1.45$ and $\\rho_0= [  4\\pi G R_C^2 ( 1 - \\frac{R_{C}}{R_{opt}}  {\\rm arctan} \\frac{R_{opt}}{R_{C}})]^{-1}  (1-\\beta) V_{opt}^2$. Let us notice that    $ V_{H, PI}$ and $V_{H,URC}$ agree with each other within few percent at most, a discrepancy that is completely negligible, in view of  the large  radial variations  of  $V(r)$ inside galaxies  and of  the large differences among the  halo velocity profiles  predicted by different  mass models.  Moreover, the  velocity profiles of Eqs. \\ref{eq:pi} and \\ref{eq:urc},  once their free parameters have been properly set, can  mimic in the regions where data are available a number of  ``theoretical'' rotation curves, not  all necessarily consistent  with observations.\\footnote{e.g. the  NFW circular velocity  with   $r_s=R_{opt}$,   in the radial  range,   $0.3< r/r_s<1.1$ is reproduced,  within a negligible  few percent discrepancy,    by  $V_{H, URC} (r; a=0.33 )$ and/or  by  $V_{H, PI}(r; R_C=0.2 R_{opt})$.}  Therefore, with the  obvious caveat that  suitable data are  available,  by investigating the properties of the density distribution of the dark halos around galaxies  by means of  the PI/URC profile   we  neither assume nor reject the theoretical   $\\Lambda$-CDM scenario. On the contrary,   by adopting certain values for  the  free parameters of the chosen  halo density profile in order  to  match the observations, we will be able to  probe  a density cusp, or   a   core  or even  something in between.  In this work,  we first obtain  the density  profile of DM halos around a significant number of galaxies of different Hubble types with high-quality kinematic data. The  intrinsic high quality of the data will allow us  to disentangle the available kinematics in its dark and luminous contributions and to constrain the DM halo parameters allowing us to  investigate  the existence of a   link between their structural  properties  and  those of the ordinary baryonic matter (i.e. stars and gas).  The sample of data used in our analysis is presented  and discussed in Sect. \\ref{sec:sample}, while in Sect. \\ref{sec:model} we describe the mass modelling. In Sect. \\ref{sec:results} we will look for a relationship between the dark  and stellar structural  parameters. Finally the conclusions will be drawn in Sect. \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} There is a mounting evidence that the  kinematics of disk galaxies can be explained in terms of the standard  disk + halo components {\\it only } if the density distribution of the latter decreases with radius  very slowly from the center out to the radius  $R_C$  inside which, instead, the density of CDM halos is predicted to fall  as a power law with an exponent between -1 and -2. Here we have investigated this ``observational\" dark matter scale $R_C$  that is absent in the current theory of galaxy formation. We found that it correlates tightly with the exponential thin disk scale length $R_D$  and with the amount of gravitational mass that it encompasses $M_{dyn}(R_C)$. The high values for the correlation number, the smallness of the scatter around the relations and especially the fact that the pairs $R_C$ and $R_D$, as well as $R_C$ and $M_{dyn}(R_C)$, are measured/derived independently from each others, ensure that we are dealing with {\\it real and physical} relationships among physical quantities.  Three different levels of consequences follow. First, the claims according to which core radii arise as a consequence of 1) serious observational errors, 2) peculiar or biased kinematics, 3) wrong mass modelling that, until today, have been  { \\it by-passed}  by considering a number of test-cases, can be now ruled out. This, on the basis that  human/observational error hardly  relate to intrinsic properties of galaxies. As an example the  suggestion that observers have badly missed the galaxy center and doing so artificially created a core, cannot explain the additional evidence that this core  is found to correlate with the galaxy disk scale length within the extraordinarily small scatter of 0.16 dex.  The second level realises that the core radius,  the quantity that defines the unexplained  feature of the distribution of the dark matter, correlates with the main quantity controlling the distribution of the luminous matter as well as the total mass within this  radius. This indicates that the DM density in galaxies has been shaped by a dark--to--luminous matter coupling, a  very challenging task  if the former is collisionless. However, let us recall that, at least in dwarf systems,  according to  Navarro et al. (1996a), large mass outflows  from  galaxies, arising from supernovae explosions,  can modify   an initial  NFW profile into a cored one,   for which  $R_C \\propto (M_D/R_D)^{0.5}$. From the results of the mass modelling given in PSS, we find $M_D \\propto R_D^3$, in excellent agreement with relationship (9). In any case, our  result points  toward some   new physics of which it  will be  a crucial benchmark.  The third level is  the realization that the core radius vs disk scale length relationship is  Hubble type free, at least   within  a band of   0.16 dex. For the objects in our sample   normal/dwarf spirals and  LSB lay on the same $\\log R_C$  vs $\\log R_D$ relationship. Moreover, pioneering analysis on the DM distribution in ellipticals (Gerhard et al. 2001) finds  halo core radii of sizes proportional to the half light radii (that correspond to 1.67 $R_D$): by analysing the mass model parameters of 21 Giant Ellipticals (shown in Fig 18 Gerhard et al. (2001), and made available by the authors) we find: \\begin{equation} \\log R_C= (1.1 \\pm 0.2 )\\  \\log R_D +(0.1 \\pm 0.4) \\end{equation} in very good agreement with Eq. (9). We caution that  the sample of this work is quite  limited to fully investigate   the morphological dependence  of DM properties and that  in Gerhard et al.(2001)   the data quality  is barely sufficient for the intrinsically complicated  mass  modelling. Nevertheless,  the emerging picture  is truly impressive,  especially  when we consider that core radii of galaxies of different Hubble types  do not  correlate  with  galaxy luminosity: e.g.  a  similar value for  $R_C$  is   found in a  $10^{11} L_{\\odot,r}$ spiral,  in a  $10^{10} L_{\\odot,r}$ LSB and  in a $10^{12} L_{\\odot,r}$ giant elliptical. The core radius  seems to  uniquely relate to the  exponential rate with  which  the stellar surface density decreases with radius.  The next  step of this line of research post-$\\Lambda$-CDM crisis  would be to thoroughly investigate this  DM distribution property in galaxies of different size and Hubble Types  (Salucci et al., 2004) and  the brightest galaxies  because of the slightly non-linear  relationship found between  $R_C$ and the stellar scale length."
    },
    "0403/astro-ph0403030_arXiv.txt": {
        "abstract": "We present the first results of the timing analysis of our {\\emph{RXTE/INTEGRAL}} monitoring campaign on GRS~1915+105. Over the 6 already performed {\\emph{RXTE}} observations, we study the presence of Low Frequency QPO (LFQPO),  and their energetic dependences. In a view to understand the QPO phenomenon, we compare the QPO properties to the spectral behaviour of the source. We propose that part of the compact jet detected during multi-wavelength observations, could produce a significant amount of hard X-rays, and hence explain the energy dependence of the amplitude of the QPOs. ",
        "introduction": "\\begin{table*} \\begin{tabular}{|ccccc|} \\hline Obs. sequence \\# & Date (MJD)&  Good times &  Revolution  \\# & {\\emph{INTEGRAL}} ref. \\\\ ({\\emph{RXTE}}) & & (s) & ({\\emph{INTEGRAL}}) & \\\\ \\hline 1 & 6-7 March 2003 (52704-05)&15768 & Rev. 48 & Hannikainen et al. 2003 \\\\ 2 & 2 April (52731)&9300 & Rev. 57 & Fuchs et al. 2003 \\\\ 3 & 9-10 April (52738-39)&25360 & Rev. 59 & Hannikainen et al. 2004 \\\\ 4 & 9 May (52768)& 14000& Rev. 69 & Hannikainen et al. 2004 \\\\ 5 & 2 November (52945) & & Rev. 122 & Solar Flares\\\\ 6 & 22-23 November(52965-66) &36100& Rev 135 & Hannikainen et al. 2004 \\\\ \\hline \\end{tabular} \\caption{log of the {\\emph{RXTE}} observation and contemporaneous {\\emph{INTEGRAL}} revolutions reported in this paper} \\label{tab:log} \\end{table*} The log of the {\\emph{RXTE}} observations analysed in this paper is reported in Table \\ref{tab:log}. All of them were simultaneous with {\\emph{INTEGRAL}} and other ground based observations. The first observation shows a new class of variability (Hannikainen et al. 2003), while during Obs. 2,3,4, the source has a steady flux in the X/Gamma rays. It shows a strong QPO, a  powerful compact jet during Obs. 2 (Fuchs et al. 2003), and a high level of radio emission during the 2 other observations (Hannikainen et al. 2004, these proceedings). ",
        "conclusions": "We found that the radio loud hard state (a.k.a. Class $\\chi_1$-$\\chi_3$) does not lead to sensible parameters when fitted with a multi colour disc black body and a power-law (see Muno et al. 2001, Trudolyubov 2001) .  A broken power-law  represents the 3-300 keV spectra well. \\\\ \\indent\tDuring the 3 steady observations high level of radio emission is observed (see  Fuchs et al. 2003, Hannikainen et al. these proceedings).  The hard X-rays above 20 keV may originate from the jet (as expected see Markoff et al. 2003; Corbel et al. 2003). The spectra of GRS 1915+105, would then fit better in the standard picture of micro-quasars states.\\\\ \\indent\tOur analysis of the spectra of LFQPOs confirms the presence of a cut-off in their energy dependence (Tomsick \\& Kaaret 2001; Rodriguez et al. 2002), with an evolving energy from $\\sim$15-20 keV in Obs. 2 to a value $ >$25 keV in the remaining Obs. (Fig.\\ref{fig:amplitude}). Although  the QPO cut-off is needed at some point (otherwise its amplitude would grow indefinitely), its evolution is unclear. It can be easily understood, however, if the jet contributes significantly to the hard X-ray, and its flux is not modulated on such a short time scale.  \\\\ \\indent\tThe fact  that a high level of radio emission (at 15 GHz) is found during obs.2, when the cut-off is clearly detected in the QPO spectrum, is compatible with this interpretation. We thus propose that a part of the hard X-ray is emitted by the compact jet, explaining both the spectral behaviour of the source and the QPO ``spectra''."
    },
    "0403/astro-ph0403564.txt": {
        "abstract": "Motivated by observations of inner halo satellite remnants like the Sgr and \\omegacen, we develop fully analytical models to study the orbital decay and tidal massloss of satellites on eccentric orbits in an isothermal potential of a host galaxy halo. The orbital decay rate is often severely overestimated if applying the ChandraSekhar's formula without correcting for (a) the evaporation and tidal loss of the satellite and (b) the contraction of satellite orbits due to adiabatic growth of the host galaxy potential over the Hubble time.  As a satellite migrates inwards, the increasing halo density affects the dynamical friction in two opposite ways: (1) it boosts the number of halo particles swept in the satellite's gravitational \"wake\", hence increasing the drag on the satellite, and (2) it boosts the tide which \"peels off\" the satellite, and reduces the amplitude of the wake. These competing processes can be modeled analytically for a satellite with the help of an empirical formula for the massloss history.  The analytical model agrees with more traditional numerical simulations of tidal massloss and dynamical friction. Rapid massloss due to increasing tides at smaller and smaller radius makes it less likely for streams or remnants of infalling satellites to intrude the inner halo (as the Sgr stream and \\omegacen) than to stay in the outer halo (as the Magellanic stream), hence any intermediate-mass central black holes of the satellites are also likely \"hung up\" at large distances as well.  It is difficult for the satellites' black holes to come close enough to merge into the supermassive black hole in the center of the host potential unless the satellites started with (i) pericenters much smaller than the typical distances to present-day observed satellites and with (ii) central density much higher than in the often seen finite density cores of observed satellites. ",
        "introduction": "Current theory of galaxy formation favors the idea that galaxies form hierarchically by merging smaller lumps or satellites.  The gravity from a moving satellite pulls behind it a \"wake\" of particles of the host galaxy (ChandraSekhar 1943, Mulder 1983). This dynamical friction dissipates the orbital energy of the satellites so that they sink deep into the host galaxy potential well, where they are disintegrated and virialized via baryonic feedbacks and tidal stripping.  These processes might have determined the density profile of virialized halo of the host galaxy (Syer \\& White 1998, Dekel et al. 2003).  There are about 150 and 300 globular clusters, and a few dozen dwarf satellites of mass $10^{6-9}\\msun$ in the Milky Way and M31 respectively. It is tempting to associate these dwarfs satellites and globular clusters as the markers/remnants of past hierarchical merging events. Indeed there are several examples of possible streams of remnants in the Milky Way (Lynden-Bell \\& Lynden-Bell 1995) including the recently found Galactic ring or Carnis Major dwarf galaxy, traced by a grouping of globular clusters (Martin et al. 2004). A giant stream is found in the Andromeda galaxy (McConnachie et al. 2003, Ferguson et al, 2002). Among these the Sagittarius dwarf galaxy stream (Ibata et al. 1997) between radius 10-50 kpc from the Galactic center is perhaps the best example. It brings in at least 5 globular clusters to the inner halo, including M54, the 2nd most massive cluster of the Milky Way. A still mysterious object is \\omegacen, with about $10^6\\lsun$ at 5 kpc from the Galactic center which has the morphology of a globular cluster, but has multiple epoches of star formation and chemical enrichment (see Gnedin et al. 2002 and the \\omegacen symposium).  Another example is the G1 cluster, a very massive globular-like object with about $10^6\\lsun$ at about 40 kpc from the M31 center (Meylan et al. 2001). A system (NGC1023-13) almost identical to G1 is also found in the S0 galaxy NGC1023, at a projected distance of about 40 kpc from the host galaxy (Larsen 2001). Freeman (1993) suggested that such systems are the remnants of nucleated dwarf satellites, with their outer tenuous dark matter and stars being removed by galaxy tides. M31 also has an unusual collection of clusters as luminous as \\omegacen within 5 kpc in projection from its double-peaked center. Given the above evidences or signs for infalling objects in our galaxy and M31, it is interesting to ask whether some of the inner globulars could have also been the result of mergers.  Beyond the Local Group, minor mergers are sometimes speculated as the mechanism to deliver massive black holes and gas material into the nucleus of an AGN to account for the directions of the jets and nuclear dusty disks (Kendall, Magorrian, \\& Pringle 2003).  A very interesting related issue is whether giant black holes acquire part of the mass by merging the smaller black holes in the nuclei of infalling satellites.  For this paper we revisit the basic theoretical questions: what is the condition for a dwarf galaxy to decay into the inner halo? What are the possible outcomes of tidal stripping of a dwarf satellite? How often do we get a system like \\omegacen or a naked black hole near the host center? The answers to these questions will help us to test the validity of the theory of hierarchical merger formation of galaxies. The key mechanism for satellites to enter the inner galaxy is dynamical friction, where the gravity of the satellite creates a wake of overdensities in the particle distribution of the host galaxy, which in turn drags the motion of the satellite with a force proportional to $m(t)^2$, where $m(t)$ is the mass of the satellite. Another process is tidal disruption, where the object sheds mass with each pericenter passage, and the remnants are littered along the orbit of the satellite. The above two processes compete with and regulate each other: orbital decay increases the tidal field, which reduces the mass of the satellite, hence slows down the orbital decay.  Some examples of these effects have been shown in Zhao (2002) in the case of \\omegacen.  The analytical formula of ChandraSekhar  is widely used for gaining insights on dynamical friction because of the time-consuming nature of the more rigorous N-body numerical simulation approach. It is a customary practice in previous works to model the orbital decay of a satellite as a point mass of a fixed mass.  However, {\\it the fixed-mass approximation is invalid} and could seriously overestimate dynamical friction because of neglecting massloss ${dm(t) \\over dt}$. It is essential in calculations of satellite orbits to model the dynamical friction and massloss together since they regulate each other.  In the past the massloss and the orbital decay are often modeled in the {\\it ab initio} fashion, resulting coupled non-linear equations without simple analytical solutions. In such models, the satellite mass is often modeled as a function of the satellite's tidal radius, hence various factors come in, including the orbital position of the satellite, the density profile of the satellite (e.g., Jiang \\& Binney 2000, Zhao 2002, Mouri \\& Taniguchi 2003, Kendall et al. 2003). However, these complications are not always necessary since the massloss history is rather similar in simulations with very different initial conditions (the mass is generally a stair-case like a function of the time), so could be parametrized in an empirical fashion, by-passing the uncertain assumptions of the satellite initial profile.  This could be useful for exploring a large parameter space of the satellite initial condition.  Another invalid approximation but common practice is to use a static potential for the host galaxy. This again is unphysical since galaxy halos do grow in hierarchical formation scenario partly because of galaxy merging, and partly because of the adiabatic contraction of the baryonic disk and bulge; galaxy rotation curves $V_{\\rm cir}(r,t)$ can change by significantly before and after the formation of baryonic disks and bulges in mass models for the Milky Way and M31 (Klypin, Zhao \\& Somerville 2002) and in generic CDM simulations with baryons (e.g. Wright 2003). The growing gravitational force tends to restrain the radial excursions of the satellite while preserving the angular momentum. The dynamical friction or drag force is also proportional to the growing density $\\rho(r,t)$ of ambient stars and dark particles in the host galaxy.  In fact, it is conceptually simple to incorporate massloss and growing potential while keeping the problem analytically tractable: the deceleration $-{dv \\over dt}$ is simply proportional to the satellite mass $m(t)$ and the ambient density $\\rho(r,t)$ of the host galaxy at the time $t$. Without massloss ChandraSekhar's formula would predict very efficient braking of the orbits, enough to make a high-mass satellite of $\\sim 10^{10}\\msun$ (the mass of the LMC or M33 sized object) to decay from a circular orbit at $\\sim 100$kpc to the very center of a high brightness galaxy in a Hubble time, delivering remnants into the inner galaxy. Here we study the effect of massloss and growing potential on the result of orbital decay, and the distribution of remnants. We present fully analytical results for calculating the decay rate for satellites on eccentric orbits in a scale-free growing isothermal potential.  The structure of the paper is following: S2 gives the analytical formulation of the problem, S3 presents results of application to globulars and dwarf satellites, S4 studies the relation between massloss and the satellite density profile. S5 discusses the progenitors of Sgr and \\omegacen, and S6 summarizes. ",
        "conclusions": "\\subsection{Effects of disk, bulge and orbital inclination}  One limitation of the current analysis is that we assume an isothermal dark matter plus stars model throughout the galaxy, hence the dynamical friction effect of the disk and bulge are not modeled accurately.  However, the disk and bulge are not important for our conclusion because we predict mainly the orbital decay in the outer halo where $j>j_{\\rm disk}=3000{\\rm kpc~km/s}$.  Inside 15 kpc, our estimation of dynamical friction by an SIS model is inaccurate only for satellites on low inclination orbits. If satellites come in random inclinations, it is more common to find high inclination orbits, for which our models should be fairly accurate even inside 15 kpc.  \\subsection{Effects of escaping stars}  In the part of our formulation where we derive the satellite mass profile, we assume a simplifying static picture that the satellite's mass is peeled off in successive layers at the shrinking tidal radius.   The picture in N-body simulations is more complicated, since satellite particles at all radii, e.g., the center of the satellite, could in principle be escaping at any time.  So eq. (26) sets only a lower limit on the initial density of the satellite.  A more rigorous model should making this correction, e.g., by introducing an empirical factor to correct this as in Jiang \\& Binney (2000).  \\subsection{Possible orbits of progenitors of the Sgr stream}  The Sgr dwarf and the Canis Major dwarf are the closest known dwarf galaxies, about 15 kpc from the center of the Milky Way and at the edge of the Milky Way disk. The Canis Major is on a (direct or retrograde) orbit slightly inclined from the plane of the Milky Way, and the Sgr is on a nearly polar orbit. Both orbits have a fairly low angular momentum with $S \\sim 20\\kpc$; the data on Sgr are more complete, and show that it oscillates between 10 kpc pericenter and 50 kpc apocenter.  Both contain several globular clusters. It is possible that the two dwarfs are the stripped-down version of a more massive object, which has dynamically decayed from the outer halo.  Interestingly a possible extension of the Sgr has been reported recently in the SDSS data near the position of the outer halo globular cluster NGC2419 (Newberg et al. 2003). There is a stream-like enhancement of halo A-colored stars at the SDSS magnitude of $g_0=20.3$ in the plane of the Sgr's orbit, corresponding to a distance of 90 kpc. If this is true, it would imply that the Sgr has changed its orbits in the past Hubble time.  There are two possible ways that this could happen.  One is that the Sgr's orbit has been deflected by a massive satellite, such as the LMC or SMC.  Indeed the orbits of the Sgr and the Magellanic Clouds do overlap at the Galactic poles, and simple timing arguments show that these systems encounter or fly by each other about 2.5 Gyrs ago at about 50 kpc on the North Galactic Pole if the rotation curve of the Milky Way is nearly flat (Zhao 1998).  The problem of this solution is that it is rare for the Sgr to receive a strong enough deflection to bring down its orbit.  Another solution is that the Sgr has been a more massive system, which orbital decayed from the outer halo (Jiang \\& Binney 2000).  {\\it Our Model B illustrates such an example of the progenitor of the Sgr} which had an initial mass comparable to the LMC ($10^{10}\\msun$) and was on an eccentric orbit with radius between $20-140\\kpc$ (cf Fig.6).  This model is similar to the Model K of Jiang \\& Binney. After a Hubble time the orbit decays to a small orbit very much like that of the Sgr with peri-to-apo ratio of $10\\kpc:50\\kpc$.  Large amount of the material of the progenitor is shed in the radius between $10-140\\kpc$, the stream near NGC2419 at 90 kpc could be part of this debris near one of the apocenters of the orbit.  Unfortunately the present model ends with a mass of $5\\times 10^9\\msun$, too large for the present-day Sgr. Some fine tuning of initial conditions and detailed N-body simulations are clearly needed to test this idea.  \\subsection{Possible orbits of the progenitor of \\omegacen}  We have mainly concentrated on the problem of getting rid of a satellite's angular momentum if it starts with a high angular momentum or orbital size $S_0\\gg 15$kpc. What would be the remnant distribution if a satellite is born with an initial orbital size $S_i<15$kpc? The stars in such a system are assembled in the inner halo from the start, e.g., by colliding an infalling gas cloud with the protogalactic gas clouds in the inner halo, (Fellhauer \\& Kroupa 2002).  Or the stars form from extragalactic gas and descend on a very radial orbit, penetrating the inner 15 kpc of the host halo from its very first pericentric passage.  An intriguing example is \\omegacen. Unfortunately our analytical model is not suited for this system because it is presently on a low-inclination eccentric retrograde orbits between 1 and 6 kpc from the Galactic center (Dinescu et al. 1999), so the contribution of dynamical friction by the disk is important.  Also hydrodynamical friction with the disk gas can play a role for an early-on partially gaseous satellite.  Nevertheless, if one applies simplisticly the tidal massloss and ChandraSekhar's dynamical friction in a spherical halo, one finds that while it seems easy to peel off a satellite galaxy to make a central star cluster, most simulations produce remnants on much larger orbits than \\omegacen (Zhao 2002). It seems some fine tuning is required to select progenitors on very low angular momentum and/or low energy orbits: the initial angular momentum needs to be low enough for the progenitor to penetrate into the inner halo or the present location of \\omegacen on its very first pericentric passage.  This means that the initial orbital size $S_i^\\omega$ of \\omegacen is in between the present value of \\omegacen $S_0 \\sim 1.25$kpc and the boundary of the inner halo $R_{\\rm disk}=15$kpc, or mathematically \\beq 1.25\\kpc < S_i^\\omega < 15\\kpc. \\eeq  Most recently there have been several very encouraging attempts to model the dynamical and star formation history of \\omegacen by nearly self-consistent N-body simulations (Mizutani et al. 2003, Tsuchiya et al. 2003, Bekki \\& Freeman 2003). All are able to produce both a reasonable mass and orbit of the \\omegacen after some trial and error with the initial parameters of the progenitor; many initial conditions lead to remnants, unlike \\omegacen, beyond 10kpc of the Milky Way center. The favored initial orbit has a small orbital size $S_i$. according to Tsuchiya et al. $j_i=60 {\\rm kpc} \\times 20 \\kms =1200{\\rm kpc}\\kms$ (or $S_i=6$kpc) and  according to Bekki \\& Freeman $j_i=25 {\\rm kpc} \\times 60 \\kms =1500{\\rm kpc}\\kms$ (or $S_i=7.5$kpc). The small orbital size seems consistent with our expectation (cf. Fig.3).  Tsuchiya et al. launch satellites with various initial mass $(0.4-1.6)\\times 10^{10}\\msun$ and with either a King profile or a Hernquist profile from 60 kpc from the Milky Way center. They choose well-aimed nearly radial orbits, with an initial perigalactic radius about 1 kpc, much more radial than the present eccentric orbit. Massloss in their King model are similar to our exponential massloss models ($n=\\infty$): rapid in the beginning, and $\\log(m)$ is roughly linear with time up to a mass of $10^8\\msun$ when the satellite has too little mass to proceed with the orbital decay. Massloss in the Hernquist model is closer to a $n=0.3$ model, linear in the beginning and rapid just before complete disruption (cf Fig.1a).  Our comparison with Tsuchiya et al.'s numerical model would be fair apart from one theoretical concern.  The progenitor in their best simulation is a two-component \"nucleated\" model with a rigid nucleus modeled by a extended-particle of $10^7\\msun$ with a half mass radius of 35pc on top of a live satellite of $0.8\\times 10^{10}\\msun$ with a Hernquist profile of half-mass radius $1.4$kpc; the dynamical friction of the Hernquist halo helps to deliver the nucleus eventually to an orbit similar to that of \\omegacen.  However, a closer examination reveals a subtle inconsistency in making the nucleus rigid:  the tidal force from the Hernquist halo beats the self-gravity of this fluffy nucleus at its half-mass radius by a factor of a few, so it could not have stayed and being moved as one piece. Nevertheless, one could have used a more contact, hence more plausible model of the rigid nucleus, say, with a total mass of $3\\times 10^6\\msun$ and a smaller half-mass radius of 7pc, which are closer to the observed mass and half-mass radius of \\omegacen. With this in mind, the formal inconsistency in Tsuchiya et al.'s best simulation seems to be harmless, and their model shows that \\omegacen could in principle be the remnant of a massive satellite on an orbit of initial apo-to-peri ratio of (about) 60kpc/2kpc."
    },
    "0403/astro-ph0403316_arXiv.txt": {
        "abstract": "We consider the stability of a configuration consisting of a vertical magnetic field in a planar flow on elliptical streamlines in ideal hydromagnetics. In the absence of a magnetic field the elliptical flow is universally unstable (the ``elliptical instability''). We find this universal instability persists in the presence of magnetic fields of arbitrary strength, although the growthrate decreases somewhat. We also find further instabilities due to the presence of the magnetic field. One of these, a destabilization of Alfven waves, requires the magnetic parameter to exceed a certain critical value. A second, involving a mixing of hydrodynamic and magnetic modes, occurs for all magnetic-field strengths. These instabilities may be important in tidally distorted or otherwise elliptical disks. A disk of finite thickness is stable if the magnetic fieldstrength exceeds a critical value, similar to the fieldstrength which suppresses the magnetorotational instability. ",
        "introduction": "The problem of momentum transport in accretion disks is widely believed to require hydrodynamic or hydromagnetic turbulence for its resolution. The origin of this turbulence may be sought in the instability of laminar solutions of the equations of hydromagnetics, solutions that are compatible with the geometry of accretion disks. The recent history of these efforts has taken the form of first recognizing such an instability mechanism, and then trying to incorporate that mechanicsm into realistic disk models.  The magnetorotational instability (MRI) mechanism, originally discovered by \\citet{vel} and \\citet{cha} and first applied to accretion disks in \\citet{bah}, is of this kind (see \\citet{mri} for a review). It appears in rotating, magnetized systems in which the specific angular momentum increases outward and in which the magnetic field is weak enough that rotational effects are not overwhelmed by magnetic tension.  A second mechanism, that of the elliptical instability considered by \\citet{jg} and others \\citep{lpk,ryug, ryugv}, is also consistent with the accretion-disk setting. This instability mechanism has been reviewed by \\citet{rk02}. In the setting considered by Goodman et. al., it appears to require a secondary in order to enforce departure from rotational symmetry of the streamlines via a tidal potential. This is certainly appropriate for binary systems but it is likely that, even in the absence of a secondary, the laminar motion in the plane of the disk would not be accurately circular, so the elliptical-instability mechanism would appear to be a candidate of considerable generality. It does not require a magnetic field. One of the conclusions of the present paper is that it further persists in the presence of a magnetic field. In the idealized setting of the present problem, the latter may be of arbitrarily large strength. However, we also argue that in the setting of a disk geometry, there may indeed be a limit on the field strength.  In this paper we therefore investigate the interaction of a vertical magnetic field with flow on elliptical streamlines, on the ground that both magnetic fields and noncircular streamlines are likely ingredients in accretion-disk settings. There are similarities with and differences from previous work on effect of magnetic fields on the elliptical instability \\citep{rk94}, which are discussed in \\S \\ref{discussion}. ",
        "conclusions": "\\label{discussion} We have explored the effect of a uniform, vertical magnetic field on the stability of planar, incompressible flow with elliptical streamlines in an unbounded medium, in the approximation of ideal magnetohydrodynamics. In the absence of magnetic fields, flows with elliptical streamlines having ellipticity parameter $\\epsilon$ [see equation (\\ref{equ:definitions})] are known to be unstable to perturbations with wavevectors that are transverse to the plane of the flow (the ``elliptical instability\"). Our first conclusion is that this elliptical instability persists in the presence of a vertical magnetic field: the latter decreases the maximum growthrate but fails to suppress the instability, no matter how large the magnetic-field parameter becomes. It can be compared with the conclusion of \\citet{rk94} that a toroidal magnetic field has a stabilizing influence. Kerswell's analysis holds for small $\\epsilon$ only and shows that the growthrate decreases with magnetic field in that limit. Our result, which holds for a vertical magnetic field, shows a similar trend for small values of the magnetic-field parameter, but this trend never results in complete stabilization of the elliptical instability with increasing magnetic field.  A second conclusion is that there are further instabilities associated with the presence of the magnetic field. One of these, for which the eigenvector is a mixture of hydrodynamic and magnetic modes, occurs for all values of the magnetic field parameter. Another, for which the eigenvector is a combination of magnetic modes only, sets in for values of the magnetic-field parameter exceeding a certain threshold value ($\\eta > \\sqrt{3}$). In all three of these instabilities, for large magnetic fields, the wave vector makes only a small angle with the plane of the unperturbed flow, reflecting the familiar tendency for dynamics to become nearly 2-dimensional in a strong, well ordered magnetic field. This is reflected in Figure \\ref{mixedmode}, which shows that as $\\eta$ increases, the unstable wedges are pushed to smaller $\\mu$. Although the unstable fraction of the $(E,\\mu)$ plane decreases with increasing $\\eta$ (except for a very slight maximum at $\\eta\\sim$ 2.18, reflecting the onset of instability between magnetic modes), the separation between the unstable wedges also decreases.  While the nonlinear evolution of the unstable system is beyond the scope of this work, the destabilization of a nearly continuous swath of parameter space may have consequences for the interactions between unstable modes. In all three cases, the maximum instability increment tends to $\\epsilon \\pi/2$, i.e., the maximum growthrate of the unstable modes tends, in dimensional units, to $\\epsilon\\Omega /4$.  All three of these instabilities may be relevant in accretion-disk settings. In systems of finite thickness $H$, however, the instability is suppressed if the Alfv\\'en speed $v_A$ exceeds a critical value of order $\\Omega H$. Magnetorotational instabilities are quenched at approximately the same fieldstrength \\citep{mri}. The growth rate of magnetoelliptical instabilities is smaller than that of magnetorotational instabilities by a factor of order $\\epsilon$, and thus they are not necessarily the primary instability in magnetized disks. They may well play a secondary role by breaking up eddies or vortices generated by other mechanisms.  Magnetoelliptic instabilities may also occur in the inner parts of barred galaxies, in which the gas flow is slightly elliptical and the magnetic field, at least in the Milky Way, has a vertical component \\citep{mor}. In such settings, the instabilities could be a source of turbulence, possibly affecting the mass supply to a central compact object."
    },
    "0403/astro-ph0403120_arXiv.txt": {
        "abstract": "We review the results obtained with the Galactic center campaigns of the BeppoSAX Wide Field X-ray Cameras (WFCs). This pertains to the study of luminous low-mass X-ray binaries (LMXBs). When pointed at the Galactic center, the WFC field of view contains more than half of the Galactic LMXB population. The results exemplify the excellent WFC capability to detect brief X-ray transients. Firstly, the WFCs expanded the known population of Galactic thermonuclear X-ray bursters by 50\\%. At least half of all LMXBs are now established to burst and, thus, to contain a neutron star as compact accretor rather than a black hole candidate. We provide a complete list of all 76 currently known bursters, including the new case 1RXS J170854.4-321857. Secondly, the WFCs have uncovered a population of weak transients with peak luminosities up to $\\sim10^{37}$~\\lum\\ and durations from days to weeks. One is the first accretion-powered millisecond pulsar SAX~J1808.4-3658. Thirdly, the WFCs contributed considerably towards establishing that nearly all (12 out of 13) luminous low-mass X-ray binaries in Galactic globular clusters contain neutron stars rather than black holes. Thus, the neutron star to black hole ratio in clusters differs from that in the Galactic disk at a marginal confidence level of 97\\%. ",
        "introduction": "Bleeker \\cite{bl03} reviewed in general terms the prospects and results of the BeppoSAX Wide Field Camera instrument package (``WFCs''). The unique capability of matching a wide field of view of $40^{\\rm o}\\times40^{\\rm o}$ per each of two identical cameras with a good angular resolution of 5\\arcmin\\ \\cite{jag97} not only led to a revolution of gamma-ray burst research (e.g., \\cite{pir03,hei03}) but also to a serious advance of our knowledge on X-ray bursts and other transient emission processes in low-mass X-ray binaries (LMXBs). Much of the relevant data was acquired during the only dedicated BeppoSAX observation program involving the WFCs as prime instrument. This encompasses observations pointed at the Galactic center during semi-yearly visibility windows that lasted from mid February to mid April and from mid August to mid October. Table~\\ref{tab1} summarizes the twelve campaigns during the six-year BeppoSAX lifetime. The combined exposure time represents 8\\% of the total BeppoSAX exposure budget. The success of this program may be anticipated by the mere fact that the field of view of one WFC emcompasses more than half the Galactic LMXB population according to pre-WFC catalogs \\cite{par95}. Another important ingredient for the success is that nearly all observations were analyzed in a near to real-time fashion. This was possible thanks to the 24 hr per day, 7 days a week, manning of the BeppoSAX Science Operations center (which was also crucial to the success of the GRB program of BeppoSAX \\cite{pir03}) and the dedicated support by the 'duty scientists'.  Tied to this program a dedicated target-of-opportunity program was in place for the BeppoSAX Narrow Field Instruments (NFIs) to follow up new transients or bursters. This program was triggered 14 times, with exposure times between 20 and 40 ksec. Whenever new transients or bursters were discovered, these were announced in IAU circulars. This happened on 20 occasions and triggered independent TOO programs on BeppoSAX, the Rossi X-ray Timing Explorer (RXTE), XMM-Newton, and ground-based radio and optical telescopes, illustrating a community service of the WFC program.  \\begin{table}[!t] \\caption{WFC observation campaigns on the Galactic center. The effective exposure times are for the position of the Galactic center. Adapted from \\cite{ha03,we03}.\\label{tab1}} \\begin{tabular}{ccrr} \\hline \\multicolumn{2}{c}{Campaign} & \\multicolumn{1}{r}{\\# obs.} & \\multicolumn{1}{r}{$t_{\\mathrm{exp}}\\,(\\rm{ks})$} \\\\ \\hline 1996 & Aug.15--Oct.29 & 67 & 1017   \\\\ 1997 & Mar.02--Apr.26 & 21 &  654   \\\\ 1997 & Sep.06--Oct.12 & 13 &  302   \\\\ 1998 & Feb.11--Apr.11 & 17 &  551   \\\\ 1998 & Aug.22--Oct.23 & 10 &  410   \\\\ 1999 & Feb.14--Apr.11 & 14 &  470   \\\\ 1999 & Aug.24--Oct.17 & 24 &  801   \\\\ 2000 & Feb.18--Apr.07 & 21 &  633   \\\\ 2000 & Aug.22--Oct.16 & 29 &  767   \\\\ 2001 & Feb.14--Apr.23 &  5 &  215   \\\\ 2001 & Sep.04--Sep.30 &  7 &  284   \\\\ 2002 & Mar.05--Apr.15 &  5 &   91   \\\\ \\hline \\end{tabular} \\end{table}  In the present paper, we provide a general overview of the accomplishments obtained through the WFC Galactic center program. We categorize the achievements in three areas: thermonuclear X-ray bursts, transients, and luminous globular cluster sources. ",
        "conclusions": "The results of the Galactic center campaigns of the BeppoSAX WFCs are plentiful. This is mainly due to its large field of view which enabled to monitor a large fraction of the Galactic LMXB population and resulted in a large exposure of roughly 7 million seconds over six years on tens of sources.  Such an exposure has not been accomplished by any other device yet. Thus, it has been possible to detect rare phenomena such as superbursts, burst-only sources and swift transients (e.g., V4641 Sgr).  The analysis of WFC data will continue and new results are expected to come out. This particularly concerns time scales of a few hours and the longest time scales, and weak signals (e.g., X-ray bursts fainter than 0.5~Crab)."
    },
    "0403/astro-ph0403250_arXiv.txt": {
        "abstract": "The Chandra Multiwavelength Project (ChaMP) has discovered a jet-like structure associated with a newly recognized QSO at redshift z=1.866. The system was 9.4\\arcmin\\ off-axis during an observation of 3C 207. Although significantly distorted by the mirror PSF, we use both a raytrace and a nearby bright point source to show that the X-ray image must arise from some combination of point and extended sources, or else from a minimum of three distinct point sources.  We favor the former situation, as three \\emph{unrelated} sources would have a small probability of occurring by chance in such a close alignment. We show that interpretation as a jet emitting X-rays via inverse Compton (IC) scattering on the cosmic microwave background (CMB) is plausible. This would be a surprising and unique discovery of a radio-quiet QSO with an X-ray jet, since we have obtained upper limits of 100 $\\mu$Jy on the QSO emission at 8.46 GHz, and limits of 200 $\\mu$Jy for emission from the putative jet. ",
        "introduction": "The objectives of the \\axaf\\ Multiwavelength Project (ChaMP) include identification and categorization of a complete, well-defined sample of serendipitous sources \\citep{Kim03, Green03}. The results will be of use, e.g., to study luminosity functions and their evolution, to quantify the newly resolved source(s) of the hard diffuse X-ray background, and to study cosmic structure and clustering of AGN and galaxies. The wide angle nature of this survey also makes it ideal to discover rare and unusual objects suitable for detailed study; e.g., lensed QSOs and X-ray jets.   Schwartz (2002a,b) has pointed out that if the jets observed in X-rays on scales of tens to hundreds of kpc are emitting via IC scattering of the CMB as suggested by \\citet{Tavecchio00} and \\citet{Celotti01}, then they will maintain the same apparent surface brightness independent of redshift, and therefore can be detected to arbitrarily large redshifts, up to the epoch at which they form.  The \\emph{Chandra} observations of such large scale jets in QSOs and powerful FR II radio sources are typically interpreted as IC/CMB emission, \\citep{Schwartz00, Harris02, Marshall01,Sambruna01,Siemiginowska02}.  All such interpretations require the assumption that the jet is either relativistically beamed with Doppler factors of order $\\delta \\sim$ 3 to 15, or that the energy density in relativistic electrons grossly exceeds the magnetic field energy density by at least two orders of magnitude. Detection of the X-ray ``beacons'' predicted by Schwartz (2002a,b) would provide additional evidence that the above assumptions are well founded.  We report the discovery of a candidate for such a system: CXOMP J084128.3+131107, (hereafter called J0841). The X-ray image shows an elongated structure. Despite the broad point response function (PSF) of the \\axaf\\ telescope at this 9.4\\arcmin\\ off-axis angle, we show that at least three point sources would be required to simulate the observed extent. We favor an interpretation of emission from the jet of an optically identified QSO which is close to the peak X-ray intensity.  We also mention alternate interpretations. Due to the small probability for three \\emph{unrelated} sources to occur by chance in this configuration, such interpretations may be even more unusual. ",
        "conclusions": "\\citet{Schwartz02} has noted that X-ray emission by IC/CMB should result in X-ray jets being cosmic beacons -- maintaining the same surface brightness at any larger redshift.  This is because the (1+z)$^{-4}$ cosmic diminution of surface brightness is exactly compensated by the (1+z)$^{4}$ increase in the energy density of the CMB with redshift. Such an effect does not depend on equipartition, or on relativistic beaming.  The low magnetic field, $\\le$2 $\\mu$G, implied by the limits to radio emission is unusual. Fields in clusters of galaxies can approach 1 $\\mu$Gauss, while typical jet fields on kpc scales are of order 10 $\\mu$Gauss. So the upper limits to magnetic field strengths derived here are somewhat weak for a jet.  However, there seems to be no fundamental physics prohibiting massive black holes to produce jets of such low internal energy density.  Selection bias against finding radio quiet X-ray jets could explain why such low magnetic field jets have not previously been noticed. Alternately, this object may have a magnetic field much weaker than the equipartition value."
    },
    "0403/astro-ph0403299_arXiv.txt": {
        "abstract": "Using evolutionary population synthesis (EPS) we present integrated spectral energy distributions (ISEDs) and absorption-line indices defined by the Lick Observatory image dissector scanner (referred to as Lick/IDS) system, for an extensive set of instantaneous burst single stellar populations (SSPs). The ages of the SSPs are in the range $1 \\leq \\tau /{\\rm Gyr} \\leq 19$ and the metallicities $-2.3 \\leq {\\rm [Fe/H]} \\leq +0.2$. Our models use the rapid single stellar evolution (SSE) algorithm of \\citet*{hur2000} for the stellar evolutionary tracks, the empirical and semi-empirical calibrated BaSeL-2.0 model of \\citet*{lej97b,lej98} for the library of stellar spectra and empirical fitting functions of \\citet{wor94b} for the Lick/IDS spectral absorption feature indices.  Applying our synthetic Lick/IDS absorption-line indices to the merit function we obtain the age and the metallicity of the central region of M32, it can be well interpreted with an instantaneous SSP with an age of $\\sim$ 6.5\\,Gyr and a metallicity similar to solar. Applying the derived age and the metallicity from the merit function to a number of index-index diagrams, we find that the plots of $H_\\beta - Fe5015$ and $H_\\beta - Fe5782$ are the best index-index diagrams from which we can directly obtain reasonable age and metallicity. ",
        "introduction": "Age effects often mask metallicity effects in the studies of stellar populations \\citep{oco76,oco86,oco94,wor92}, and separating them is a cumbersome affair \\citep*{ren86,buz92,buz93}. The age-metallicity degeneration originates from the fact that increasing either the metallicity or the age makes the integrated spectral energy distribution (ISED) of a single stellar population (SSP) redder \\citep*{bre96}.   Previous studies showed that it is difficult to break this degeneration only by broad-band colours \\citep{ari96,wor94a}. In order to solve this question, some studies used spectral information, instead of colours only, in the evolutionary population synthesis (EPS) models \\citep*{bre94,bru93,bru96,bru2003,buz89,kod97,tan96,vaz99a,wor94a}. With the development of these models including spectral information, the spectral resolution has been improved from about 20\\,\\AA \\ to 2\\,\\AA \\ \\citep{vaz99a}.   Furthermore, these spectral information has been translated to the line strengths either by empirical fitting polynomials \\citep{gor93,wor94b,wor97} or approaches other than the Lick fitting functions \\citep{pel89,buz95}. The inclusion of absorption feature indices in EPS models could add the power of diagnostics to the study of stellar populations \\citep{wor94a}, therefore some studies have included spectral absorption feature indices into EPS models attempting to understand the stellar populations. Some earlier EPS studies contained spectral indices include no metallicity dependence \\citep{bru83,tin72a,tin72b,tin76} or are otherwise limited in scope \\citep*{aar78,fro80,mou78,tri92}. Recently, the spectral absorption indices have been combined systematically in EPS studies, in those models some model builders use the absorption feature indices at intermediate resolution (9\\,\\AA) of Lick system \\citep*{bru96,bru2003,jon95,kur99,pel89,vaz96,vaz99a,wor94a}, which has been accepted by many investigators and widely used in their studies; other model builders use the indices at considerably higher resolution (FWHM $\\sim$ 2\\,\\AA) in Rose system \\citep{buz92,buz93,gor93,ros94,vaz99a,vaz2001}.   SSPs are assemblies of chemically homogeneous and coeval single stars, the star formation history of any stellar system can be described by a superposition of SSP models with different ages and metallicities. Studying SSPs can help us to understand the evolution of clusters and galaxies, the distribution of metallicities, and to quantify the star formation history of the galaxies. In this paper we also investigate SSPs with a systematic and self-consistent set of stellar evolution models by \\citet{pol98}, spectral library -- the BaSeL-2.0 model \\citep{lej97a,lej97b,lej98} and the spectral absorption feature indices of Lick system \\citep{gor93,wor94b}.   The outline of the paper is as follows. In Section 2, we describe our SSPs models and algorithm. In Section 3, we give the results and discussion. In Section 4, we apply the observed spectral indices to the merit function to determine the age and the metallicity for SSP-like assembly. Finally we present the summary and conclusions in Section 5. ",
        "conclusions": "In this paper we use the EPS technique to present the ISEDs at intermediate resolution (10\\,\\AA \\ in the ultraviolet and 20\\,\\AA \\ in the visible) and the Lick/IDS spectral absorption indices for an extensive set of instantaneous burst SSPs of various ages and metallicities ($1 \\leq \\tau \\leq 19\\,$Gyr, $-2.3 \\leq {\\rm [Fe/H]} \\leq 0.2$). In our EPS models we adopted the rapid SSE algorithm, the empirical and semi-empirical calibrated BaSeL-2.0 spectral library and the fitting functions for the absorption-line indices.   As checks on our models, we compare our ISEDs with those of Worthey and the high resolution spectrum of NGC 6838, also compare the Lick/IDS absorption feature indices with the values of other EPS models and those of Galactic and M31 globular clusters and the central region of M32. We conclude that: (a) The disagreement between our ISEDs with those of Worthey presents in two wavelength ranges, one in the far-ultraviolet region, the other in the visible and infra-red regions. The discrepancy of ISEDs in the far-ultraviolet region is mainly caused by the inclusion of hotter PN stars in our models, while the discrepancy of ISEDs in the visible and infra-red regions is caused by the choice of stellar evolutionary models and spectral library. The inclusion of cooler TPAGB/PPN stars can produce a significant deviation in ISEDs for young SSPs, but almost no effect for old SSPs in the visible and infra-red regions. (b) The integrated spectrum of NGC 6838 can be fitted by a SSP with solar metallicity and an age of $\\sim$ 12.6\\,Gyr. (c) When comparing our synthetic Lick/IDS absorption-line indices with the values in the literatures and observational data, we find that our models are generally in agreement with other models and match real globular clusters in most index-index diagrams. $TiO_{\\rm 1}$ is, however, systematically smaller.   At last using the merit function $F$, we find the that spectral indices of the central region of M32 can be fitted with a single stellar population with an age of $\\sim$ 6.5\\,Gyr and a solar-like metallicity with $F=1.00$. Applying the age and the metallicity of M32 to some index-index diagrams, we see that $H_\\beta - Fe5015$ and $H_\\beta - Fe5782$ are better diagrams to estimate directly the age and metallicity from index-index plots."
    },
    "0403/astro-ph0403585_arXiv.txt": {
        "abstract": "We analyse near-infrared {\\em HST}/NICMOS $F110W (J)$ and $F160W (H)$ band photometry of a sample of 27 $i'$-drop candidate $z\\simeq6$ galaxies in the central region of the HST/ACS Ultra Deep Field (HUDF).  The infrared colours of the 20 objects not affected by near neighbours are consistent with a high redshift interpretation. This suggests that the low redshift contamination of this $i'$-drop sample is smaller than that observed at brighter magnitudes where values of 10-40\\% have been reported.  The $J-H$ colours are consistent with a slope flat in $f_\\nu$ ($f_\\lambda\\propto\\lambda^{-2}$), as would be expected for an unreddened starburst.  There is, however, evidence for a marginally bluer spectral slope ($f_\\lambda\\propto\\lambda^{-2.2}$) which is perhaps indicative of an extremely young starburst ($\\sim10$ Myr old) or a top heavy initial mass function and little dust. The low levels of contamination, median photometric redshift of $z\\sim6.0$ and blue spectral slope, inferred using the near-infrared data, supports the validity of the assumptions in our earlier work in estimating the star formation rates and, that the majority of the i-drop candidates galaxies lie at $z\\sim6$. ",
        "introduction": "\\label{sec:intro}  In recent years the observational horizon has expanded rapidly and radically for those observing distant galaxies.  Large format red-sensitive detectors on wide field imaging instruments, the new generation of 8m class telescopes and the refurbished Hubble Space Telescope ({\\em HST}), have pushed the limits to which we can routinely detect star-forming distant galaxies progressively from redshifts of one to beyond $z=6$.  At the highest redshifts currently accessible, narrow band emission lines searches using the Lyman-$\\alpha$ line have moved on from redshifts of 4 \\citep{hm96,fu03}, to 5.7 \\citep{hu99} and now reach to $z\\sim6.5$ \\citep{ko03,hu02}.  Photometric redshifts \\citep[e.g.][]{lan96,co97,fe99} are now routinely used to process large datasets and identify high redshift candidates. An application of this method, the continuum based Lyman-break photometric technique pioneered at $z\\sim3$ by \\citep{guh90,st95}, has been extended progressively to z$\\sim4$ and $z\\sim5$ \\citep[e.g.][]{st99,br04}, spectroscopically confirmed, and using $i'$-band drop selection further extended to $z\\sim6$ \\citep{st03,st04a,st04b,bo03a,bb03,bo04,di04,yan03}.  The `classical' Lyman break technique used by Steidel et al.\\ (1996) uses three filters, one redward of rest frame Lyman-$\\alpha$ ($\\lambda_\\mathrm{rest}>1216$\\AA), a second in the spectral region between rest-frame Lyman-$\\alpha$ and the rest-frame Lyman limit (912\\AA) and a third at $\\lambda_\\mathrm{rest}<912$\\AA. At z$\\sim$3 the technique relies on the ubiquitous step in the spectra of the stellar component of galaxies at 912\\AA\\ due to photospheric absorption, supplemented by optically-thick Lyman limit system absorption caused by neutral hydrogen in the galaxy in question or in the intervening IGM.  At higher redshifts the evolution in the Lyman-$\\alpha$ forest absorption, particularly in the spectral region 912-1216\\AA\\ means that the effective break migrates redward to the Lyman-$\\alpha$ region.  Recently this search technique has been extended to $z\\sim6$ by various authors and $i'$-drop samples have been used to constrain the star formation history of the universe (as derived from rest-frame UV luminous, starbursting stellar populations, e.g. Stanway, Bunker \\& McMahon 2003; Giavalisco et al.\\ 2004; Bouwens et al.\\ 2004; Bunker et al.\\ 2004).  Any such estimation, however, is dependent on the characteristic properties of the $i'$-drop population. Although Steidel and coworkers \\citep{st99,st01} have successfully shown that a Lyman-break technique cleanly selects galaxies at intermediate redshift ($z\\approx3-4$), the $i'$-drop varient of the technique lacks a second colour to constrain the redshift distribution and is subject to contamination from both low redshift elliptical galaxies and cool dwarf stars.  Until recently, these factors have been essentially unconstrained by observational data.  The {\\em HST}/ACS Ultra Deep Field (HUDF, Beckwith, Somerville and Stiavelli 2004), imaging an 11 arcmin$^2$ region of the sky to faint magnitudes in the $F435W(B)$, $F606W(v)$, $F775W(i')$ and $F850LP(z')$ bands, has now allowed a clean sample of $i'$-drop objects to be defined and constraints to be placed upon its luminosity function \\citep{b04}. The {\\em HST}/NICMOS treasury programme, complimentary to the HUDF, has imaged the central region of the ACS field to faint magnitudes at wavelengths of 1.1 and 1.6 microns, allowing the infrared properties of the $i'$-drop population (in particular their rest-frame ultra-violet spectral slope, faint-end contamination and luminosity at 1500\\AA) to be determined. In this paper we discuss the infrared properties of the $i'$-drop population defined in \\citet{b04}, and their implications for the nature of these objects.  We adopt a $\\Lambda$-dominated, `concordance' cosmology with $\\Omega_{\\Lambda}=0.7$, $\\Omega_{M}=0.3$ and $H_{0}=70\\,h_{70} {\\rm km\\,s}^{-1}\\,{\\rm Mpc}^{-1}$. All magnitudes in this paper are quoted in the AB system \\citep{og83} and the \\citet{ma95} prescription, extended to $z=7$, is used where necessary to estimate absorption due to the intergalactic medium. Beyond $z=7$ absorption is assumed to decrease the transmitted flux by 98\\% \\citep[the decrement observed in the flux of $z>6$ quasars identified by the Sloan Digital Sky Survey, ][]{fa03} . ",
        "conclusions": "\\label{sec:conc}  \\begin{enumerate} \\item The near-infrared properties of a faint sample of $i'$-drop objects (defined by Bunker et al.\\ 2004) has been investigated, using the deep $F110W$ and $F160W$ images of the HST/ACS Ultra Deep Field.  \\item The infrared colours of all objects not suffering from contamination by near neighbours are consistent with a high redshift interpretation. This suggests that the contamination of this $i'$-drop sample is smaller than that observed at brighter magnitudes.  \\item The spectral slopes inferred from near infrared colours are consistent with $\\beta=-2.0$ although there is evidence for a marginally steeper spectral slope of $\\beta=-2.2\\pm0.2$.  \\item These steep spectral slopes suggest that the dust extinction of the sample is small and that the $i'$-drop population may comprise galaxies with young starbursts - possibly triggered by mergers (several close systems of galaxies are observed in this sample).  \\item The low levels of contamination, and the steep spectral slope, inferred from near-infrared data supports the validity of our previous star formation density estimates based on $i'$-drop samples. \\end{enumerate}  \\subsection*{Acknowledgements}  Based on observations made with the NASA/ESA Hubble Space Telescope, obtained from the Data Archive at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS 5-26555. These observations are associated with programs \\#9803 and \\#9978. We thank Paul Hewett, Richard Ellis and Malcolm Bremer for useful discussions. ERS acknowledges a Particle Physics and Astronomy Research Council (PPARC) studentship supporting this study.  We also thank Steve Beckwith and the HUDF team for making these high quality images publically available."
    },
    "0403/astro-ph0403066_arXiv.txt": {
        "abstract": "In this paper we present a multiresolution-based method for period determination that is able to deal with unevenly sampled data. This method allows us to detect superimposed periodic signals with lower signal-to-noise ratios than in classical methods. This multiresolution-based method is a generalization of the wavelet-based method for period detection that is unable to deal with unevenly sampled data, presented by the authors in \\cite{otazu}. This new method is a useful tool for the analysis of real data and, in particular, for the analysis of astronomical data. ",
        "introduction": "\\label{intro}  The detection of periodic signals in astronomical data has been usually addressed by classical Fourier-based or epoch folding methods. These methods have different problems when dealing with non-sinusoidal periodic signals or with very low signal-to-noise ratios. When the analysed data set contains several periodic signals, the behavior of classical period determination methods highly depends on intrinsic signal characteristics (see, for example, \\citealt{lafler}; \\citealt{jurkevich}; \\citealt{stellingwerf}; \\citealt{lomb}; \\citealt{scargleA}, \\citeyear{scargleB}; \\citealt*{roberts}; \\citealt{press}). The reader is referred to the introduction of \\cite{otazu} (hereafter Paper~I) for a general discussion. To avoid this problem, we presented there, as a preliminary step, a wavelet-based approach that only works with evenly time sampled data.  In astronomy, however, this is not a usual situation, since data are mostly acquired on irregular intervals of time. In such a case there are two possibilities: resample the data into a new evenly sampled data set, or use a method able to deal with the original unevenly sampled data set. In the first case we are forced to modify the original data, which necessarily implies a loss of information. Moreover, this is not always possible if the temporal gaps are larger than some of the periods present in the data. In order to avoid these problems, a technique capable to deal with unevenly sampled data is needed.   The present paper is a natural extension of Paper~I that allows to work on unevenly time sampled data. We show how the methodology of multiresolution decomposition (similar to the wavelet decomposition philosophy) is very well suited to this problem, since it is completely oriented towards decomposing functions into several frequential characteristics.  As in Paper~I, the main objective is to isolate every signal present in our data and to analyse them separately, avoiding their mutual influences. In Section~\\ref{multiresolution} we outline some concepts in multiresolution analysis and their similarities with wavelet theory that are relevant to the stated problem. In Section~\\ref{period} we propose an algorithm to detect each of the periodic signals present in a data set by combining multiresolution analysis decomposition with classical period determination methods. In Sections~\\ref{simulated} and \\ref{results} we present some examples of synthetic data we used to test the algorithm and the results we obtained. We summarise our conclusions in Section~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  In this paper we have presented a multiresolution-based method for period determination able to deal with unevenly sampled data. This constitutes a significant improvement with respect to the wavelet-based method presented in Paper~I, which is unable to deal with unevenly sampled data. The overall performance of the present method is similar to the wavelet-based one, in the sense that it allows us to detect superimposed periodic signals with lower signal-to-noise ratios than in classical methods. We stress that one advantage of the present method over classical methods is the simultaneous detection of a period in more than one multiresolution plane, allowing to improve the confidence of a given detection. Moreover, since here we are not forced to lose or modify the information when averaging or interpolating the original data, we can reach higher noise-to-signal ratios than in the wavelet-based method described in Paper~I.  We note that the multiresolution decomposition scheme that we have used can be interpreted as a particular case of scale-space filtering. In order to improve isolation of periodic features, more general approaches could be used to perform this decomposition. In this context, anisotropic diffusion schemes proposed by \\cite{perona} could be useful if properly tuned."
    },
    "0403/hep-ph0403003_arXiv.txt": {
        "abstract": "In a model of TeV right-handed (RH) neutrino by Krauss, Nasri, and Trodden, the sub-eV scale neutrino masses are generated via a 3-loop diagram with the vanishing see-saw mass forbidden by a discrete symmetry, and the TeV mass RH neutrino is simultaneously a novel candidate for the cold dark matter. However, we show that with a single RH neutrino it is not possible to generate two mass-square differences as required by the oscillation data. We extend the model by introducing one more TeV RH neutrino and show that it is possible to satisfy the oscillation pattern within the modified model. After studying in detail the constraints coming from the dark matter, lepton flavor violation and the muon anomalous magnetic moment, and the neutrinoless double beta decay, we explore the parameter space and derive predictions of the model. Finally, we study the production and decay signatures of the TeV RH neutrinos at TeV $e^+ e^-/\\mu^+ \\mu^-$ colliders. ",
        "introduction": "One of the most natural way to generate a small neutrino mass is via the see-saw mechanism \\cite{see-saw}.  There are very heavy right-handed neutrinos, which are gauge singlets of the standard model (SM), and so they could have a large majorana mass $M_R$.  After electroweak symmetry breaking, a Dirac mass term $M_D$ between the right-handed and the left-handed neutrinos can be developed. Therefore, after diagonalizing the neutrino mass matrix, a small majorana mass $\\sim m_D^2/M_R$ for the left-handed neutrino is obtained.  This is a very natural mechanism, provided that $M_R \\sim 10^{11-13}$ GeV. One drawback of this scheme is that these right-handed neutrinos are too heavy to be produced at any terrestrial experiments.  Therefore, phenomenologically there are not many channels to test the mechanism. Although it could be possible to get some hints from the neutrino masses and mixing, it is rather difficult to reconstruct the parameters of the right-handed neutrinos using the low energy data \\cite{kim}.  Another natural way to generate a small neutrino mass is via higher loop processes, e.g., Zee model \\cite{zee}, with some lepton number violating couplings.  However, these lepton number violating couplings are also subject to experimental constraints, e.g., $\\mu \\to e \\gamma, \\tau \\to e\\gamma$. In the Zee model, there are also extra scalars whose masses are of electroweak scale, and so can be observed at colliders \\cite{gl}.  On the other hand, recent cosmological observations have established the concordance cosmological model where the present energy density consists of about $73\\%$ of cosmological constant (dark energy), $23\\%$ (non-baryonic) cold dark matter, and just $4\\%$ of baryons. To clarify the identity of the dark matter remains a prime open problem in cosmology and particle physics. Although quite a number of promising candidates have been proposed and investigated in detail, other possibilities can never be neglected.  Recently, Krauss, Narsi, and Trodden \\cite{krauss} considered an extension to the SM, similar to the Zee model, with two additional charged scalar singlets and a TeV right-handed neutrino.  They showed that with an additional discrete symmetry the Dirac mass term between the left-handed and right-handed neutrinos are forbidden and thus avoiding the see-saw mass.  Furthermore, the neutrino mass can only be generated at three loop level, and sub-eV neutrino masses can be obtained with the masses of the charged scalars and the right-handed neutrino of order of TeV. Phenomenologically, this model is interesting because the TeV right-handed neutrino can be produced at colliders and could be a dark matter candidate.  In this work, we explore in details the phenomenology of the TeV right-handed (RH) neutrinos. We shall extend the analysis to three families of left-handed neutrinos and explore the region of the parameters that can accommodate the present oscillation data. In the course of our study, we found that the model in Ref. \\cite{krauss} with a single RH neutrino cannot explain the oscillation data, because it only gives one mass-square difference.  We extend the model by adding another TeV RH neutrino, which is slightly heavier than the first one.  We demonstrate that it is possible to accommodate the oscillation pattern. We also obtain the relic density of the RH neutrino, and discuss the possibility of detecting them if they form a substantial fraction of the dark matter.  We also study the lepton number violating processes and the muon anomalous magnetic moment, and the production at leptonic colliders. In particular, the pair production of $N_1 N_2, N_2 N_2$ at $e^+ e^-/\\mu^+\\mu^-$ colliders gives rise to very interesting signatures. The $N_2$ so produced will decay into $N_1$ plus a pair of charged leptons inside the detector.  Thus, the signature would be either one or two pairs of charged leptons plus a large missing energy.  The organization is as follows.  We describe the model in the next section. In Sec. III, we explore all the phenomenology associated with the TeV RH neutrino. In Sec. IV, we discuss the signatures in collider experiments. Section V is devoted to the conclusion. ",
        "conclusions": "In this paper, we have discussed a model that explains the small neutrino mass and dark matter in the universe at the same time. Such a model was proposed by Krauss {\\it et al.} as a modification of Zee model. However, our study revealed that their original model is unfortunately not capable of explaining the neutrino oscillation pattern.  We have extended the model by introducing another right-handed neutrino.  We succeed in showing that such an extension is possible to achieve the correct neutrino mixing pattern.  A prediction of this model is the normal mass hierarchy. In addition, the undiscovered mixing angle $\\theta_{13}$ is relatively large, because of the requirement of a mild cancellation between the parameters for a small $\\theta_{13}$ and a sensible coupling of the charged scalar, $\\lambda_s$.  The relic density of the lightest right-handed neutrino has also been revisited. Under the constraint by WMAP we found that the mass of the right-handed neutrino cannot be as large as TeV but only of order $1\\times 10^2$ GeV, after a careful treatment of the calculation.  In addition, other constrains including the muon anomalous magnetic moment, radiative decay of muon, neutrinoless double beta decay have also been studied.  With all the constraints we are still able to find a sensible region of parameter space.  Finally, our improved model has an interesting signature at leptonic colliders via pair production of right-handed neutrinos, in particular $N_1 N_2$ and $N_2 N_2$.  The $N_2$ so produced will decay into $N_1$ plus two charged leptons.  Thus, the signature is either one or two pairs of charged leptons with a large missing energy. Hence, this model can be tested not only by neutrino experiments but also by collider experiments."
    },
    "0403/astro-ph0403700_arXiv.txt": {
        "abstract": " ",
        "introduction": "At the heart of a successful theory of galaxy formation must be a detailed physical understanding of the dissipational processes which form spiral galaxies.  The disk is the defining stellar component of disk galaxies, and understanding its formation is in our view the most important goal of galaxy formation theory.  Although much of the information about the pre-disk state of the baryons has been lost in the dissipative process, some tracers are likely to remain.  What do we mean by the reconstruction of early galactic history? We seek a detailed physical understanding of the sequence of events which led to the Milky Way. Ideally, we would want to associate components of the Galaxy to elements of the protocloud -- the baryon reservoir which fueled the stars in the Galaxy.  For many halo stars, and some outer bulge stars, this may be possible with phase space information. But for much of the bulge and the disk, secular processes cause the populations to become relaxed (i.e.  the integrals of motion are partially randomized).  In order to have any chance of unravelling disk formation, we must explore chemical signatures in the stellar spectrum.  The exponential thin disk, with a vertical scale height of about 300 pc, is the most conspicuous component in edge-on disk galaxies.  The thin disk is believed to be the end product of the quiescent dissipation of most of the baryons and contains almost all of the baryonic angular momentum. Many disk galaxies show a second fainter disk component with a longer scale height (typically about 1 kpc); this is known as the thick disk \\citep{sd00,db02}. The Milky Way has a significant thick disk \\citep{gr83}:  its surface brightness is about 25\\% of the thin disk's surface brightness, and its stars are significantly more metal poor than the stars of the thin disk.  The galactic thick disk is currently believed to arise from heating of the early stellar disk by one or more accretion events although other possible origins have been discussed \\citep{gwk89,gwj95}. It is composed of only old stars, with ages greater than 10 Gyr, equivalent to forming at a redshift of $z\\geq1$ \\citep{rw00}. Furthermore,  the thick disk is chemically distinct from the thin disk at low metallicity \\citep{kf98,fbl03}, in the sense that [$\\alpha$/Fe] is enhanced relative to the thin disk, although solar values are seen at higher metallicity. The thick disk may be one of the most significant components for studying signatures of galaxy formation because it presents a 'snap frozen' relic of the state of the (heated) early disk.  In the discussion which follows, we make reference to the KAOS project (see http://www.noao.edu/kaos) which proposes a highly multiplexed wide-field multi-object spectrograph for the Gemini Observatory. Full specifications for the proposed instrument can be found at the above web site. ",
        "conclusions": ""
    },
    "0403/astro-ph0403536_arXiv.txt": {
        "abstract": "We combine the latest observations of disk galaxy photometry and rotation curves at moderate redshift from the FORS Deep Field (FDF) with simple models of chemical enrichment. Our method describes the buildup of the stellar component through infall of gas and allows for gas and metal outflows. In this framework, we keep a minimum number of constraints and we search a large volume of parameter space, looking for the models which best reproduce the photometric observations in the observed redshift range ($0.5<z<1$). We find the star formation efficiency to correlate well with $\\vrot$  so that massive disks are more efficient in the formation of stars and have a smaller spread in stellar ages. This trend presents a break at around $\\vrot\\sim 140$~km~s$^{-1}$. Galaxies on either side of this threshold have significantly different age distributions. This break has been already suggested by several authors in connection with the contribution from either gravitational instabilities or supernova-driven turbulence to star formation. The gas infall timescale and gas outflows also present a correlation with galaxy mass, so that massive disks have shorter infall timescales and smaller outflow fractions. The model presented in this paper suggests massive disks have formation histories resembling those of early-type galaxies, with highly efficient and short-lived bursts, in contrast with low-mass disks, which have a more extended star formation history. The ages correlate well with galaxy mass or luminosity, and the predicted gas-phase metallicities are consistent with the observations of local and moderate redshift galaxies. One option to explain the observed shallow slope of the Tully-Fisher relation at intermediate redshift could be small episodes of star formation in low-mass disks. ",
        "introduction": "Quantifying the star formation history of galaxies constitutes one of the major unsolved issues towards a complete understanding of galaxy formation. Within the current paradigm of structure formation, disk galaxies represent the building blocks from which the whole zoo of galaxies is formed. A detailed analysis of galaxy formation histories is hindered by the complexity of star formation. By coupling simple assumptions about star formation and chemical enrichment to population synthesis models which track the evolution of well-defined stellar populations, one can constrain the possible scenarios of star formation in a semi-quantitative way. Furthermore, by exploring galaxies at high redshift one takes advantage of the lookback times which allow us to probe back into the past histories of these galaxies.  Disk galaxies have been  complicated  to analyze, mostly because of their ongoing star formation, dust and complex geometry which prevent us from getting a clear picture of their underlying stellar populations. Qualitatively, the photo-spectroscopic properties of disk galaxies such as our own Milky Way are compatible with a weak and continuous star formation rate (SFR) along with a stronger early episode of star formation whose amplitude correlates well with galaxy type, being stronger in disks of earlier types (e.g. Kennicutt 1998a).  Studies of disks at moderate and high redshift seem to yield a similar formation scenario. Lilly et al. (1998) observed a sample of $341$ galaxies selected from the CFRS and the LDSS redshift surveys, out to $z\\simlt 1$, finding bluer $U-V$ colours and stronger [OII] emission at moderate redshift, which suggests an increase in the SFR by a factor of about $3$ from local galaxies out to $z\\sim 0.7$. However, the size function was not found to change with redshift, and the larger changes are associated with the smaller galaxies (however, see Giallongo et al. 2000; Vogt 2000). This is indicative of earlier formation of massive disks, in agreement with the lower gas fractions found in these systems (e.g. Bell \\& de~Jong 2000). Boissier \\& Prantzos (2001) presented a detailed model of chemical enrichment in disk galaxies calibrated to the Milky Way (Boissier \\& Prantzos 1999). The application of this model to disk galaxies at moderate redshift showed that large disks have already completed their evolution by $z\\sim 1$, whereas low-mass disks undergo a later evolution, resulting in a steepening of the slope of the Tully-Fisher relation with redshift. A similar conclusion is reached by Ferreras \\& Silk (2001) when comparing sets of simple chemical enrichment models over a wide range of parameter space using the $z=0$ optical and NIR colour--$\\vrot$ relation as a constraint. The star formation efficiency was found to correlate with galaxy mass, a point which is questioned by Boissier et al. (2001). However, this controversy hinges on the degeneracy between star formation efficiency and infall timescale. Both groups agree on the need for a rapid buildup of the stellar component in massive disks, a property that can be achieved either by a high efficiency of star formation or by a shorter infall timescale.  The evolution of the Tully-Fisher relation with redshift is a powerful discriminator of star formation histories. Vogt (2000) presented a sample of $100$ faint disk galaxies in the range $0.2<z<1$ and no change was found in either shape or slope with respect to  local samples. A zero point offset $<0.2-0.3$~mag was found in the rest frame $B$ band. Milvang-Jensen et al. (2003) explored cluster/field differences in the TF relation at moderate redshift on a smaller sample of $8$ disk galaxies in cluster MS1054-03 ($z=0.83$) as well as $19$ field spirals in the $z=0.15-0.9$ range. At fixed rotation velocity, they found cluster spirals to be $\\sim 0.5-1$~mag brighter. However, their sample is too small to draw a firm conclusion on the evolution of the Tully-Fisher relation. In a larger sample of disk galaxies in 3 clusters at moderate redshift ($z\\sim 0.5$) Ziegler et al. (2003) did not find significant differences with respect to field disks.  \\begin{figure} \\includegraphics[width=3.5in]{tfz_f1.eps} \\caption{Colour-redshift plots of the sample of FDF moderate redshift disk galaxies used in this paper. The filled and empty dots represent galaxies with $\\log\\vrot >2.2$ and $\\leq 2.2$, respectively. The lines correspond to the colours of simple stellar populations from the latest models of Bruzual \\& Charlot (2003), for three ages (from top to bottom): 8, 1, and 0.1~Gyr. For each age two metallicities are considered: $Z_\\odot$ (solid) and $Z_\\odot /3$ (dashed). Notice the bluer colours -- which are most sensitive to young stars -- suggest young stellar ages, whereas the red/NIR colours imply older ages. } \\label{fig:CCDs} \\end{figure}  The work on $113$ disk galaxies in the range $0.1<z<1.0$ from the FORS Deep Field on which this work is based (B\\\"ohm et al. 2004) revealed a striking result, namely a shallow slope of the rest frame $B$ band Tully-Fisher relation with respect to local disks. Slowly rotating disks ($\\vrot\\simlt 100$~km~s$^{-1}$) at moderate redshift appear up to $2$~magnitudes brighter than their low-redshift counterparts. One should question whether the samples at low and moderate redshift correspond to a simple one-to-one mapping. Indeed, the sample of Lilly et al. (1998) presents a large fraction of bright and small disks at $z\\sim 1$ which do not have local counterparts.  In this paper, we combine photometric data from the FDF disk sample with a model which convolves galactic chemical evolution histories with stellar population synthesis in order to relate the photometric observations with a small set of phenomenological parameters which describe the major drivers of star formation in galaxies. We describe the observations in \\S2 and our model in \\S3. The results are discussed in \\S4 and the resulting ages and metallicities are presented in \\S5. Finally, in \\S6 we discuss our results and give the conclusions. Throughout this paper we use a $\\Lambda$CDM cosmology with $\\Omega_m = 0.3$ and $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$, as suggested by the analysis of the angular power spectrum of the Cosmic Microwave Background observed by WMAP (Spergel et al. 2003). For this cosmology, the age of the Universe at $z=0$ and $z=1$ corresponds to $13.5$ and $5.8$~Gyr, respectively. ",
        "conclusions": "A simple phenomenological description of the star formation history allows us to explore a large number of possible histories only to be constrained by the observations. In this paper we have used the moderate redshift sample of disk galaxies from the FORS Deep Field. Comparing the aperture photometry in several passbands between data and simple models of chemical enrichment across a wide spectral range ($U\\cdots K_s$) we have found an increase of the star formation efficiency -- as defined in equation~\\ref{eq:schmidt}  -- with rotation velocity. A clear threshold appears around $\\vrot\\sim 140$~km~s$^{-1}$ so that slowly rotating disks present uniformly low star formation efficiencies and younger stellar populations with a larger spread of stellar ages. More massive disks show a wide range of efficiencies. A similar conclusion was reached by Ferreras \\& Silk (2001) on the analysis of the local sample of UMa cluster disks from Verheijen (2001). The thorough analysis of observational data across a wide spectral range -- from the radio to the UV -- probing the stellar and gas component of a sample of 928 nearby disk galaxies led Gavazzi, Pierini \\& Boselli (1996) to conclude that, to first order, mass is the driving parameter of galaxy evolution. They suggested that a positive correlation between star formation efficiency and galaxy mass should be invoked to explain the observations. Furthermore, this trend is consistent with the analysis of Dalcanton et al. (2004) who also find a threshold in $\\vrot$ when exploring the thickness of dust lanes in disks. They relate this threshold to a tradeoff between the two main causes of turbulence in the ISM, namely supernova feedback and gravitational instabilities. The former is the dominant factor in low-mass disks and produces a higher velocity dispersion in the gas, which results in lower star formation efficiencies. In contrast, Dalcanton et al. suggest turbulence in massive disks to be mainly caused by gravitational instabilities. The lower velocity dispersion implied by these instabilities results in an increased efficiency of star formation. Along similar lines, Silk (2001) suggested a Schmidt-like star formation law: \\be \\psi\\sim\\Omega\\rho_g\\Big( \\frac{\\sigma_g}{v_{\\rm SN}}\\Big) , \\ee where $\\Omega$ is the disk angular velocity, $\\sigma_g$ is the cold gas velocity dispersion and $v_{\\rm SN}$ is a characteristic velocity defined from supernova feedback as the specific momentum injected by supernovae per unit star formation rate. In this framework, the results from figure~\\ref{fig:params} suggest a major change in the ratio $(\\sigma_g/v_{\\rm SN})$ at around $\\vrot\\sim 140$~km~s$^{-1}$  This bimodal trend has been explored in a large dataset of $\\sim 120,000$~galaxies drawn from the SDSS (Kauffmann et al. 2003). Low-mass galaxies were found to have younger stellar populations and low concentrations, with some of these systems having experienced recent star formation. At stellar masses above $\\sim 3\\times 10^{10} M_\\odot$, galaxies feature older populations and the high concentrations typical of spheroidal systems. In the light of our models, this spheroidal-like behaviour of large disks reveals the high star formation efficiency as expected in bulges or elliptical galaxies. Furthermore, a recent paper (Kannappan 2004) has explored the gas-to-star mass ratios using $u-K$ colour as a proxy. The sample of $\\sim 35,000$ galaxies was extracted from the SDSS and 2MASS databases and a similar bimodal trend was found, with a threshold around a similar stellar mass, i.e. $2-3\\times 10^{10}M_\\odot$. Our results confirm this bimodal trend in disk galaxies out to $z\\sim 1$.  \\begin{figure} \\includegraphics[width=3.5in]{tfz_f6.eps} \\caption{The Tully-Fisher relation at moderate redshift. The filled and hollow circles correspond to the $z<0.75$ and $z\\geq 0.75$ FDF subsamples, respectively. A few characteristic error bars in $\\log\\vrot$ are shown. The stars and triangles represent the ($z\\simgt 0.4$) data from Vogt et al. (1997) and the field disks from Milvang-Jensen et al. (2003), respectively. The solid line is a linear fit to the FDF data. The fit to the local sample of Pierce \\& Tully (1992) is also shown as a dashed line, with a shaded region representing the scatter. } \\label{fig:TFz} \\end{figure}  The star formation histories proposed in this paper for disk galaxies allow us to determine reliable evolution corrections and rest-frame luminosities. Figure~\\ref{fig:TFz} shows the Tully-Fisher relation with respect to rest-frame $M_B$. The model absolute luminosities are very similar to those already presented by Ziegler et al. (2002) and B\\\"ohm et al. (2004) for which the $k$ corrections were obtained using the population synthesis models of M\\\"oller et al. (2001). As a comparison, we also show (at $z\\simgt 0.4$) the field sample of Milvang-Jensen et al. (2003; triangles) as well as the sample from the DEEP Groth Strip Survey (Vogt et al. 1997; stars). The solid line is a fit to our data, whereas the local Tully-Fisher relation is shown as a dashed line and the scatter is represented by the shaded region (Pierce \\& Tully 1992). The FDF sample clearly gives a shallower Tully-Fisher relation with respect to the local sample. The small size and scatter in the sample of Milvang-Jensen et al. (2003) prevents us from making any estimate of their mismatch with respect to the local TFR. Vogt et al. (1997) give a result which is more compatible with the observations at $z=0$. However, their selection favours large disks which biases the sample towards early-type systems. B\\\"ohm et al. (2004) suggested this slope change could be caused  by a mass-dependent luminosity evolution. In this paper we do not attempt to relate the bright and slowly rotating disks from the FDF sample with local low-mass disk galaxies. In fact, a na\\\"\\i ve one-to-one mapping of local and high redshift Tully-Fisher relations would result in an anticorrelation between star formation efficiency and galaxy mass. We only use the aperture photometry of the FDF disk galaxies in order to constrain their star formation histories. The rotation velocities are then used to assign a ``size'' parameter to each galaxy. The evolution of the slowly rotating disks from our analysis to $z=0$ results in a population of bright disks which are not seen in local surveys. Hence, our results conclude that these objects do not have local counterparts, in agreement with Lilly et al. (1998) who find a population of small and bright disks at $z\\simgt 0.7$ which do not have counterparts in the local galaxy census. A simple estimate of the contribution from young stars can be done if we consider the combination of two simple stellar populations, an old one which represents the underlying population and a young one with a fractional {\\sl mass} contribution $f_Y$ which corresponds to the recent episode of star formation. We can then relate the actual luminosity of the galaxy ($L$) to the luminosity of the underlying old stellar component ($L_O$), namely \\be \\frac{L}{L_O}=\\frac{M\\Upsilon^{-1}}{M_O\\Upsilon_O^{-1}}= 1 + \\Big( \\frac{f_Y}{1-f_Y}\\Big)\\frac{\\Upsilon_O}{\\Upsilon_Y}, \\ee where $\\Upsilon =M/L$ is the mass-to-light ratio, and the subscripts $O$ and $Y$ denote the old and young component, respectively. Let us consider a small fractional contribution, say $f_Y=0.1$. Assuming a metallicity $\\log (Z/Z_\\odot )=-0.5$ for both components, and an age of $6$~Gyr for the old population -- corresponding to a formation redshift $z_F\\sim 5$ for a galaxy at $z=0.75$ -- this 10\\% contribution in mass from young stars translates into $\\Delta M_B = M_B - M_{B,O} = 2.5\\log (L_{B,O}/L_B) = -1.9$~mag for a young component with $t_Y=0.1$~Gyr and $-0.95$~mag for $t_Y=0.5$~Myr. A couple of Gyr after the episode has subsided (i.e. $t_O=8$~Gyr and $t_Y=2$~Gyr) the same 10\\% contribution implies $\\Delta M_B=-0.36$~mag. Hence, small episodes of star formation could account for the presence of these bright and slowly rotating disks. We should mention that no strong evolution is expected in $\\vrot$ from mass accretion from $z=1$ to the present time. In our adopted $\\Lambda$CDM cosmology accretion is mostly suppressed at low redshift by the growth factor in linear theory. Indeed, the disk galaxy models of van~den~Bosch (2002) show that the rotation curves do not change more than $\\sim 5$\\% in this redshift range.  Nevertheless, this does not invalidate the correlation between $\\vrot$ and the star formation efficiency, but rather, strengthens it. The progenitors of the local slowly rotating disks without any recent star formation will be below the detection threshold of the FDF sample, which implies even lower efficiencies should be considered at low $\\vrot$, in agreement with the analysis of Ferreras \\& Silk (2001). Their prediction of a {\\sl steepening} of the slope of the TFR with redshift implies disks with $\\vrot\\simlt 100$~km~s$^{-1}$ at $z\\sim 1$ should be fainter than $M_B\\sim -18$, a result which is also consistent with the models of Boissier \\& Prantzos (2001).  It is intriguing that there is evidence  for very modest episodes of late star formation even in nearby early-type galaxies (Trager et al. 2000). Theory strongly suggests that such episodes should increase with redshift. This is commonly accepted as plausible in galaxy compact groups and protoclusters, where the merger frequency is high. However a comparably strong statement in the environs of typical disk galaxies is puzzling. Very few dwarfs are found near the Milky Way or M31 relative to the predictions of the commonly accepted CDM model. The usual explanation is that companion dwarfs have been stripped of gas and indeed of stars, as suggested by the evidence for tidal streams (Ferguson et al. 2002) and the ongoing disruption of the Sagittarius dwarf galaxy (e.g. Majewski et al. 2003). Such episodes must have been more frequent in the past if theory --  which predicts an order of magnitude more dwarfs than are observed -- is to be reconciled with the observations.  Dwarf disruption will include both compression of gas in the dwarf core as well as gas ejection into the gas reservoir of the host galaxy. Both effects will result in transient episodes of star formation. These enhanced phases of star formation, although individually modest and short-lived, could account for as much as half of the current stellar mass. Such minibursts are also inferred to have occurred in the Milky Way from  the past history of star formation inferred from analysis of chromospheric age indicators in the local stellar population (Rocha-Pinto et al. 2000). The inferred stages of episodic brightening at a lookback time of $5 $~Gyr or more could help account for the brightening that we infer of small disks in the past. We conclude that more data are badly needed at moderate redshift in the $M_B\\sim -18$ range if we are to develop a full understanding of the Tully-Fisher relation.  \\appendix"
    },
    "0403/astro-ph0403646_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:1}   The cosmic \\hbox{X-ray} background (XRB) is largely due to accretion onto supermassive black holes integrated over cosmic time. Thus, extragalactic \\hbox{X-ray} surveys offer the potential to contribute substantially to our understanding of the physics of Active Galactic Nuclei (AGN) as well as the evolution of the AGN population. Such surveys have dramatically advanced over the past four years, largely due to the flood of data from the {\\it Chandra X-ray Observatory\\/} (hereafter \\chandra) and the {\\it X-ray Multi-Mirror Mission-Newton\\/} (hereafter \\xmm). The superb \\hbox{X-ray} mirrors and charge-coupled device (CCD) detectors on these observatories provide  \\begin{description}  \\item{$\\bullet$} Sensitive imaging spectroscopy in the \\hbox{$\\approx 0.5$--10~keV} band, with up to \\hbox{50--250} times (depending upon the energy band considered) the sensitivity of previous \\hbox{X-ray} missions. \\hbox{X-ray} surveys have finally reached the depths needed to complement the most sensitive surveys in the radio, submillimeter, infrared, and optical bands.  \\item{$\\bullet$} \\hbox{X-ray} source positions with accuracies of \\hbox{$\\approx$~0.3--3$^{\\prime\\prime}$}. These high-quality positions are essential for matching to (often faint) multiwavelength counterparts.  \\item{$\\bullet$} Large source samples allowing reliable statistical inferences to be drawn about the overall extragalactic \\hbox{X-ray} source population. In a fairly deep \\chandra\\ or \\xmm\\ observation, $\\simgt 100$--200 sources can be detected.  \\end{description}  \\noindent The extragalactic survey capabilities of \\chandra\\ and \\xmm\\ are complementary in several important respects. The sub-arcsecond imaging of \\chandra\\ provides the best possible source positions, and with long exposures \\chandra\\ can achieve the highest possible sensitivity at energies of \\hbox{$\\approx$~0.5--6~keV}; unlike the case for \\xmm, even the deepest \\chandra\\ observations performed to date do not suffer from significant source confusion. \\xmm, in comparison, has a substantially larger photon collecting area than \\chandra, allowing efficient \\hbox{X-ray} spectroscopy. In addition, \\xmm\\ has better high-energy response than \\chandra\\ and can carry out the deepest possible surveys from \\hbox{$\\approx$~7--10~keV}. Even \\xmm, however, does not cover the peak of the \\hbox{X-ray} background at \\hbox{20--40~keV}. Finally, the field of view for \\xmm\\ is $\\sim 2.5$ times that of \\chandra.  \\begin{table}[t!] \\centering \\caption{Some Deep X-ray Surveys with \\chandra\\ and \\xmm} \\label{tab:1}       % \\begin{tabular}{lrl} \\hline\\noalign{\\smallskip} Survey Name & Exposure & Representative Reference or Note  \\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} \\noalign{\\chandra} \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} \\chandra\\ Deep Field-North          & 1950~ks \\hspace{0.1 in} & D.M. Alexander et~al., 2003, AJ, 126, 539 \\\\ \\chandra\\ Deep Field-South          &  940~ks \\hspace{0.1 in} & R. Giacconi et~al., 2002, ApJS, 139, 369 \\\\ HRC Lockman Hole                    &  300~ks \\hspace{0.1 in} & PI: Murray \\\\ Extended CDF-S                      &  250~ks \\hspace{0.1 in} & PI: Brandt \\\\ Groth-Westphal                      &  200~ks \\hspace{0.1 in} & PI: Nandra \\\\ Lynx                                &  185~ks \\hspace{0.1 in} & D. Stern et~al., 2002, AJ, 123, 2223 \\\\ LALA Cetus                          &  177~ks \\hspace{0.1 in} & PI: Malhotra \\\\ LALA Bo\\\"otes                       &  172~ks \\hspace{0.1 in} & J.X. Wang et~al., 2004, AJ, 127, 213 \\\\ SSA13                               &  101~ks \\hspace{0.1 in} & A.J. Barger et~al., 2001, AJ, 121, 662 \\\\ 3C295                               &  100~ks \\hspace{0.1 in} & V. D'Elia et~al., 2004, astro-ph/0403401 \\\\ Abell~370                           &   94~ks \\hspace{0.1 in} & A.J. Barger et~al., 2001, AJ, 122, 2177 \\\\ SSA22 ``protocluster''              &   78~ks \\hspace{0.1 in} & L.L. Cowie et~al., 2002, ApJ, 566, L5 \\\\ ELAIS                               &   75~ks \\hspace{0.1 in} & J.C. Manners et~al., 2003, MNRAS, 343, 293  \\\\ WHDF                                &   75~ks \\hspace{0.1 in} & PI: Shanks \\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} \\noalign{\\xmm} \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} Lockman Hole                  &  766~ks \\hspace{0.1 in} &  G. Hasinger et~al., 2001, A\\&A, 365, L45 \\\\ \\chandra\\ Deep Field-South    &  317~ks \\hspace{0.1 in} &  A. Streblyanska et~al., 2004, astro-ph/0309089 \\\\ 13~hr Field                   &  200~ks \\hspace{0.1 in} &  M.J. Page et~al., 2003, AN, 324, 101 \\\\ \\chandra\\ Deep Field-North    &  180~ks \\hspace{0.1 in} &  T. Miyaji et~al., 2003, AN, 324, 24 \\\\ Subaru Deep                   &  100~ks \\hspace{0.1 in} &  PI: Watson \\\\ ELAIS S1                      &  100~ks \\hspace{0.1 in} &  PI: Fiore \\\\ Groth-Westphal                &   80~ks \\hspace{0.1 in} &  T. Miyaji et~al., 2004, astro-ph/0402617 \\\\ \\noalign{\\smallskip}\\hline \\end{tabular} \\\\ \\raggedright {The Extended \\chandra\\ Deep Field-South is comprised of four fields (each 250~ks), the \\xmm\\ ELAIS S1 survey is comprised of four fields (each 100~ks), and the \\chandra\\ ELAIS survey is comprised of two fields (each 75~ks). The \\xmm\\ Subaru Deep survey also has seven flanking fields (each $\\approx 50$~ks). Only the first $\\approx 100$~ks of the \\xmm\\ Lockman Hole data have been published at present.} \\end{table}  \\chandra\\ and \\xmm\\ have resolved $\\simgt 80$--90\\%\\ of the \\hbox{0.5--10~keV} XRB into discrete sources, extending earlier heroic efforts with missions including \\rosat, \\asca, and \\sax. The main uncertainties in the precise resolved fraction are due to field-to-field cosmic variance (which leads to spatial variation in the XRB flux density) and instrumental cross-calibration limitations. With the recent advances, attention is now focused on (1) understanding the nature of the \\hbox{X-ray} sources in detail and their implications for AGN physics, and (2) understanding the cosmological evolution of the sources and their role in galaxy evolution. In this review, we briefly describe the key \\chandra\\ and \\xmm\\ extragalactic surveys to date (\\S2) and detail some of their implications for AGN physics and evolution (\\S3). In \\S3 we highlight two topics of current widespread interest: (1) X-ray constraints on the AGN content of luminous submillimeter galaxies, and (2) the demography and physics of high-redshift $(z>4)$ AGN as revealed by X-ray observations. We also discuss prospects for future \\hbox{X-ray} surveys with \\chandra, \\xmm, and upcoming missions (\\S4).  Throughout this paper, we adopt $H_0=70$~km~s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm M}=0.3$, and $\\Omega_{\\Lambda}=0.7$ (flat cosmology).  \\begin{figure}[t!] \\centering \\includegraphics[height=10cm]{fig01.ps} \\caption{Adaptively smoothed image of the 2~Ms \\hbox{CDF-N}, constructed from data in the \\hbox{0.5--2~keV} (red), \\hbox{2--4~keV} (green), and \\hbox{4--8~keV} (blue) bands. Nearly 600 sources are detected in the $\\approx 448$~arcmin$^2$ field. The regions covered by the HDF-N and GOODS-N surveys are denoted. Adapted from D.M. Alexander, F.E. Bauer, W.N. Brandt, et~al., 2003, AJ, 126, 539.} \\label{fig:1} \\end{figure} ",
        "conclusions": ""
    },
    "0403/astro-ph0403193_arXiv.txt": {
        "abstract": "Recently, \\citet{z10} claimed to have discovered a galaxy at a redshift $z=10$, and identified a feature in its spectrum with a hydrogen Ly$\\alpha$ emission line. If this identification is correct, we show that the intergalactic medium (IGM) around the galaxy must be significantly ionized; otherwise, the damping wing of Ly$\\alpha$ absorption by the neutral IGM would have greatly suppressed the emission line.  We find either that the large-scale region surrounding this galaxy must have been largely reionized by $z=10$ (with a neutral fraction $\\la 0.4$) or that the stars within the galaxy must be massive ($\\ga 100\\, M_\\odot$), and hence capable of producing a sufficiently large \\ion{H}{2} region around it. We generalize these conclusions and derive the maximum \\Lya line flux for a given UV continuum flux of galaxies prior to the epoch of reionization. ",
        "introduction": "\\citet{z10} have reported the discovery of a galaxy at a redshift $z=10$, which is gravitationally magnified by a factor $\\sim 25-100$ by the foreground galaxy cluster Abel 1835. The redshift identification is supported by the spectroscopic detection of an emission line at $1.338\\mu$m, the photometric signature of an absorption trough at shorter wavelengths, and the lensing geometry. Here, we examine the far-reaching implications of associating this spectral feature with the Ly$\\alpha$ emission line from a galaxy at $z=10$. In particular, we show that if hydrogen in the intergalactic medium (IGM) around this galaxy were neutral, then the damping wing of its \\Lya absorption would have greatly suppressed the \\Lya emission line of the galaxy.  Our results provide a strong incentive for obtaining higher quality data on this galaxy.  More generally, we derive the maximum \\Lya line flux for a given UV continuum flux of galaxies prior to the epoch of reionization. Future detection of bright \\Lya lines from high-redshift galaxies can be used to improve the current lower limit on the redshift of reionization.  Throughout, we adopt values for cosmological parameters derived from recent observations of the cosmic microwave background \\citep{WMAP, uros}. For the contributions to the energy density, we assume ratios relative to the critical density of $\\Omm=0.27$, $\\Oml=0.73$, and $\\Omega_b=0.046$, for matter, vacuum energy, and baryons, respectively.  We also choose a Hubble constant $H_0=72\\mbox{ km s}^{-1}\\mbox{Mpc}^{-1}$, and a primordial scale-invariant ($n=1$) power-law power spectrum with $\\sigma_8=0.9$, where $\\sigma_8$ is the root-mean-square amplitude of mass fluctuations in spheres of radius $8\\ h^{-1}$ Mpc. ",
        "conclusions": "Figure 1 shows that if the galaxy discovered by \\citet{z10} indeed has a redshift $z=10$, then either its stars are very massive ($\\ga 100M_\\odot$) and hence capable of creating a large \\ion{H}{2} region around it, or the large-scale IGM around it has already been mostly reionized (with a neutral fraction $\\la 0.4$). The ratio of the galaxy's UV continuum flux in two bands is consistent with the first possibility, but better determination of the continuum spectrum or detection of other recombination lines (Bromm, Kudritzki, \\& Loeb 2001; Tumlinson, Giroux, \\& Shull 2001; Oh, Haiman, \\& Rees 2001) is required to make a convincing case.  For example, a stellar population with a top-heavy IMF will produce substantial flux in He II Ly$\\alpha$ ($303.9 $\\AA) and He II $\\lambda 1640$.  A search for features in the spectrum of a redshift $z=10$ galaxy at $\\lambda_{\\rm obs}= 3343$\\AA\\ and $1.8\\mu$m would constrain the first interpretation. The discovery of other galaxies may test the second possibility, since the damping wing of a neutral IGM limits the maximum \\Lya line flux of any galaxy prior to reionization based on its UV continuum flux.  The imprint of the damping wing is insensitive to a variety of complicating factors. Clumping of gas and its infall onto the galaxy have a minor effect, since the damping wing averages the IGM over large physical scales ($\\sim 1$ Mpc).  Moreover, the characteristic size of the \\ion{H}{2} region around a bright galaxy at $z\\sim 10$ is an order of magnitude larger than the scale of its infall region. The characteristic peculiar velocity of galaxies at $z\\sim 10$ is an order of magnitude smaller than the Hubble velocity at the edge of their \\ion{H}{2} regions, and can only significantly change the effect of resonant absorption on the line profile. We have ignored resonant absorption from \\ion{H}{1} within the \\ion{H}{2} region since it would not suppress even the blue wing of the \\Lya emission line if the galaxy had a peculiar velocity (in the direction of recession) larger than the line width (thus allowing the line to appear nearly symmetric); such a peculiar velocity may not be improbable within small samples because it makes the line more easily observable and adds a selection bias.  In any case, the addition of resonant absorption would only strengthen our limits on the \\Lya line flux before reionization. Dust extinction in the host galaxy would affect ionizing photons (and thus the \\Lya line flux) more than it would the continuum at longer wavelengths. Based on the observed UV continuum, we would have predicted a weaker line with the inclusion of extinction, and so our argument would have been stronger.  We have considered the lowest (most probable) value within the range of lensing magnification factors inferred by \\citet{z10}. A higher value would again strengthen our limits. For example, a magnification factor of 100 would imply that the line was 4 times weaker intrinsically, but the predicted line from the UV continuum would then be weakened by more than a factor of 4 since the smaller \\ion{H}{2} region would imply a damping wing that suppressed more of the line flux.  Current observational constraints on reionization provide an inconsistent picture.  On the one hand, the large-scale polarization anisotropies of the cosmic microwave background measured by WMAP imply a reionization redshift of 10-20 (Kogut et al.\\ 2003), while on the other hand, the extent of the \\ion{H}{2} regions around the highest redshift quasars indicates\\footnote{The evidence for a possible Gunn-Peterson (1965) trough in these quasar spectra (White et al. 2003) is still being debated (Songaila 2004).} a significantly neutral IGM at $z\\sim 6.4$ (Wyithe \\& Loeb 2004). Moreover, the IGM should have been cooler than observed at $z\\sim 3$-4 if hydrogen had been fully reionized at $z\\ga 9$ (Theuns et al.\\ 2002; Hui \\& Haiman 2003). Theoretically, it is possible that the ionization fraction evolved in a complex, non-monotonic fashion owing to an early episode of Pop III star formation (Wyithe \\& Loeb 2003; Cen 2003).  Existing observations do not enable us to discriminate between the two interpretations of a positive detection of a \\Lya emission line from the $z=10$ galaxy of \\citet{z10}. Sokasian et al.\\ (2003a,b) have studied the process of reionization using cosmological simulations with radiative transfer; they have shown that for a Scalo-type IMF, reionization occurs over an extended redshift interval.  For models in which the \\Lya optical depth at $z\\approx 6$ matches the value inferred from the SDSS quasars, a substantial fraction of the mass in the universe ($>30 \\%$) is already completely ionized by $z=10$.  Since reionization proceeds ``inside-out'' affecting overdense regions first within the simulated volume (see also Gnedin 2000), we would expect the environment of the \\citet{z10} galaxy to be ionized.  If the evolution had been more complex and included sources with population III stars prior to $z=12$, an even larger fraction of the mass would have been ionized by $z=10$ (Wyithe \\& Loeb 2003; Sokasian et al.\\ 2003b). This picture could change, however, if reionization occurred ``outside-in'' (e.g., Miralda-Escud\\'e, Haehnelt \\& Rees 2000). As shown by Barkana \\& Loeb (2004), the fluctuations in the real universe should be far larger than indicated in existing simulations because of their limited box size. In reality, small fluctuations on scales larger than the box are greatly amplified through biased galaxy formation at high redshifts.  The various scenarios for reionization can be tested with more data on \\Lya emitting galaxies (supplementing recent work by Hu et al. 2002, 2004; Malhotra \\& Rhoads 2002; Kodaira et al.\\ 2003; Santos et al.\\ 2003; Kneib et al.\\ 2004; Stanway et al.\\ 2004; Barton et al.\\ 2004), SDSS quasars (Fan et al.\\ 2003), or gamma-ray burst afterglows (Loeb 2003; Barkana \\& Loeb 2004a). Over the next decade, it may also be possible to map directly the neutral hydrogen in the IGM through its $21\\,$cm line (see recent discussions by Zaldarriaga, Furlanetto, \\& Hernquist 2003; Morales \\& Hewitt 2003; Loeb \\& Zaldarriaga 2003; Gnedin \\& Shaver 2003; Sokasian, Furlanetto \\& Hernquist 2003) with forthcoming instruments such as LOFAR ({\\it http://www.lofar.org/}) or SKA ({\\it http://www.skatelescope.org/})."
    },
    "0403/astro-ph0403470_arXiv.txt": {
        "abstract": "{ Isolated HAEBE stars are believed to represent an intermediate stage of objects between young stellar objects surrounded by massive, optically thick, gaseous and dusty disks and Vega like stars surrounded by debris disks. The star \\aba \\ is already known for being surrounded by an intermediate-stage dust disk emitting a fairly large infrared and (sub-)millimetric excess. Until now, the outer disk structure has only been resolved at millimeter wavelengths and at optical wavelength coronographic imaging. We have obtained 20~$\\mu$m images which show an unexpected ellipse-shaped disk structure in emission at a distance of about 260 AU from the central star.  Large azimuthal asymmetries in brightness can be noticed and the center of the ellipse does not coincide with the star.  A simple, pure geometrical model based on an emission ring of uniform surface brightness, but having an intrinsic eccentricity succeeds in fitting the observations.  These observations give for the first time clues on a very peculiar structure of pre-main-sequence disk geometry, i.e. a non uniform increase in the disk thickness unlike the common usual sketch of a disk with a constant flaring angle.  They provide also valuable informations on the disk inclination as well as its dust composition; at such a large distance from the star, only transient heating of very small particles can explain such a bright ring of emission at mid-infrared wavelengths.  Finally, the increase of thickness inferred by the model could be caused by disk instabilities; the intrinsic eccentricity of the structure might be a clue to the presence of a massive body undetected yet. ",
        "introduction": "Herbig Ae/Be (HAEBE) stars represent a class of intermediate mass, pre-main-sequence (PMS) stars, first described as a group by \\cite{herbig1960}. The circumstellar (CS) disks found around these stars are believed to be the sites of on-going planet formation. By studying the characteristics and evolution of the CS disk and its dust composition, valuable insights can be obtained into the processes leading to the formation of planets, and put constraints on disk and planet formation models. Infrared spectroscopy obtained with the Infrared Space Observatory (ISO) has given us insight into the dust composition of sample of isolated HAEBE systems \\citep[e.g.][]{abaur,processing,HerbigOverview}. While these spectra reveal a rich mineralogy, no direct information concerning the spatial distribution of the different dust species can be inferred from the ISO data. Most studies so far have used the available spectral energy distributions (SEDs) to put constraints on the spatial distribution of the CS material. Models for passively heated disks surrounding PMS stars are successful in reproducing the ISO spectra \\citep[e.g.][]{dullemond2001, dominik2003}, but these models cannot be uniquely constrained from SED fitting alone \\cite{degenaracy}, for this spatially resolved imaging, as presented in this paper, is required. \\begin{figure*}[t] \\vspace*{-3cm} \\parbox{11.0cm}{ \\begin{center} \\vspace*{2.5cm} \\resizebox{13.0cm}{!}{\\includegraphics[angle=90]{image.ps}} \\end{center} } \\hspace*{-0.5cm} \\parbox{6.5cm}{ \\begin{center} \\vspace*{2.5cm} \\resizebox{7cm}{10cm}{\\includegraphics[bb= 5cm 10.3cm 16cm 29.2cm,clip]{fig1_bis.ps}} \\end{center} }  \\label{fig-image} \\caption{Image of the AB Aurigae disk at 20.5~$\\mu$m. The left figure shows the deconvolved image, with a pixel scale of 0.3\\arcsec/pixel.  Clearly visible is a resolved central emission region surrounded by a ring like structure.  The panels on the right show the normalized intensity profiles along a cut through the CAMIRAS images of 2 reference stars and AB~Aur. The upper right panel shows the comparison between two observations of reference stars, demonstrating the stability of the PSF. The panel on the lower right shows a comparison of AB~Aur with the PSF, clearly showing that the central emission is extended and that a ring-like emission structure is also detected in non deconvolved data.  } \\end{figure*}  Among isolated HAEBE systems, the disk around \\aba \\ is one of the most interesting and studied. Its star has a probable age of 2~Myr \\citep{mario1997}, indicating that, according to current planet formation theories, planet building could still being ongoing in this system.  ISO spectra show strong PAH emission bands, and emission from silicates, and carbonaceous dust grains \\citep{mario_abaur,abaur}.  Though the inferred grain sizes are differing from interstellar grains, the dust around AB~Aur seems to be relatively unprocessed, indicating an evolutionary young system. This seems also to be confirmed by combined ISO-SWS and sub-millimeter observations of H$_2$ and CO rotational lines, demonstrating that the disk still has a large gas content \\citep{Thi2001}.  Its disk has been resolved in the millimeter range by \\citet{mannings1997}, showing a structure consistent with a Keplerian disk.  Its surrounding nebulae and outer disk structure were also studied in the visible range using broad-band coronographic observations \\citep{grady_stis_abaur}, showing a disk with spiral-shaped structures. Near-IR interferometric \\citep{millan1999, millan2001, eisner2003}, have resolved the inner parts of the disk, showing it to be consistent with a passive disk with an inner hole, seen at a low ($\\le 40\\degr$) inclination angle. Recent mid-IR imaging \\citep{chen2003}, has also resolved the inner structure, showing it to originate from thermal emission from dust grains heated by the stellar radiation field near the central star. Here we present mid-IR imaging, not only resolving the thermal emission from dust close to the central star, but also an emission structure at the outer parts of the disk. In the next sections we will describe these observations and propose a pure geometrical model for explaining our results, unlike the paradigm according to which gravitational interactions with massive bodies are inferred when observing brightness asymmetries in dusty disks. ",
        "conclusions": "We have presented in this paper a thermal infrared image at 20.5 $\\mu$m of the \\aba \\ dusty disk.  The deconvolved image shows an inner resolved structure containing the majority (95 \\%) of the thermal flux at this wavelength and an unexpected ring-like ellipse-shaped structure. This ring is located at an average distance of 260 AU from the star, but its most interesting features lie in large scale asymmetries and a center that is offset with respect to the star.  When modeling the ring structure using a purely geometrical model (in which asymmetries are essentially produced by inclination and shadowing effects in a flared disk), we end with the conclusion that we are in presence of an {\\bf intrinsically ellipse-shaped} ring that could trace the presence of puffed-up matter from a {\\bf flared dust disk} containing transiently heated particles.  Self-shadowing instabilities perturbating the disk vertical thickness could produce such a ring. However, the origin of its intrinsic eccentricity remains unclear; massive bodies could be acting in gravitationally structuring the disk. Further detailed modelings still need to be investigated to fully understand the physics of this intriguing structure; further observations, including higher resolution imaging in several PAH bands and spatially resolved mid-IR spectroscopy, are needed to better understand this object."
    },
    "0403/astro-ph0403187_arXiv.txt": {
        "abstract": "We have monitored the phase of the main X-ray pulse of the Crab pulsar with the Rossi X-ray Timing Explorer (RXTE) for almost eight years, since the start of the mission in January 1996. The absolute time of RXTE's clock is sufficiently accurate to allow this phase to be compared directly with the radio profile.  Our monitoring observations of the pulsar took place bi-weekly (during the periods when it was at least 30 degrees from the Sun) and we correlated the data with radio timing ephemerides derived from observations made at Jodrell Bank. We have determined the phase of the X-ray main pulse for each observation with a typical error in the individual data points of 50~$\\mu$s.  The total ensemble is consistent with a phase that is constant over the monitoring period, with the X-ray pulse leading the radio pulse by 0.0102$\\pm$0.0012 period in phase, or 344$\\pm$40~$\\mu$s in time.  The error estimate is dominated by a systematic error of 40~$\\mu$s in the radio data, arising from uncertainties in the variable amount of pulse delay due to interstellar scattering and instrumental calibration.  The statistical error is 0.00015 period, or 5~$\\mu$s. The separation of the main pulse and interpulse appears to be unchanging at time scales of a year or less, with an average value of 0.4001$\\pm$0.0002 period.  There is no apparent variation in these values with energy over the 2-30 keV range.  The lag between the radio and X-ray pulses may be constant in phase (i.e., rotational in nature) or constant in time (i.e., due to a pathlength difference).  We are not (yet) able to distinguish between these two interpretations. ",
        "introduction": "For many years it has been assumed that the main pulse and interpulse of the Crab pulsar (PSR~B0531+21) are perfectly lined up in phase over the full range of the electro-magnetic spectrum. Even though there have been reports in the past that this alignment may not be as perfect as generally assumed, the absolute calibration of spacecraft clocks was not sufficiently accurate to allow a precise measurement of the phase difference. The most compelling result predating the Rossi X-ray Timing Explorer (RXTE) observations was presented by \\citet{masn1994}, based on Figaro~II observations covering 0.15 to 4.0~MeV, that were made in 1986 and 1990.  While discounting the 1986 result which has a considerable uncertainty due to potential errors in the dispersion measure, we consider the 1990 result (the $\\gamma$-ray pulse leading the radio pulse by 375$\\pm$148~$\\mu$s) fairly reliable, though the error is probably underestimated.  The precision with which absolute time can be determined with the RXTE clock allows us to time X-ray pulses with an accuracy better than 10~$\\mu$s, depending on pulse shape, as shown by \\citet{rots1998b}. At the same time, the Crab pulsar monitoring program at Jodrell Bank provides timing ephemeris data, reduced to infinite frequency. These (monthly) timing ephemerides represent fits to the daily time-of-arrival measurements with rms residuals of order 20-50~$\\mu$s. This allows us to measure and monitor the radio to X-ray phase difference of the pulses with an error of about one milli-period. We have reported on these results in the past (\\citet{rots1998a}, \\citet{rots1998c}, \\citet{rots2000a}, \\citet{rots2000b}).  At optical wavelengths, \\citet{sanw1999} has reported a time delay of 140~$\\mu$s (optical leading the radio), but the details are not easily accessible.  \\citet{shea2003} report that, in the case of giant radio pulses, the optical pulse in the wavelength range 600-750~nm is leading the radio pulse by 100$\\pm$20~$\\mu$s. \\citet{roma2001}, on the other hand, claim that the optical (355-825~nm) and radio peaks are coincident within 30~$\\mu$s, based on test observations with a prototype transition-edge sensor detector. However, it is not clear whether the timing calibration of the instrument was complete at the time.  \\citet{ulme1994} presented results from OSSE observations (50-100~keV), indicating that the hard X-ray to radio lag was $<$30$\\pm$30~$\\mu$s.  It would appear that the estimate of their errors was too optimistic. \\citet{nola1993} present pulse profiles but no absolute phases. \\citet{kuip2003} report on INTEGRAL data, covering 6-50~keV, and deriving a time delay (the radio trailing) of 280$\\pm$40~$\\mu$s for a single epoch.  The precise timing of the pulses in the different wavelength regimes has important repercussions for the understanding of the nature and spatial origin of the emission processes that give rise to the pulses in different parts of the spectrum.  \\citet{roma1995} have suggested that, while the radio precursor comes from the polar cap, the pulse and interpulse originate in the outer gap in the magnetosphere, with the higher energy pulses being generated at significantly greater height.  Thus, measuring the pulse shapes and the absolute timing throughout the electromagnetic spectrum places important constraints on the shape of the the outer gap and on the height in the magnetosphere at which the radiation is generated.  In this paper we will present the results of the RXTE monitoring campaign of the Crab Pulsar from the start of that mission.  We have adopted the radio nomenclature for the features in the pulse profile (main pulse, bridge, and interpulse). ",
        "conclusions": "The X-ray main pulse leads its radio counterpart by about 344$\\pm$40~$\\mu$s (systematic error); the statistical error is 5~$\\mu$s.  This is more than twice the time difference of 140~$\\mu$s that \\citet{sanw1999} determined for the optical B band and three times the 100~$\\mu$s measured by \\citet{shea2003} in red light. The time or phase difference appears to have been constant over the past eight years, but the data are not accurate enough to distinguish between the two. The errors in individual measurements range up to 50~$\\mu$s. We should caution the reader when making comparisons with results in other wave bands.  First, various authors have used different definitions of the pulse phase.  In our estimation, our own definition agrees with that used in the radio band, as does the definition of \\citet{shea2003}, but that is probably not true for most other reports.  Second, the scatter in values for individual observations is fairly large (a milliperiod) and may be intrinsic.  Greater accuracy can only be achieved with a statistically significant set of observations.  As we have mentioned, there is a systematic error of up to 40~$\\mu$s in the offset due to uncertainties in the interstellar scattering and the calibration of the radio equipment.  This error may change on timescales of a few months, but since we do not know whether (and if so, by how much) the error is reduced by averaging, we quote 40~$\\mu$s as the final uncertainty for the radio to X-ray timing of the pulse. However, such an error does not affect the comparison with results from other wavebands provided they are all approximately contemporaneous and use the same Jodrell Bank timing ephemeris records.  On the other hand, it also appears that most (if not all) errors for the results at other wavebands quoted in the Introduction are seriously underestimated by ignoring the systematic error.  The phase difference between the two X-ray pulses is constant at 0.400 period, within the measurement errors. It is also equal to the phase difference between the radio main pulse and interpulse, within the measurement error.  It may be of interest to note that in the X-ray pulse profile the trailing edges of the pulse as well as the interpulse are distinctly steeper than their leading edges.  This does not appear to be the case for the optical interpulse.  If the X-ray to radio lag were a true phase lag, attributable to the (radial) energy distribution across a cone, with the pulses occurring near the cone edges, one would expect the placement to be symmetrical, i.e., one X-ray pulse to be leading, the other trailing. As it stands, both X-ray pulses are leading by the same amount. The simplest explanation for this phenomenon is that we are dealing with a time delay reflecting a pathlength difference: the radio pulses originate approximately 100~km closer to the surface of the neutron star, as already suggested by \\citet{masn1994}."
    },
    "0403/astro-ph0403378_arXiv.txt": {
        "abstract": "Quartessence is one of the alternatives to $\\Lambda$-CDM that has lately attracted considerable interest. According to this unifying dark matter/energy scenario, the Universe evolved from an early non-relativistic matter-dominated phase to a more recent accelerated expansion phase, driven by a single fluid component. Recently, it has been shown that some problems of the quartessence model, such as the existence of instabilities and oscillations in the matter power spectrum, can be avoided if a specific type of intrinsic entropy perturbation is considered. In the present article we explore the role of skewness in constraining this non-adiabatic scenario. We show that non-adiabatic quartessence and quintessence have different signatures for the skewness of the density distribution on large scales and suggest that this quantity might prove helpful to break possible degeneracies between them. ",
        "introduction": "According to the current standard cosmological model, the dynamics of the universe would be dominated by two unknown components: dark-matter (DM), responsible for structure formation, and dark-energy (DE), that causes the accelerated expansion. Although there are several candidates for both DM and DE, there is still no evidence of either of them in laboratory physics. From the point of view of simplicity, it would be interesting to explore the possibility that a single component plays the role of both DE and DM, reducing from two to one the unknown constituents of the universe. A model that provides a single description of DE and DM through \\textquotedblleft unifying-dark-matter\\textquotedblright\\ or simply quartessence \\cite{makler03} has attracted a lot of interest recently. A prototype of this model is given by the quartessence Chaplygin model (QCM) \\cite{kamenshchik01}.  Both the background and linear fluctuations were extensively studied for QCM, and were compared to observational data. The generalized Chaplygin gas (as quartessence) appears to be compatible with all available data regarding the expansion history (see e.g. ref. \\cite{makler03b} and refs. therein). For adiabatic perturbations, a linear analysis was done for the CMB \\cite{amendola03}\\ and LSS \\cite{beca03}. In this case, only QCM models close to the ``$\\Lambda$CDM limit'' are allowed. Recently, it was shown that problems (pointed out in \\cite{sandvik02}), such as the existence of instabilities and oscillations in the matter power spectrum of QCM, can be avoided if a specific type of intrinsic entropy perturbation is considered. Such non-adiabatic model is consistent with the 2dF power spectrum for any value of the model parameters in the permitted interval, as long as the effective shape parameter assumes certain values \\cite{reis03b}. An \\textquotedblleft averaging problem\\textquotedblright\\ was also pointed out as a shortcoming of quartessence \\cite{avelino03}. However, it is straightforward to show that the above mentioned non-adiabatic quartessence does not suffer this kind of problem \\cite{obs}.  Thus, up to the present time, we can say that adiabatic quartessence is disfavored by the data, but it is not possible to distinguish non-adiabatic quartessence from concordance models like $\\Lambda $CDM and quintessence using the previously considered observables. However, as we shall see, measurable differences in the predictions of these models appear clearly in the nonlinear regime, in particular in the skewness of the matter distribution in large scales. In our investigation we specifically consider three different quartessence models. All of them have the $\\Lambda $CDM model as a limiting case for the background solution. The analysis of these three cases indicates that our result should be applicable to more generic quartessence models.  While most studies of the nonlinear regime deal with the clumping of pressureless fluid (DM), in the case of quartessence it is imperative that one includes the effects of pressure. Accordingly, we apply, and somewhat extend to include relativistic pressure, a method for the computation of density cumulants developed in refs. \\cite{bernardeau92,bernardeau94,fosalba98}. ",
        "conclusions": ""
    },
    "0403/astro-ph0403652_arXiv.txt": {
        "abstract": "{ We present a spectroscopic analysis of two galaxy clusters at z$\\approx0.2$, out to $\\sim4$\\,Mpc. The two clusters VMF73 and VMF74 as identified by \\citet{VMFJQH98} were observed with multiple object spectroscopy using MOSCA at the Calar Alto 3.5\\,m telescope. Both clusters lie in the \\mbox{ROSAT} Position Sensitive Proportional Counter field R285 and were selected from the X-ray Dark Cluster Survey \\citep{GBCZ04} that provides optical $V$- and $I$-band data. VMF73 and VMF74 are located at respective redshifts of z$=0.25$ and z$=0.18$ with velocity dispersions of 671 km\\,s$^{-1}$ and 442 km\\,s$^{-1}$, respectively. Both cluster velocity dispersions are consistent with Gaussians. The spectroscopic observations reach out to $\\sim2.5$ virial radii. Line strength measurements of the emission lines H$_\\alpha$ and [O\\,II]$\\lambda$3727 are used to assess the star formation activity of cluster galaxies which show radial and density dependences. The mean and median of both line strength distributions as well as the fraction of star forming galaxies increase with increasing clustercentric distance and decreasing local galaxy density. Except for two galaxies with strong H$_\\alpha$ and [O\\,II] emission, all of the cluster galaxies are normal star forming or passive galaxies. Our results are consistent with other studies that show the truncation in star formation occurs far from the cluster centre. ",
        "introduction": "\\label{sec:Intro} Galaxy properties such as colour, morphology and spectral characteristics are functions of redshift as well as of galaxy environment. Examples for cluster specific redshift dependences are the Butcher-Oemler effect \\citep{BO78a} and the change of the ``morphological mix'' in clusters with redshift. Local clusters like the Coma cluster are dominated by S0 galaxies, which form only a small fraction in higher redshift clusters at z$\\sim0.5$, whereas spiral galaxies are more abundant in distant than in local clusters. This observation has raised the question whether the spiral galaxies in distant clusters are the progenitors of S0 galaxies in local clusters, and which kind of processes may be responsible for a transformation of one morphological type into another \\citep[e.g.][]{DOCSE97,PSDCB99}. Further evidence for an evolution of galaxy properties with redshift is provided by studies which show that the universal average star formation rate (SFR) has been rapidly declining since z$\\sim$1 \\citep[e.g.][]{LLHD96,MFDGS96,BSIK99,SPF01}. However, at a given epoch, cluster galaxies always show suppressed star formation compared with the field population at the same redshift \\citep[e.g.][]{BMYCE97}. Models of hierarchical structure formation predict a continuous accumulation of material and substructure to more and more massive galaxies, groups and clusters as time progresses \\citep[e.g.][ and references therein]{Kauff96,KCDW99b,CLBF00}. Therefore, it may be possible to link the decline in the global star formation rate to the growth of large scale structure in the universe.\\\\ A number of studies have investigated the connection between the cluster environment and morphological gradients \\citep[e.g.][]{Dress80,PG84,WG93,DOCSE97,DML01}. To characterize the cluster environment, the clustercentric distance of a galaxy as well as the local projected galaxy density are useful as independent parameters. Both of them, however, are projected parameters. The morphological dependence of cluster galaxies on the environment has been studied based on the clustercentric radius e.g. by \\citet{WG91,WG93} and based on local galaxy density by \\citet{Dress80,DOCSE97} and others. It has been discussed which parameter is more appropriate, i.e. whether galaxy morphology is influenced rather by global properties, characterized by the clustercentric distance, or by local properties like the local galaxy density. However, in centrally concentrated, regular and presumably relaxed clusters the local galaxy density is closely correlated with the clustercentric radius \\citep{DOCSE97,GNMBG03}. A correlation between morphological gradients and clustercentric distance or local galaxy density has been found, such that the dense cluster core is populated predominantly by ellipticals and S0 galaxies. Towards the cluster outskirts, the fraction of spirals increases while the elliptical and S0 fraction decreases \\citep{Dress80}.\\\\ In a similar way, the star formation activity of galaxies depends on the environment. This dependence has been investigated by various studies, which reveal a lower level of star formation in cluster galaxies compared with the field population at the same redshift \\citep{ASHCY96,BSMYC98,BBSZD02}. The results of these studies suggest that star formation is suppressed in cluster galaxies over a wide range of cluster masses. To a certain extent it is not surprising to find lower SFRs of galaxies in high density environments as these dense regions like cluster cores are dominated by galaxy types with an intrinsically lower SFR. However, recent studies show that even within the same galaxy type the SFR is lower in the cluster population compared to the surrounding field \\citep{BSMYC98,GNMBG03}.\\\\ The present analysis focuses on star formation properties of galaxies in clusters at z\\,$\\sim0.2$ out to large distances from the cluster centre. Studies of \\mbox{CNOC\\,1} data explored clusters out to $\\sim1$ virial radius (R$_v$) and found that mechanisms responsible for the low SFR of cluster galaxies are likely to start acting far out of the cluster core in the infall region, where field galaxies attracted by the cluster potential experience the influence of the cluster environment for the first time \\citep{BMYCE97,BSMYC98}. It has been shown that galaxies as far out as $\\sim2$\\,R$_{v}$ may have crossed the cluster core within a Hubble time \\citep{RS98,BNM00}. This suggests that galaxies in the cluster outskirts may have undergone significant evolution on their way through the cluster core and that cluster specific processes start acting on infalling galaxies even beyond 2\\,R$_v$. It has been discussed, for instance, whether the cluster environment induces starbursts in infalling galaxies which lead to a rapid consumption of the gas supply and a subsequent passive evolution of these galaxies \\citep{DG83,BD86,ASHCY96}. \\\\ In the local universe, the cluster environment has been explored from the densest central regions to large radial distances and into the field by studies of \\mbox{2dFGRS} and \\mbox{SDSS} data \\citep{LBDCB02,GNMBG03,BEMLB04}. At intermediate and high redshift, only few comparable studies have yet been conducted \\citep{ASHCY96,KSNOB01}. Galaxy properties, however, depend strongly on redshift. Therefore, to understand the evolution of galaxies, it is essential to investigate different environments at different redshifts. In this study we present observations of two low-mass clusters at z$\\sim0.2$ as part of a larger program to study the correlation between star formation and environment at this redshift.\\\\ The outline of the paper is as follows. In chapter 2, the selection of objects, observations and data reduction are described. The data analysis and results are presented in chapter 3. In chapter 4, the results are discussed and compared with previous studies, and chapter 5 is a summary of our findings. Throughout this paper, we use a cosmology of H$_0=70$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_m=0.3$ and $\\Omega_{\\Lambda}=0.7$ ",
        "conclusions": "\\label{sec:Dis} \\subsection{Comparison with earlier work} \\label{sec:comp} The relation between star formation properties of galaxies and their environment has been investigated by various studies of the local universe \\citep{LBDCB02,GNMBG03} and at intermediate redshift \\citep[e.g.][]{BBSZD02}. In all of these studies, the clustercentric distance as well as the local galaxy density are used to characterize the cluster environment.\\\\ \\citet{LBDCB02} and \\citet{GNMBG03} have analysed the star formation activity of \\mbox{2dFGRS} and \\mbox{SDSS} galaxies in the redshift range $0.05\\leq$\\,z\\,$\\leq0.1$. Both studies derive SFRs from the H$_{\\alpha}$ flux and analyse star formation properties as a function of clustercentric radius as well as of local galaxy density. \\citet{GNMBG03} investigate furthermore the environmental dependence of the H$_{\\alpha}$ and [O\\,II] emission line strengths. The trends found in the present study are qualitatively similar to the ones detected in these earlier studies. Star formation activity decreases with decreasing clustercentric radius and increasing local galaxy density. The median of both line strength distributions in the sample of VMF73 and VMF74 increases from $\\sim0$\\,\\AA\\ in the centre to $\\sim5$\\,\\AA\\ in the outer regions. This is in good agreement with \\citet{GNMBG03}. The same study finds density dependences of the star formation activity such that the median of the H$_\\alpha$ and [O\\,II] line strength distributions decreases with increasing local galaxy density. The density estimates used by \\citet{GNMBG03} are different from the ones in the present study. \\citet{GNMBG03} apply a background subtraction using redshifts while in the present analysis a statistical background subtraction was applied. As a result, the densities in Figure~\\ref{fig:densEm} are much larger than those of \\citet{GNMBG03} due to the difference in projection. In the VMF73 and VMF74 clusters, we see a qualitatively similar effect in H$_\\alpha$ and [O\\,II]; however, the sample size is too small to warrant a detailed comparison.\\\\ In a redshift range comparable to the one of the present analysis, \\citet{BBSZD02} have carried out a study of ten galaxy clusters with low X-ray luminosity. Galaxy clusters of low X-ray luminosity have been chosen in order to study the environmental dependence of galaxy SFRs and possible mechanisms suppressing the star formation in low mass clusters which are considered the progenitors of more massive clusters in models of hierarchical structure formation. The cluster sample of VMF73 and VMF74 shows an increase in the mean and median of the [O\\,II] line strength distribution from the centre to the outer regions which is consistent with \\citet{BBSZD02}. However, there is large scatter in the data of VMF73 and VMF74 and the observed trends in mean and median are of limited statistical significance. A stronger trend is found in the increase of the galaxy fraction with W$_0$([O\\,II])\\,$>5$\\,\\AA\\ from $\\sim15$\\% in the centre to $\\sim50$\\% in the outer fields. \\citet{BBSZD02} find an increase in the fraction of star forming galaxies from $\\sim10$\\% in the centre to $\\sim30$\\% at $\\sim1$\\,R$_v$. Density trends in the sample of VMF73 and VMF74 have been detected such that the mean and median of the [O\\,II] line strength distribution decrease with increasing density, while \\citet{BBSZD02} find that the mean reaches $\\sim6$\\,\\AA\\ only in the lowest density regions. Both studies detect similar density trends in the fraction of galaxies with W$_0$([O\\,II])\\,$>5$\\,\\AA\\@.\\\\ As a further indicator for star formation, galaxy colours in the two clusters VMF73 and VMF74 were investigated. We find a radial trend in the colours such that with increasing radius the average colours become bluer. At large radii beyond $\\sim1.5$\\,R$_v$, there are few galaxies with a dominant old, red population. On the other hand, there are some rather blue galaxies without emission found in the central regions. This population may consist of galaxies which had ongoing star formation at the time of their infall into the cluster environment. On their way through the cluster, the diffuse gas in the halo may have been lost which therefore cannot replenish the fuel for star formation in the disk and the SFR slowly declines \\citep{BNM00}. However, this mechanism predicts a systematically lower SFR of star forming galaxies in high density environments compared with similar galaxies in lower densities. A recent study of SDSS and 2dF data finds that the lower level of star formation in high density environments is largely due to the smaller fraction of star forming galaxies in these environments, while there is little difference in the star forming galaxy populations in high and low density regimes \\citep{BEMLB04}. The sample of VMF73 and VMF74 is too small to draw any strong conclusions from the population of blue galaxies without H$_\\alpha$ emission. Figure~\\ref{fig:vicHab} suggests that there is a small difference in the H$_\\alpha$ emission between blue, star forming galaxies in the inner and outer fields. Furthermore, in the VMF73 and VMF74 sample a distinctive colour change is visible at a clustercentric radius of $\\sim1$\\,R$_v$ where the red sequence of both clusters disappears. This feature is reminiscent of the work of \\citet{KSNOB01} who find an abrupt colour change at a characteristic density of 2 galaxies h$^2_{50}$ Mpc$^{-2}$ in the cluster Cl~0939+4713 at z$=0.41$.  \\subsection{Physical Implications} The aim of this investigation has been to explore the star formation properties of galaxies over a wide range of densities with a special focus on galaxies in large distances from the cluster core. The difference between the present study and the two previous ones cited above is that we investigated cluster galaxies in a higher redshift regime than did \\citet{GNMBG03} analysing \\mbox{SDSS} data. Furthermore, this work reaches out to larger clustercentric distances than the study of X-ray faint clusters at intermediate redshift \\citep{BBSZD02}. In regions as far out as $\\sim3-4$\\,R$_v$ from the cluster centre, galaxies from the surrounding field are thought to fall continuously into the cluster attracted by its gravitational potential and experience for the first time the influence of the cluster environment. Therefore, galaxies in these infall regions which denote a transition between the low density field and the cluster are expected to have intermediate properties between the field and cluster galaxy population. The results of \\citet{GNMBG03} and \\citet{BBSZD02} are consistent with this picture in finding an increase in line strengths and SFRs with increasing clustercentric distance and decreasing density. The present study finds similar radial and density dependences of galaxy star formation. The change of galaxy colours with radius is a further indication for some kind of transformation of cluster galaxies. The observed disappearance of the red sequence at a radius of $\\sim1$\\,R$_v$ is similar to an abrupt colour break reported by \\citet{KSNOB01} at a characteristic density of 2 galaxies h$^2_{50}$ Mpc$^{-2}$. \\citet{KSNOB01} find that this density corresponds to subclumps outside the cluster centre. This may be a further hint that galaxy transformation first occurs in groups and subclumps far from the centre.\\\\ Several processes are being discussed in the literature as candidates for suppressing star formation in cluster galaxies. The radial dependence of star formation properties suggests that this mechanism starts to act about 2.5\\,R$_v$ from the cluster core and possibly even beyond. This result rules out mechanisms that are effective only in the highest density regimes, such as ram pressure stripping, as being solely responsible for the decreasing SFRs throughout the cluster. Although more than one process may be involved and ram pressure stripping may still play a role in the cluster core, it is likely that in the outer regions a more subtle mechanism is at work which does not halt the star formation abruptly. Apparently, this mechanism starts to affect galaxies relatively early after their first contact with the cluster environment, for even galaxies far out of the central regions show a reduced star formation. The scenario of strangulation \\citep{BNM00} in which the diffuse gas in the dark halo of an infalling galaxy is lost may be an explanation for the reduced SFRs in cluster galaxies. As discussed above, this mechanism predicts lower SFRs of blue galaxies in high density regions compared with blue galaxies in lower density regimes. However, based on the blue galaxy population in the VMF73 and VMF74 cluster sample alone, we cannot draw any strong conclusions about the mechanism suppressing star formation in the cluster environment. The observed decrease in SFRs is compatible with strangulation, although it is possible that other effects such as galaxy-galaxy interactions play a role or were important in the past and influence galaxy evolution in small groups even before a galaxy enters the cluster environment.  \\subsection{Outlook} The reduction and analysis of the remaining four galaxy clusters will increase the sample size and enable us to investigate the environmental dependence of galaxy star formation properties with more precision and statistical significance. With the complete sample of six galaxy clusters, both the SFR-radius relation and the SFR-density relation can be studied without the limitations of low number statistics and thus help us to understand the evolution of galaxies in the cluster environment.\\\\ Additionally, it would be interesting to study the dependence of star formation activity on the X-ray luminosity of clusters. Groups and low mass X-ray faint clusters are considered the progenitors of more massive clusters in hierarchical merging models. Studying clusters over a range of masses is therefore an important step to investigate the growth of structure in the universe."
    },
    "0403/astro-ph0403008_arXiv.txt": {
        "abstract": "s{We model the spectral energy distribution of the ultrasoft broad-line AGN RE J2248-511 with Comptonised accretion disc models. These are able to reproduce the steep optical and ultrasoft X-ray slopes, and the derived black hole mass is consistent with independent mass estimates. This AGN displays properties of both broad and narrow line Seyfert 1 galaxies, but we conclude that it is intrinsically a `normal' Seyfert 1 viewed at high inclination angle.} ",
        "introduction": "RE J2248-511 is a nearby ($z=0.101$) EUV-selected Seyfert galaxy discovered by the \\emph{ROSAT} Wide Field Camera\\cite{Po}. Further observations showed this source to have a strong soft X-ray excess and spectral variability at optical\\cite{Ma} \\cite{Gru} and soft X-ray\\cite{emp95} \\cite{aad} wavelengths. However, none of these observations were simultaneous, so the existence of an optical to soft X-ray big blue bump (BBB) could not be confirmed.  The soft X-ray spectrum resembles those of narrow-line Seyfert 1 galaxies (NLS1's), but this AGN has high velocity optical emission lines\\cite{Ma} with H$\\beta$ FWHM$\\sim$2900 km s$^{-1}$, which classifies it as a Seyfert 1. Studies of \\emph{ROSAT} PSPC slopes in AGN had concluded that sources with both steep soft X-ray continuum slopes and broad optical emission lines are not found in nature\\cite{BBF} \\cite{WB} \\cite{Gru}. This makes RE J2248-511 an unusual and ideal case in which to examine the relationship between the X-ray and optical continua, specifically the interpretation of the BBB as thermal emission from an accretion disc. We describe the results of a multiwavelength monitoring campaign of RE J2248-511, consisting of X-ray observations from the \\emph{XMM} satellite with supporting quasi-simultaneous optical observations made at the South African Astronomical Observatory (SAAO) and archival multiwavelength data. ",
        "conclusions": "We have shown that Comptonised accretion disc models, which treat the optical to soft X-ray emission as a single BBB, are able to comprise the majority of this flux. From these we derive constraints on the black hole mass which are fully consistent with the mass we obtain through the independent photoionisation technique. This implies that thermal disc emission is the likely origin of the optical, UV and some of the soft X-ray continuum. An accretion disc does not, however, constitute the soft excess as observed with the pn, but it could be the origin of the `ultrasoft' component below 0.25 keV observed with \\emph{ROSAT}. The 0.3-2 keV soft component observed with \\emph{XMM} has a blackbody-like shape, but is too hot to be blackbody disc emission. A model including Comptonisation of soft photons in a hot plasma provides a good fit. Between the \\emph{XMM} observations, the X-ray spectral shape of RE J2248-511 remained approximately constant, while comparison with previous X-ray data shows long-term variability. This demonstrates that the `soft' state is a long-lived phase, and not, for example, a rapid flaring of the disc.  Is RE J2248-511 an intermediate class of object linking the Seyfert 1's with the NLS1's, or a true Seyfert 1 seen with a particular observational bias? \\\\ Contrary to the proposed NLS1 scenario\\cite{PDO}, the data are best fitted with high black hole masses ($\\sim10^{8}$ M$_{\\odot}$), whilst still favouring the high accretion rates suggested for NLS1 galaxies. Therefore, it is not necessary for a black hole to have a low mass for the formation of an ultrasoft X-ray excess. Comparison of Comptonised accretion discs to the SED show that the orientation of the disc is close to face-on, allowing us to see a greater surface area of the accretion disc. If most of the EUV emission arises in the accretion disc then this source would appear EUV-bright compared with similar Seyfert galaxies viewed more edge-on. In the hard X-rays RE J2248-511 resembles a normal Seyfert 1 far more than a NLS1\\cite{BME}. The soft X-ray flux measured with \\emph{XMM} is typical of a Seyfert 1\\cite{TP}, and its slope is similar to that found in the PG bright quasar sample \\cite{Laor}. This source also follows the observed correlation between Balmer linewidth and soft X-ray slope\\cite{emp92} \\cite{Laor} if the `ultrasoft' part of the X-ray spectrum (below 0.25 keV) is excluded. Variability in the soft X-ray excess is a property which RE J2248-511 shares with the NLS1 galaxies, but changes are often more dramatic in NLS1's than observed here.  The strength and broad wavelength span of the BBB in RE J2248-511 is unusual for a broad-line AGN, but we find this may be explained by a face-on Comptonised accretion disc, likely to be accreting at a high rate onto a 10$^{7.5}$-10$^{8.5}$ M$_{\\odot}$ black hole. We propose that this source is intrinsically a normal Seyfert 1, but shares the ultrasoft X-ray excess property with the NLS1's because it is observed at a higher inclination angle than the majority of Seyfert 1 galaxies."
    },
    "0403/astro-ph0403522_arXiv.txt": {
        "abstract": "Here we report our discovery of a band of blue luminescence (BL) in the Red Rectangle (RR) nebula. This enigmatic proto-planetary nebula is also one of the brightest known sources of extended red emission as well as of unidentified infra-red (UIR) band emissions. The spectrum of this newly discovered BL is most likely fluorescence from small neutral polycyclic aromatic hydrocarbon (PAH) molecules. PAH molecules are thought to be widely present in many interstellar and circumstellar environments in our galaxy as well as in other galaxies, and are considered likely carriers of the UIR-band emission. However, no specific PAH molecule has yet been identified in a source outside the solar system, as the set of mid-infra-red emission features attributed to these molecules between the wavelengths of 3.3 $\\mu$m and 16.4 $\\mu$m is largely insensitive to molecular sizes. In contrast, near-UV/blue fluorescence of PAHs is more specific as to size, structure, and charge state of a PAH molecule. If the carriers of this near-UV/blue fluorescence are PAHs, they are most likely neutral PAH molecules consisting of 3-4 aromatic rings such as anthracene (C$_{14}$H$_{10}$) and pyrene (C$_{16}$H$_{10}$). These small PAHs would then be the largest molecules specifically identified in the interstellar medium. ",
        "introduction": "The family of emission bands at 3.3, 6.2, 7.7, 8.6, 11.2, \\& 12.7 $\\mu$m (called the unidentified infra-red (UIR) bands) is found in almost all astrophysical environments including the diffuse interstellar medium (ISM), the edges of molecular clouds, reflection nebulae, young stellar objects, HII regions, star forming regions, some C-rich Wolf-Rayet stars, post-AGB stars, planetary nebulae, novae, normal galaxies, starburst galaxies, most ultra-luminous infra-red galaxies and AGNs \\citep[see][and references therein]{pet04}. Approximately 20-30\\% of the Galactic IR radiation is emitted in these UIR bands and 10-15\\% of the interstellar carbon is contained in the UIR carriers \\citep{sw95}, indicating that the carriers represent an abundant component of the ISM. These UIR bands are the signatures of aromatic C-C and C-H fundamental vibrational and bending modes, and are generally attributed to a family of PAH molecules containing 50-100 carbon atoms \\citep{hony01,cook98,lp84,ver01,atb85,sel84}. Although the presence of PAHs in space is widely accepted  by most astronomers, their specific sizes  and ionization states remain elusive. On absorption of a far-UV photon, a PAH molecule usually undergoes a transition to an upper electronic state. If the molecule undergoes iso-energetic transitions to highly vibrationally excited levels of the ground state, then the molecule relaxes through a series of IR emissions in the C-C and C-H vibrational and bending modes. These transitions are largely independent of size, structure and ionization state of the molecule.  However, \\emph{electronic} fluorescence, a transition from the upper excited level to the ground state, is more specific \\citep{reyle00}. In particular, the wavelength of the first electronic transition seen in neutral PAHs is closely dependent on the size of the molecular species. In general, this fluorescence wavelength increases with the molecular weight of the molecule (Figure 1). An observation of fluorescence in an astronomical source in the UV/visible range offers the possibility of estimating the size of the PAH molecules, which the observation of the UIR band emission in the same source does not. \\begin{figure} \\plotone{f1.eps} \\caption{The first electronic transition in the fluorescence spectra \\citep{berl65,peaden80,reidel1,reidel2} of neutral PAH molecules as a function of their molecular weight. We distinguish between condensed and non-condensed PAHs, with condensed PAHs considered to be more stable under interstellar conditions.\\label{fig1}} \\end{figure} ",
        "conclusions": "Another important factor which should be considered when identifying possible carriers is fluorescence efficiency. Although many different fluorescing species may exist with similar abundances, the spectra of the most efficient ones will be dominant. Laboratory studies have shown that on comparison of emission rates per molecule, from PAH molecules placed at 1 AU from the Sun, anthracene and pyrene show 10-100 times more efficiency than naphthalene (C$_{10}$H$_8$) and phenanthrene (C$_{14}$H$_{10}$) \\citep{brech94}. Additional support for small PAHs as sources of the observed BL is found in the spatial correlation between the BL and the PAH 3.3 $\\mu$m emission in the RR and the presence of the distinct PAH ionization discontinuity in the spectrum of the central source, HD44179 (U. Vijh. A. Witt, \\& K. Gordon, in preparation)."
    },
    "0403/astro-ph0403714_arXiv.txt": {
        "abstract": "The observed number counts of quasars may be explained either by long-lived activity within rare massive hosts, or by short-lived activity within smaller, more common hosts.  It has been argued that quasar lifetimes may therefore be inferred from their clustering length, which determines the typical mass of the quasar host. Here we point out that the relationship between the mass of the black-hole and the circular velocity of its host dark-matter halo is more fundamental to the determination of the clustering length.  In particular, the clustering length observed in the 2dF quasar redshift survey is consistent with the galactic halo -- black-hole relation observed in local galaxies, provided that quasars shine at $\\sim 10$--$100\\%$ of their Eddington luminosity. The slow evolution of the clustering length with redshift inferred in the 2dF quasar survey favors a black-hole mass whose redshift-independent scaling is with halo circular velocity, rather than halo mass. These results are independent from observations of the number counts of bright quasars which may be used to determine the quasar lifetime and its dependence on redshift.  We show that if quasar activity results from galaxy mergers, then the number counts of quasars imply an episodic quasar lifetime that is set by the dynamical time of the host galaxy rather than by the Salpeter time.  Our results imply that as the redshift increases, the central black-holes comprise a larger fraction of their host galaxy mass and the quasar lifetime gets shorter. ",
        "introduction": "The Sloan Digital Sky Survey (York et al.~2000) and the 2dF quasar redshift survey (Croom et al.~2001a) have measured redshifts for large samples of quasars, and determined their luminosity function over a wide range of redshifts (Boyle et al.~2000; Fan et al.~2001a,b; Fan et al.~2003).  The 2dF survey has also been used to constrain the clustering properties of quasars (Croom et al. 2001b). It has been suggested that quasars have clustering statistics similar to optically selected galaxies in the local universe, with a clustering length $R_0\\sim8$Mpc. The large sample size of the 2dF quasars also provided clues about the variation of clustering length with redshift (Croom et al. 2001b) and apparent magnitude (Croom et al.~2002). Quasars appear more clustered at high redshift, although with a relatively mild trend. There is also evidence that more luminous quasars may be more highly clustered.  As pointed out by Martini \\& Weinberg~(2001) and Haiman \\& Hui~(2001), the quasar correlation length determines the typical mass of the dark matter halo in which the quasar resides. One may therefore derive the quasar duty-cycle by comparing the number density of quasars with the density of host dark matter halos. The quasar lifetime then follows from the product of the duty-cycle and the lifetime of the dark-matter halo in between major mergers, although there is a degeneracy between the lifetime and the quasar occupation fraction or beaming.  Preliminary results suggested quasar lifetimes of $t_{\\rm q}\\sim10^6-10^7$ years, consistent with the values determined by other methods (see Martini~2003 for a review), including the transverse proximity effect (Jakobsen et al.~2003) and counting arguments relative to the local population of remnant supermassive black holes (SMBHs; see Yu \\& Tremaine~2002). Kauffmann \\& Haehnelt~(2002) used a detailed semi-analytic model (Kauffmann \\& Haehnelt~2000) to predict the correlation length of quasars and its evolution with redshift. They found that their model reproduces the correct correlation length as well as its redshift evolution for present-day quasar lifetimes of $t_{\\rm q}\\sim10^7$ years.  In this work we argue that the correlation length of quasars is fundamentally determined by the relation between the masses of SMBHs and their host galactic halo rather than by the quasar lifetime. We find that the amplitude of the correlation length depends on the product of the SMBH--halo relation and the typical fraction of the Eddington luminosity at which quasars shine. Moreover, we show that the evolution of the correlation length is sensitive to how the SMBH -- halo relation evolves with redshift, and therefore to the physics of SMBH formation and quasar evolution.  The paper is organized as follows. In \\S~\\ref{cf} and \\S~\\ref{MbhMhalo} we discuss the calculation of the correlation function of quasars with a particular apparent magnitude $B$, and the different scenarios for the evolution of the SMBH -- galaxy halo relation.  In \\S~\\ref{cfobs} we compare fiducial model correlation functions to the results of the 2dF quasar redshift survey (Croom et al.~2001b,c).  We find that the shallow evolution of the correlation length implies that SMBH mass has a redshift-independent scaling with the circular velocity of the host dark matter halo rather than with its mass. The ranges of the normalizations in the SMBH -- galaxy halo relation and of the fraction of Eddington allowed by observations of the observed quasar correlation function are explored in \\S~\\ref{epseta}. In \\S~\\ref{lf} we examine the quasar lifetime by requiring that observational constraints from both the correlation and luminosity functions be satisfied simultaneously. Finally, in \\S~\\ref{disc} we summarize our results and discuss a very simple, physically motivated model that satisfies all constraints with no free parameters. Throughout the paper we adopt the set of cosmological parameters determined by the {\\em Wilkinson Microwave Anisotropy Probe} (WMAP, Spergel et al. 2003), namely mass density parameters of $\\Omega_{m}=0.27$ in matter, $\\Omega_{b}=0.044$ in baryons, $\\Omega_\\Lambda=0.73$ in a cosmological constant, and a Hubble constant of $H_0=71~{\\rm km\\,s^{-1}\\,Mpc^{-1}}$. ",
        "conclusions": "\\label{disc}  In this paper we have demonstrated that the relationship between the mass of a SMBH and its host dark-matter halo is {\\it the} fundamental quantity that determines the clustering of quasars. In particular the $M_{\\rm bh}$--$M_{\\rm halo}$ relation and its evolution with redshift, may be determined directly from the evolution in the correlation length of quasars, independent of consideration of the quasar luminosity function or assumptions about the quasar lifetime.  Beginning with the locally observed $M_{\\rm bh}$--$M_{\\rm halo}$ relation, we have shown that the observed correlation length of quasars in the 2dF quasar redshift survey is consistent with SMBHs that shine at $\\sim 10$--$100\\%$ of their Eddington luminosity during their bright quasar episode. Moreover, the evolution of the clustering length with redshift is consistent with a SMBH mass whose redshift-independent scaling is with the circular velocity of the host dark matter halo (see equation~\\ref{eps}). In contrast, it appears that a relation for the SMBH mass whose redshift-independent scaling is with the mass of the host dark matter halo (equation~\\ref{eps2}) would have resulted in too much evolution of the clustering length with redshift.  The portion of the 2dF quasar survey from which the data used in this study were drawn, contains $\\sim10^4$ quasars. Upon completion, the Sloan Digital Sky Survey will have spectra for ten times this number of quasars, spread over a larger redshift range (York et al.~2000). This will allow more accurate determination of the correlation length and its evolution, and hence provide more accurate constraints the SMBH population. The selection of quasars for the spectroscopic survey of SDSS at $z<3$ is restricted to $i^*<19$. Thus, most of the low redshift SDSS quasars will fall at luminosities that are brighter than the characteristic break in the luminosity function (e.g. Boyle et al.~2000). These bright quasars have a lower abundance than expected from the number of halos merging at $z\\la 2$. The reason for this low abundance may be either that the lifetime (or occupation fraction) of bright quasars is lower than expected at $z\\la 2$, or alternatively that the fraction of the Eddington limit at which the quasars shine is small. The 2dF quasar sample straddles the luminosity function break. We have shown that the similarity in the clustering statistics of sub-samples of quasars having different ranges in luminosity suggest that the lifetime (or occupation fraction) of bright quasars is lower than expected at $z\\la 2$, but that these quasars also shine near their Eddington luminosity. Measurements of the correlation length as a function of quasar luminosity in SDSS can help further distinguish between these two possibilities.    As we have shown in a previous paper (Wyithe \\& Loeb~2003), feedback--regulated growth of SMBHs during active quasar phases implies $M_{\\rm bh}\\propto v_c^5$, consistent with the inferred relation for galactic halos in the local universe (Ferrarese 2002). Assuming $\\eta=1$ as suggested by observations of both low and high-redshift quasars (Floyd~2003; Willott, McLure \\& Jarvis~2003), we have shown that only $\\sim7\\%$ of the Eddington luminosity output needs to be deposited into the surrounding gas in order to unbind it from the host galaxy over the dynamical time of the surrounding galactic disk. The power-law index of 5 in the $M_{\\rm bh}$--$v_c$ relation, is larger than the values of 4--4.5 inferred from the local relation between $M_{\\rm bh}$ and the stellar velocity dispersion $\\sigma_\\star$ (Merritt \\& Ferrarese 2001; Tremaine et al. 2002); the difference originating from the observation that the $v_c$--$\\sigma_\\star$ relation is shallower than linear for stellar bulges embedded in cold dark matter halos (Ferrarese~2002). Feedback--regulated growth also implies that the $M_{\\rm bh}$--$v_{\\rm c}$ relation should be independent of redshift, as implied by the slow evolution in the clustering length of quasars.  A very simple model can therefore be constructed to describe the correlation length of quasars and its evolution in the 2dF survey, as well as the number counts of quasars out to very high redshift ($z\\sim 6$). The model includes four simple assumptions, but no free parameters: {\\it (i)} the locally observed $M_{\\rm bh}$--$M_{\\rm halo}$ relation extends to high redshifts through equation~(\\ref{eps}) with the SMBH mass scaling as halo circular velocity to the 5-th power; {\\it (ii)} SMBHs shine near their Eddington luminosity with a universal spectrum (Elvis et al.~1994) during luminous quasar episodes; (iii) quasar episodes are associated with major galaxy mergers; and (iv) the quasar lifetime is set by the dynamical time of the host galaxy.  The assumption of a SMBH mass that scales with halo circular velocity independently of redshift is supported by the observations of Shields et al.~(2003) that there is no evolution in the $M_{\\rm bh}$--$\\sigma_\\star$ relation out to $z\\sim3$. The assumption that quasars shine near their Eddington limit is supported by observations of high and low redshift quasars (Floyd~2003; Willott, McLure \\& Jarvis~2003).  Overall, the sample of available data on the clustering properties and the number counts of quasars is most readily explained by a quasar lifetime that is set by the dynamical time of the host galaxy rather than by the Salpeter (1964) $e$-folding time for the growth of its mass."
    },
    "0403/astro-ph0403591_arXiv.txt": {
        "abstract": "{ A sample of 229 nearby thick disk stars has been used to investigate the existence of an age-metallicity relation (AMR) in the Galactic thick disk. The results indicate that that there is indeed an age-metallicity relation present in the thick disk. By dividing the stellar sample into sub-groups, separated by 0.1\\,dex in metallicity, we show that the median age decreases by about 5--7\\,Gyr when going from [Fe/H]\\,$\\approx -0.8$ to [Fe/H]\\,$\\approx -0.1$. Combining our results with our newly published $\\alpha$-element trends for a local sample of thick disk stars, that show signatures from supernovae type Ia (SN\\,Ia), we can here draw the conclusion that the time-scale for the peak of the SN\\,Ia rate is of the order 3--4 Gyr in the thick disk. The tentative evidence for a thick disk AMR that we present here also has implications for the thick disk formation scenario; star-formation must have been an ongoing process for several billion years. This is further discussed here and appear to strengthen the hypothesis that the thick disk originates from a merger event with a companion galaxy that puffed up a pre-existing thin disk. ",
        "introduction": "Chemical evolution of stellar populations is an important part of any model of galaxy formation and evolution. Many studies in the past decades show how we are able to further refine and constrain models of Galactic chemical evolution by combining kinematics and elemental abundances of local dwarf stars (e.g. Chiappini et al.~\\cite{chiappini}; Matteucci~\\cite{matteucci}; Edvardsson et al.~\\cite{edvardsson}; Feltzing \\& Gustafsson~\\cite{feltzing1998}; Feltzing et al.~\\cite{feltzing2003}; Bensby et al.~\\cite{bensby}; Reddy et al.~\\cite{reddy}). As evidenced by the cited articles our understanding of chemical evolution is mainly driven by the studies in the solar neighbourhood but have far reaching impact for our interpretation of integrated light from other galaxies (e.g. Matteucci~\\cite{matteucci}).  However, it is not only the elemental abundances and kinematics of the stars that are of importance when we want to further improve the models of galaxy formation and evolution but also the ages of the stars (see e.g. Edvardsson et al.~\\cite{edvardsson}; Raiteri et al.~\\cite{raiteri}; Pilyugin \\& Edmunds~\\cite{pilyugin}; Berczik~\\cite{berczik}).  Many studies have found there to be a clear relation between the ages and the metallicities of the solar neighbourhood disk stars (Twarog~\\cite{twaroga}, \\cite{twarogb}; Rocha-Pinto et al.~\\cite{rochapinto}; Meusinger et al.~\\cite{meusinger}). In contrast to this Edvardsson et al.~(\\cite{edvardsson}) found no particular evidence for an age-metallicity relation in the Galactic disk in the solar neighbourhood and Feltzing et al.~(\\cite{feltzing}) confirmed this. Feltzing et al.~(\\cite{feltzing}) also showed how dangerous selection effects could be and how an artificial age-metallicity relation can be created (see their Figs.~13 and 14).  Gilmore \\& Reid~(\\cite{gilmore}) showed that our galaxy is host to two kinematically distinct disk structures. The ``new'' disk was dubbed the thick disk and was found to have a mean metallicity around $-0.6$\\,dex (Wyse \\& Gilmore~\\cite{wyse}) and a scale-height of 800--1300\\,pc (e.g. Reyl\\'e \\& Robin~\\cite{reyle}; Chen~\\cite{chen2}) while the thin disk has a mean metallicity of around $-0.1$\\,dex and a scale height of 100--300 pc (e.g. Gilmore \\& Reid~\\cite{gilmore}; Robin et al.~\\cite{robin}). Recent studies have shown that stars selected to belong to either the thin or the thick disk show different trends for the elemental abundances (e.g. Fuhrmann~\\cite{fuhrmann}; Feltzing et al.~\\cite{feltzing2003}; Bensby et al.~\\cite{bensby}; Bensby et al.~\\cite{bensby_syre}; Reddy et al.~\\cite{reddy}; Prochaska et al.~\\cite{prochaska}; Mashonkina \\& Gehren~\\cite{mashonkina}).  The question then arises: could it be so that the lack of a relation between ages and metallicities for stars in the solar neighbourhood is in fact a population effect? That is, are we looking at a mixture of stars from (at least) two populations that have different chemical enrichment histories?  It thus appears natural to, yet again, revisit the question of an age-metallicity relation in the solar neighbourhood. In the study presented here we will address the question of a relation between ages and metallicities for stars that are kinematically selected to resemble the thick disk closely.  The paper is organized as follows: in Sect.~\\ref{sect:kin} we describe the stellar sample and the kinematical selection criteria and investigates if there are potential biases present in the sample. In Sect.~\\ref{sect:alpha} we describe the choice of $\\alpha$-enhancement used in the isochrones when deriving the stellar ages. In Sect.~\\ref{sec:amr} we derive ages from stellar isochrone fitting and find that there is a possible age-metallicity relation present in the thick disk. In Sect.~\\ref{sec:discussion} we discuss the implications this tentative age-metallicity relation have on the star-formation history of the thick disk, on the time-scale of SN\\,Ia rate in the thick disk, and on our understanding of the origin and evolution of the thick disk. Finally, in Sect.~\\ref{sec:summary} we give a short summary. ",
        "conclusions": " \\begin{enumerate} \\item \tThere is an age-metallicity relation present in the Galactic thick disk, indicating that it has had an ongoing star formation for a time-period of up to 5\\,Gyr (this is model dependent). \\item\tThe thick disk age-metallicity relation in combination with the abundance trends for $\\alpha$-elements in the thick disk, that show signatures from SN\\,Ia, indicate that the time-scale for the peak of the SN\\,Ia rate in the thick disk is of the order 3--4\\,Gyr. \\item\tThe quite long star formation period in the thick disk strengthens the hypothesis that the thick disk formed as a result of an ancient merger event between the Milky Way and a companion galaxy. \\end{enumerate}  Studies of the thick disk using nearby stars will always be be subject to uncertainties due to the overlapping velocity-, and metallicity distributions of the thin and thick disk. At distances well above $Z\\approx1.5$\\,kpc from the Galactic plane the thick disk is the dominant stellar population. Determining accurate metallicities and $\\alpha$-abundances for a larger sample of dwarf stars at high $Z$ would therefore enable accurate age determinations that would verify the existence (or non-existence) of the thick disk AMR deduced from nearby stars."
    },
    "0403/astro-ph0403558_arXiv.txt": {
        "abstract": "It has been conjectured that the distribution of magnifications of a point source microlensed by a randomly distributed population of intervening point masses is independent of its mass spectrum. We present {\\it gedanken} experiments that cast doubt on this conjecture and numerical simulations that show it to be false. ",
        "introduction": "Every  investigation of microlensing  at high  optical depth  that has explored the effect of multiple microlens mass components  has led to the  conclusion  that the  magnification  probability distribution  is independent of  the spectrum of  microlens masses.  The recent effort by \\citet{2001MNRAS.320...21W} is typical.  While it was not their principal result, they comment in passing \\begin{quotation} ``...   we  confirm   the  finding  of  \\citet{1992ApJ...386...19W}  and \\citet{1995MNRAS.276..103L}  that  the  magnification distribution  is independent of the mass function.'' \\end{quotation} This conjecture has important consequences regarding the more general applicability of microlensing studies that are limited to a single mass component.  While galaxies have stars with a range of masses, restricting to a single component makes analytic calculations more tractable \\citep[e.g.][]{ 1986MNRAS.223..113P,1987ApJ...319....9S,1997ApJ...489..508K} and greatly decreases the number of cases that must be simulated numerically \\citep[e.g.][]{1992ApJ...386...19W,1995MNRAS.276..103L,2001MNRAS.320...21W}. If true, the conjecture simplifies things considerably.  Both theoretical and experimental lines of evidence lead to this conclusion, which has struck many investigators as obvious.  On the experimental side, simulations like those carried out by \\citet{2001MNRAS.320...21W} and their predecessors produce magnification histograms for different mass distributions that appear to be indistinguishable for fixed surface mass density and shear.  On  the   theoretical  side,  the  high  magnification   tail  of  the magnification   probability  distribution   has  been   shown   to  be independent       of      the       microlens       mass      spectrum \\citep{1987ApJ...319....9S}.    Moreover,  \\citet{1992A&A...258..591W} showed that the average  number of positive parity microimages depends only upon  the surface mass density (or equivalently the convergence) and the shear.   Since the scale free nature  of gravity requires that the magnification  probability distribution for a point  source be the same for  microlensing by a single  mass of any size,  it would appear strange if a mixture of two masses (at constant convergence and shear) produced a different magnification probability distribution.  There is, however, at least one argument against this apparently obvious conclusion, which we detail in $\\S$ 2 below.  It suggests that the magnification probability distribution {\\it does} depend upon the mass spectrum.  The argument suggests that the dependence would show up in a highly magnified negative parity macroimage -- typically one of a close pair of images in a quadruply imaged quasar like PG1115+080.  We have carried out lensing simulations of such an image (at constant convergence and shear) for a variety of different cases.  In Figure 1 we show simulations with two populations of point masses.  The first component is comprised of $1.000 \\msun$ objects referred to hereafter as ``micro-lenses.''  The second component is comprised of $0.005 \\msun$ objects referred to hereafter as ``nano-lenses.''  The designations and mass scale are arbitrary but are intended to convey the sense that the micro-lenses are very much smaller than the lensing galaxy and that the nano-lenses are very much smaller than the micro-lenses.  The eight panels of Figure 1 show magnification histograms obtained by varying the mass fractions in the micro-lensing component, with the remaining fraction in the nano-lensing component.  For the sake of comparison, we reproduce in each panel the result for a pure micro-lensing component.  As the fraction contributed by micro-lenses decreases to 20\\% and 10\\% the histogram broadens out and develops a second peak.  But as it decreases further to 0\\%, the magnification distribution narrows and ends up looking like the 100\\% case (modulo finite source effects and sample variance).  Unless our simulations are faulty, the conjecture is false.  In $\\S$ 2 we put forward  a qualitative argument for the dependence of the microlensing  probability distribution  on the mass  spectrum.  In $\\S$ 3 we  give details of the numerical  simulations that confirm the effect.  In $\\S$ 4 we offer a qualitative interpretation of our results.  In $\\S$ 5 we discuss some astrophysical consequences. ",
        "conclusions": ""
    },
    "0403/astro-ph0403244_arXiv.txt": {
        "abstract": "Young star clusters with masses similar to those of classical old globular clusters are observed not only in starbursts, mergers or otherwise disturbed galaxies, but also in normal spiral galaxies. Some young clusters with masses as high as $\\sim10^6 \\, \\msun$ have been found in the disks of isolated spirals. Dynamical mass estimates are available for a few of these clusters and are consistent with Kroupa-type IMFs.  The luminosity (and possibly mass-) functions of young clusters are usually well approximated by power-laws. Thus, massive clusters at the tail of the distribution are naturally rare, but appear to be present whenever clusters form in large numbers.  While bound star clusters may generally form with a higher efficiency in environments of high star formation rate, many of the apparent differences between clusters in starbursts and ``normal'' galaxies might be simply due to sampling effects. ",
        "introduction": "It may be worth recalling some of the main properties of the Milky Way open cluster system. The census of open clusters is still highly incomplete beyond distances of a few kpc from the Sun, although the situation is improving with new surveys such as 2MASS (see e.g.\\ the contributions by Carpenter and Hanson in this volume). The luminosity function of Milky Way open clusters was analysed by \\citet{van84}, who found it to be well modelled by a power-law $N(L)dL \\propto L^{-1.5}dL$ over the range $-8 < M_V < -3$. However, they also noted that extrapolation of this luminosity function would predict about 100 clusters as bright as $M_V=-11$ in the Galaxy, clearly at odds with observations, and thus suggested some flattening of the LF slope at higher luminosities.  The brightest known young clusters (e.g.\\ NGC~3603, $h$ and $\\chi$ Per) have absolute $V$ magnitudes of $M_V\\sim-10$, corresponding to total masses of several thousand \\msun . Recently, there have been claims that the Cyg OB2 association might be an even more massive cluster \\citep{kno00}, but this object is probably too diffuse to be a bound star cluster (though it does have a compact core). There are, however, a number of old ($>1$ Gyr) open clusters in the Milky Way with masses of $\\sim10^4$ \\msun\\ \\citep{friel95}. These objects are likely to have lost a significant fraction of their total mass over their lifetimes, and may thus originally have been even more massive. They serve to illustrate that, even in the Milky Way, the distinction between globular and open clusters is not always clear-cut.  It has been recognized for about a century that the Magellanic Clouds, and the LMC in particular, host a number of ``blue globular clusters'' \\citep{shap30}. % Among the most massive of these is NGC~1866, with a mass of around $10^5 \\, \\msun$ and an age of $\\approx100$ Myr \\citep{fis92,van99}. An older example is NGC~1978 with similar mass but an age of 2--3 Gyr, clearly demonstrating that at least some such clusters can survive for several Gyrs.  The interaction between the LMC and the Milky Way has probably affected the star formation history of the LMC, which is known to be bursty with major peaks in the star formation rate correlating with perigalactic passages \\citep{smeck02}. One might argue, then, that the formation of YMCs in the LMC could be induced by interaction with the Milky Way. However, the LMC is not the only example even in the Local Group of a galaxy that hosts YMCs.  Another well-known example is M33, which does not display evidence for a bursty cluster (and, presumably star-) formation history \\citep{cs82,cs88}.  \\citet{chan99,chan01} have identified many more star clusters in this galaxy, though not all are particularly massive.  With the launch of HST it became possible to investigate more crowded and/or distant systems in detail and attention started to shift towards more extreme starbursts, including a large number of merger galaxies (e.g.\\ Whitmore, this volume). It is now clear that luminous, young star clusters often form in very large numbers in such galaxies, and this has led to suggestions that formation of ``massive'' star clusters might require special conditions such as large-scale cloud-cloud collisitions \\citep{js92}.  However, the question remains to be answered why some non-interacting galaxies also contain YMCs, whereas apparently the Milky Way does not. YMCs are now being found in an increasing number of non-interacting galaxies, posing a severe challenge for formation scenarios which require special conditions. ",
        "conclusions": "It is becoming increasingly clear that ``massive'' star clusters can form in a wide variety of galaxies, and not just in mergers or otherwise disturbed galaxies. With the possible exception of some dwarf galaxies, the luminosity distributions of young star clusters generally appear to be power-laws. If cluster luminosities are sampled at random from a power-law distribution, the most luminous clusters will naturally be rare, but so far there is no evidence for a statistically significant upper cut-off. In other words, very luminous (and massive) clusters appear to form whenever clusters form in large numbers.  % This is illustrated by the fact that young star clusters with masses up to $\\sim10^6 \\, \\msun$ have been identified in the disks of several apparently normal, isolated spiral galaxies with rich cluster systems.  These galaxies, such as NGC~5236, NGC~6946 are characterised by high star formation rates, but these do not generally appear to be triggered by interactions with other galaxies. Dynamical mass estimates are now available for a small number of these clusters, and the mass-to-light ratios are compatible with standard Kroupa-type IMFs. There is every reason to be optimistic that important clues to the formation of classical globular clusters may be obtained by studying their younger counterparts in the Local Universe."
    },
    "0403/astro-ph0403134_arXiv.txt": {
        "abstract": "{We discuss the lifetimes and evolution of dense cores formed as turbulent density fluctuations in magnetized, isothermal molecular clouds. We consider numerical simulations in which we measure the cores' magnetic criticality and Jeans stability in relation to the magnetic criticality of their ``parent clouds'' (the numerical boxes). In subcritical boxes, dense cores do not form, and collapse does not occur. In supercritical boxes, some cores collapse, being part of larger clumps that are supercritical from the start, and whose central, densest regions (the cores) are initially subcritical, but rapidly become supercritical, presumably by accretion along field lines. Numerical artifacts are ruled out. The time scales for cores to go from subcritical to supercritical and then collapse are a few times the core free-fall time, $\\tfc$. Our results suggest that cores are out-of-equilibrium, transient structures, rather than quasi-magnetostatic configurations. }  \\resumen{Discutimos los tiempos de vida y la evoluci\\'on de los n\\'ucleos densos formados como fluctuaciones turbulentas de densidad en nubes moleculares isot\\'ermicas magnetizadas. Consideramos simulaciones num\\'ericas en las que medimos la criticalidad magn\\'etica y la estabilidad de Jeans de los n\\'ucleos, en relaci\\'on a la criticalidad magn\\'etica de sus nubes ``madre'' (los dominios de integraci\\'on num\\'erica). En dominios subcr\\'iticos, no se forman n\\'ucleos densos, y no ocurre colapso gravitacional. En dominios supercr\\'iticos, algunos n\\'ucleos se colapsan, formando parte de grumos m\\'as grandes que son supercr\\'iticos desde el inicio, y cuyas regiones m\\'as densas (los n\\'ucleos) son inicialmente subcr\\'iticas, pero r\\'apidamente se tornan supercr\\'iticos, presumiblemente por acreci\\'on a lo largo de las l\\'ineas de campo magn\\'etico. Descartamos la posibilidad de artefactos num\\'ericos. Las escalas de tiempo en las que los n\\'ucleos transitan desde el estado subcr\\'itico, pasando por el supercr\\'itico, hasta finalmente el colapso, son de unas cuantas veces su tiempo de ca\\'ida libre, $\\tfc$. Nuestros resultados sugieren que los n\\'ucleos son estructuras transientes, fuera de equilibrio, y no configuraciones cuasi-magneto-hidrost\\'aticas. }   \\addkeyword{ISM: Molecular clouds} \\addkeyword{ISM: Turbulence} \\addkeyword{Stars: Formation}   \\begin{document} ",
        "introduction": "\\label{sec:intro}  The prevailing view (which we hereafter refer to as the ``standard (magnetic support) model'' of star formation; see, e.g., the reviews by Shu, Adams \\& Lizano 1987; McKee et al.\\ 1993) concerning low-mass-star-forming clumps is that they are quasi-static equilibrium configurations with so-called ``subcritical'' mass-to-magnetic-flux ratios, so that the clumps are supported against their self-gravity by the magnetic field in the direction perpendicular to it, and by a combination of thermal and turbulent pressures along the field. Under ideal MHD conditions, the magnetic field is ``frozen'' into the plasma, and the magnetic flux is conserved. Under the additional assumption that the clump's mass is also constant, then the mass-to-flux ratio is a fixed parameter of the clump, which therefore cannot collapse if this ratio is subcritical (meaning that the core's self-gravity is never enough to overwhelm the magnetic support). However, because the cold molecular gas is only partially ionized, the process known as ambipolar diffusion causes a loss of magnetic flux from the clumps on time scales long compared to their free-fall time, allowing them to contract and form denser cores that will ultimately collapse.  However, it well known that molecular clouds are supersonically turbulent (e.g., Larson 1981; Blitz \\& Williams 1999), and it is becoming increasingly accepted that the cores within them are the density fluctuations induced by the turbulence (\\BP, \\VS\\ \\& Scalo 1999; see also the reviews by \\VS\\ et al.\\ 2000; Mac Low \\& Klessen 2004). In this context, the cores have a highly dynamical origin (supersonic compressions), and their masses are hardly fixed. Moreover, it is natural to ask whether they can settle into hydrostatic equilibria, an event which requires the equilibria to be stable (or ``attracting'', in the language of nonlinear phenomena). Otherwise, the dynamic density fluctuations will just ``fly past'' the equilibrium state on their way to collapse, or else ``rebound'' and merge back with their environment, if they do not quite reach the equilibrium point. In this case, the cores' lifetimes should be much shorter than in the standard model, probably comparable to their free-fall times. In the present contribution, we argue in favor of this scenario. To this end, we discuss the formation and evolutionary time scales of cores that form in numerical simulations of isothermal, compressible MHD turbulence, in relation to the magnetic criticality of the whole computational box. Here we present a brief overview. For a full, detailed discussion, see \\VS\\ et al.\\ (2004). ",
        "conclusions": "\\label{sec:conclusions}  In this contribution, we have presented a discussion and numerical simulations of the formation, nature, and lifetimes of dense cores in magnetized clouds. We argued that the mass-to-flux ratio of a cloud puts an upper bound to that of any clump or core within it, and so the only way to form supercritical cores is within supercritical clouds (neglecting ambipolar diffusion). Moreover, numerical simulations of marginally subcritical clouds show that no gravitationally bound cores form in this case, while in supercritical clouds the gravitationally bound cores that form occur inside clumps that are supercritical, and rapidly become supercritical themselves and collapse, in timescales that do not exceed twice their local free-fall time $\\tfc$. We also showed that the mass-to-flux ratio of structures defined through a density threshold (as would be the case of cores observed in a single molecular line, which requires the density to be larger than a certain value in order to be excited) does not remain constant, because the core is continuously connected to its parent structure, from which it can accrete mass. Our results support the notion that clumps and cores are out-of-equilibrium, transient structures, and that a class of ``failed'' cores should exist, that will not form stars."
    },
    "0403/physics0403086_arXiv.txt": {
        "abstract": "The issue of asymmetric uncertainties resulting from fits,  nonlinear propagation and systematic effects is reviewed. It is shown that, in all cases, whenever a published result is given with asymmetric uncertainties, the {\\it value} of the physical quantity of interest {\\it  is biased} with respect to what would be obtained using at best all experimental and theoretical information that contribute to evaluate the combined  uncertainty. The probabilistic solution to the problem is provided both in exact and in approximated forms. ",
        "introduction": "We often see published results in the form $$\\mbox{`best value'}\\ ^{+\\Delta_+}_{-\\Delta_-}\\,,$$ where $\\Delta_+$ and $\\Delta_-$ are {\\it usually} positive.\\footnote{For examples of measurements having $\\Delta_+$ and $\\Delta_-$ with all combinations of signs, see public online tables of Deep Inelastic Scattering results.\\cite{HERA} I want to make clear since the very beginning that it is not my intention to blame experimental or theoretical teams which have reported in the past asymmetric uncertainty, because we are all victims of a bad tradition in data analysis. At least, when asymmetric uncertainties have been given, there is some chance to correct the result, as described in Sec.~\\ref{sec:thumb}. Since some asymmetric contributions to the global uncertainties almost unavoidably happen  in complex experiments, I am more worried of collaborations that never arrive to final asymmetric uncertainties, because I must imagine they have symmetrised somehow the result but, I am afraid, without applying the proper shifts to the `best value' to take into account asymmetric contributions, as it will be discussed in the present paper.} As firstly pointed out in Ref.~\\cite{ConMirko} and discussed in a simpler but more comprehensive way in Ref.~\\cite{BR}, this practice is far from being acceptable and, indeed, could bias the believed value of important physics quantities. The purpose of the present paper is, summarizing and somewhat completing the work done in the above references, to remind where asymmetric uncertainty stem from and to show why, as they are usually treated, they bias the value of physical quantities, either in the published result itself or in subsequent analyses. Once the problems are spotted, the remedy is straightforward, at least within the Bayesian framework (see e.g. \\cite{BR}, or \\cite{RPP_GdA} and \\cite{RPP_Dose} for recent reviews). In fact the Bayesian approach is conceptually based on the intuitive idea of probability, and formally grounded on the basic rules of probability (what are usually known as the probability `axioms' and the `conditional probability definition') plus logic. Within this framework many methods of `conventional' statistics are reobtained, as approximations of general solutions, under well stated conditions of validity. Instead, in the conventional, frequentistic approach {\\it ad hoc} formulae, prescriptions and un-needed principles are used, often without understanding what is behind these methods -- before a `principle' there is nothing!  The proposed Bayesian solutions to cure the troubles produced by the usual treatment of asymmetric uncertainties is to step up from approximated methods to the more general ones (see e.g. Ref.~\\cite{BR}, in particular the top down approximation diagram of Fig.~2.2). In this paper we shall see, for example, how $\\chi^2$ and minus log-likelihood fit `rules' can be derived from the Bayesian inference formulae as approximated methods and what to do when the underlying conditions do not hold. We shall encounter a similar situation regarding standard formulae to propagate uncertainty.  Some of the issues addressed here and in Refs. \\cite{ConMirko} and \\cite{BR} have been recently brought to our attention by Roger Barlow~\\cite{Barlow}, who proposes frequentistic ways out. Michael Schmelling had also addressed questions related to `asymmetric errors', particularly related to the issue of weighted averages~\\cite{Schmelling}. The reader is encouraged to read also these references to form his/her idea about the spotted problems and the proposed solutions.  In Sec. \\ref{sec:propagation} the issue of propagation of uncertainty is briefly reviewed at an elementary level (just focusing on the sum of uncertain independent variables -- i.e. no correlations considered) though taking into account asymmetry in probability density functions (p.d.f.) of the {\\it input} quantities. In this way we understand what `might have been done' (we are rarely in the positions to exactly know ``what has been done'') by the authors who publish asymmetric results and what is the danger of improper use of such a published `best value' -- {\\it as is} -- in subsequent analyses. Then, Sec.~\\ref{sec:sources} we shall see in where asymmetric uncertainties stem from and what to do in order to overcome their potential troubles. This will be done in an exact way and, whenever is possible, in an approximated way. Some rules of thumb to roughly recover  sensible probabilistic quantities (expected value and standard deviation) from results published with asymmetric uncertainties will be given in Sec.~\\ref{sec:thumb}. Finally, some conclusions will be drawn. ",
        "conclusions": "Asymmetric uncertainties do exist and there is {\\it no way to  remove them artificially}. If they are not properly treated, i.e. using prescriptions that do not have a theoretical ground but are more or less rooted in the physics community, the published result is biased. Instead, if they are properly treated using probability theory, in most cases of interest the final result is practically symmetric and approximately Gaussian, with expected value and standard deviations which take into account the several shifts due to individual asymmetric contributions. Note that some of the simplified methods to make statistical analyses had a {\\it raison d'\\^etre}  many years ago, when the computation was a serious limitation. Now it is not any longer a problem  to evaluate, analytically or numerically, integrals of the kind of those appearing e.g. in Eqs.(\\ref{eq:prop_general}), (\\ref{eq:theta_E}) and  (\\ref{eq:theta_sigma}).  In the case the final uncertainty remains asymmetric, the authors should provide detailed information about the `shape of the uncertainty', giving also most probable value, probability intervals, and so on. But the best estimate of the {\\it expected value and standard deviation should be always given} (see also the {\\it ISO Guide}~\\cite{ISO}).  To conclude, I would like to leave the final word to my preferred quotation with whom I like to end seminars and courses on probability theory applied to the evaluation and the expression of uncertainty in measurements: \\begin{quote} {\\sl \\small ``Although this {\\rm Guide} provides a framework for assessing uncertainty, it cannot substitute for critical thinking, intellectual honesty, and professional skill. The evaluation of uncertainty is neither a routine task nor a purely mathematical one; it depends on detailed knowledge of the nature of the measurand and of the measurement. The quality and utility of the uncertainty quoted for the result of a measurement therefore ultimately depend on the understanding, critical analysis, and integrity of those who contribute to the assignment of its value.''}\\cite{ISO} \\end{quote}  \\vspace{1.0cm} \\noindent It is a pleasure to thank {\\it Superfaber} (Fabrizio Fabbri in {\\it hepnames}) for helpful discussions on the subject and for his {\\it super}vision of the manuscript."
    },
    "0403/astro-ph0403302_arXiv.txt": {
        "abstract": "{ We present high resolution optical spectroscopy of three candidate members of the Taurus-Auriga star forming region. Based on the spectral type, the strength, profile and width of the H$\\alpha$ line, the detected lithium at 6708 \\AA, the location of these objects in a H-R diagram and the comparison with similar objects belonging to young stellar associations, we determine that they are bona fide members of the SFR, with about $\\sim$3 Myr, have masses at or below the substellar limit and, at least in one case, there is active accretion from a circum(sub)stellar disk. This result suggests that high mass brown dwarfs go through a Classical TTauri phase and form like stars, from colapse and fragmentation of a molecular cloud. ",
        "introduction": "Brown dwarfs, objects unable to fuse hydrogen in an stable manner (i.e., with masses below at about 0.072 M$_\\odot$, Baraffe et al. 1998), pose an important problem to the theory of stellar formation. Several formation mechanisms have been proposed, including formation like a star (from collapse and fragmentation of a molecular cloud, Padoan P., \\& Nordlund 2004) to a planet-like process (from a circumstellar disk) or as stellar ``embryos'', ejected from multiple systems before they are able to accrete enough matter (Reipurth \\& Clarke 2001; Bate et al. 2002). These proposed mechanisms have different implication in their formation, their evolution and their properties. In particular, if brown dwarfs are  created like stars, it seems that they should go through an phase of active accretion  from a circum(sub)stellar disk, such as  low mass stars do (Shu et al$.$ 1987).  Actually, since the last couple of years, different groups have presented indirect and direct evidences of active accretion in a handful of young brown dwarfs belonging to several nearby star forming regions (SFR), open clusters and moving groups (Fern\\'andez \\& Comer\\'on 2001; Muench et al. 2001; Natta \\& Testi 2001; Natta et al. 2001, 2002; Testi et al. 2002; Jayawardhana et al. 2002ab, 2003ab; Muzerolle et al. 2003; Barrado y Navascu\\'es et al. 2002, 2003, 2004; Barrado y Navascu\\'es \\& Mart\\'{\\i}n 2003; Mohanty \\& Basri 2003; Mohanty, Jayawardhana \\& Barrado y Navascu\\'es 2003; Comer\\'on et al. 2003)   In this paper, we enlarge the sample by collecting high resolution spectroscopy of three proposed members of the Taurus-Auriga SFR, located at 140 pc and about 1-3 Myr (for an update on Taurus, see Luhman et al. 2003 and references therein). ",
        "conclusions": "\\subsection{Spectral types}  Spectral types were derived by comparing with several spectral templates, by  using the order around 7050 \\AA, corresponding to a TiO band, following Mohanty et al. (2004), since this range is very sensitive to effective temperature. Errors can be estimated as half a subclass. Our values are very close to those obtained by Brice\\~no et al. (2002) from low resolution spectra (see Table 1).        \\setcounter{figure}{0} \\begin{figure} \\centering \\includegraphics[width=8.2cm]{barradoF1.ps} \\caption{ Spectra around HeI6678 \\AA{} and LiI6707.8 \\AA. } \\end{figure}   \\subsection{On the lithium abundance, age, and mass}  We have detected lithium 6707.8 \\AA{} in two out of the three targets (KPNO-Tau-05 and KPNO-Tau-08, Figure 1). The spectrum of KPNO-Tau-03 has worse quality. With some caveats, the visual inspection indicates that this feature is also present. Note that   Brice\\~no et al. (2002)  states that they detected lithium  with their low resolution spectrum (at higher signal-to-noise ratio). This element is easily destroyed in the stellar interior, being its surface abundance dependent on mass and age (as well as other second order parameters). In fact, brown dwarfs which are more massive than about 0.060 M$_\\odot$ do deplete lithium in a time scale of few tens of million years (D'Antona \\& Mazzitelli 1994; Burrows et al. 1997; Baraffe et al. 1998). The Taurus-Auriga complex has an age between 1  and 3 Myr. Since our three candidate members have spectral types between M6 and M7, they should have masses equal or larger than 0.06 M$_\\odot$, according to Baraffe et al. (1998). Therefore, the detection of this alkali clearly indicates that these objects are in the  pre-Main Sequence (PMS). Moreover, from the statistical point of view, the likelihood of having three late-M, PMS interloper whose spectral and photometric properties coincide with the Taurus sequence is negligible. Therefore, we have to conclude that they, indeed, belong to the association.         \\setcounter{figure}{1} \\begin{figure} \\centering \\includegraphics[width=8.2cm]{barradoF2.ps} \\caption{ Lithium equivalent width versus the spectral type. The solid line corresponds to the upper envelope of the values measured in young open clusters. The long-dashed line delimits the areas for weak-line and post-T~Tauri stars (adapted from Mart\\'{\\i}n 1997 and Mart\\'{\\i}n \\& Magazz\\`u 1999). Short-dashed and dotted lines correspond to the cosmic abundances --A(Li)=3.1-- from gravities of Logg=4.5 and 4.0, respectively (curves of growth from Zapatero Osorio et al$.$ 2002). All Pleiades and IC2391  members with measured lithium equivalent width are shown as open circles and triangles --upper limits--. Note the lithium depletion boundary at  $\\sim$M5.5. Sigma Orionis low mass stars and brown dwarfs appear as solid squares (Zapatero Osorio et al$.$ 2002). The big asterisks represent KPNO-Tau-3, 5 and 8. } \\end{figure}  Figure 2 displays values of the lithium equivalent width (W) measured for our targets (large asterisks), pre-main sequence members of the $\\sim$5 Myr Sigma Orionis cluster (Zapatero Osorio et al. 2002, solid squares), and stellar and substellar members of IC2391 --53 Myr-- and the Pleiades --125 Myr-- from Barrado y Navascu\\'es et al. (1999, 2004), Soderblom et al. (1993), Garc\\'{\\i}a-L\\'opez et al. (1994), Jones et al. 1996, Stauffer et al. (1998) and Jeffries et al. (1999). The solid line describe the maximum W(Li) measured in cluster stars. The long-dashed curve is an update version of the criterion used by Mart\\'{\\i}n (1997) and Mart\\'{\\i}n \\& Magazz\\`u (1999) to distinguish between Weak-line and post- TTauri stars. Finally, the dotted and short dashed lines correspond to the Zapatero Osorio et al. (2002) curves of growth for Log\\,g=4.0 and 4.5 and an abundance of A(Li)=3.1 (i.e., cosmic abundance). We have labeled in the plot different sections. The W(Li) for KPNO-Tau-5 is lower than the measured values in the other two Taurus objects. This fact might be due to the presence of optical veiling, about r(6700)$\\sim$0.2, although we do not think this is the case. On one hand, no other sign of accretion has been found in this object (see next section). Moreover, this equivalent width  is compatible with the dispersion found in Sigma Orionis and the  --undepleted- values characteristic of cluster brown dwarfs. For the other two objects, no veiling seems to be present either. In the lithium equivalent width dispersion is real both in Taurus and Sigma Orionis cluster (the same effect is present in the Lambda Orionis association, Barrado y Navascu\\'es, Stauffer \\& Jayawardhana 2004), it might imply diferences in the surface abundances due to additional mixing mechanisms in the stellar and substellar interior, different to pure convection (such as those proposed in solar mass stars), differences in the structure due to different stellar and substellar parameters (such as fast rotation in the case of pre-Main sequence stars, Mart\\'{\\i}n \\& Claret  1996), or to induced effects related to  activity and/or rotation over the spectral feature itself or the surrounding continuum (see the discusion in Barrado y Navascu\\'es et al. 2001 and references therein).  As a summary, based on these data, mainly from the detection  of lithium and accretion when present (see next section), we can conclude that these three objects are young and, indeed, belong to the Taurus region. In the case of KPNO-Tau-5, due to its spectral type (M7) its substellar nature is well established. The other two are located at the substellar borderline and their nature is not so firmly established due to uncertainties in the models and the spectral type determination.  Once membership to the stellar association has been proved, we have derived the bolometric magnitudes from $Ic$ and $Ks$, the distance modulus $(m-M)_0$=5.731 (140 pc, Kenyon, Dobrzycka \\& Hartmann 1994), the reddening derived by Brice\\~no et al. (2002) and the bolometric corrections by Comer\\'on et al. (2000) and Tinney et al. (1993) for these two bands ($Ic$ and $Ks$, respectively). Masses were computed based on models by Baraffe et al. (1998). Effective temperatures were obtained using several scales, namely Luhman (1999) for intermediate gravity and Leggett (2000, 2001). All the measured and derived values are listed in Table 1.        \\setcounter{figure}{2} \\begin{figure} \\centering \\includegraphics[width=8.2cm]{barradoF3.ps} \\caption{HR diagram of KPNO-Tau-3 and 8 (on top of each other) and KPNO-Tau-5. The values represented by asterisks and open stars were derived from $Ic$ and $Ks$ magnitudes, respectively. } \\end{figure}   Figure 3 displays a HR diagram. The isochrones and evolutionary tracks --solid and dashed lines, respectively-- are from Baraffe et al. (1998). Asterisks and  open stars correspond to bolometric luminosities obtained from $Ic$ and $Ks$, respectively. For this particular diagram, we made use of the effective temperature  scale by Luhman (1999). Other temperature scales, such as that from Leggett et al. (2001), would shift the location of these three objects (two of them are almost on top of each other, the small shift in Teff is arbitrary) to the right hand-side, making them younger and less massive. Note, however, that Luhman' scale has been tunned specifically for the model we are using here. In any case, regardless the election of models (alternative models are, for instance, Burrows et al. 1997; D'Antona \\& Mazzitelli 1994, 1997, 1998; Chabrier  et al. 2000; Baraffe et al. 2002), bolometric corrections and effective temperature scale, these members are at or below the substellar frontier and have an age of about 3 Myr.        \\setcounter{figure}{3} \\begin{figure} \\centering \\includegraphics[width=8.2cm]{barradoF4.ps} \\caption{Color-color diagrams. The position of the Taurus BD candidates are indicated with large asterisks. Crosses indicate the position of classical TTauri stars belonging to Orion stellar population (Herbig \\& Bell 1988). The thick-solid and dashed lines correspond to the locii of the main sequence stars (from Bessell \\& Brett 1988; Kirkpatrick et al$.$ 2000; Leggett et al. 2001) and CTT stars (Meyer et al$.$ 1997, Barrado y Navascu\\'es et al. 2003), respectively.   } \\end{figure}     \\subsection{Near infrared photometry, H$\\alpha$ emission and accretion}  None of our three targets   seems to have neither the presence of forbidden lines, characteristics of outflows, nor  near infrared infrared excesses, coming from an accretion disk. Figure 4 compares the colors $(Ic-J)$ and $(H-Ks)$, and includes  the loci for dwarfs --Bessell \\& Brett 1988; Kirkpatrick et al. 2000; Leggett et al. 2001)-- and Classical TTauri stars (Meyer et al. 1997, Barrado y Navascu\\'es et al. 2003) and Classical TTauri stars from Orion (Herbig \\& Bell 1988). As can be seen, the Taurus objects have photometric properties --data from 2MASS-- similar to those MS stars of similar spectral type and, as stated in the previous paragraph, no near IR excess is seen. This fact, by itself, cannot  prove or disprove the presence of a circum(sub)stellar disk, since a hole can be present  or the disk temperature can be very cold. Moreover, the IR excess might depend on the orientation of the disk. See, for example, the case of LS-RCrA~1, M6.5 brown dwarf belonging to RCrA dark cloud (Barrado y Navascu\\'es, Mohanty \\& Jayawardhana 2004). Additional observations and longer wavelengths --more sensitive to cooler disks- both from ground-based telescopes and space-borne instruments such as those in Spitzer Space Telescope, can help to shed some light in this issue.     \\setcounter{figure}{4} \\begin{figure} \\centering \\includegraphics[width=8.2cm]{barradoF5.ps} \\caption{ H$\\alpha$ profiles. The instrumental profile is included as a dashed line (top spectrum). } \\end{figure}   We have detected and measured H$\\alpha$ in emission for these three objects. The profiles of this feature are displayed in Figure 5. Note the possible asymmetry in KPNO-Tau-8, the double peak in KPNO-Tau-5 and KPNO-Tau-3, and the width of the line (310 km/s), typical of accreting brown dwarfs (White \\& Basri 2003; Jayawardhana et al. 2003). In fact, this last object is above the criterion defined by Barrado y Navascu\\'es \\& Mart\\'{\\i}n (2003) which discriminate between accreting and non accreting objects, as Figure 6 clearly indicates. We note, however, that this criterion depends on low resolution spectroscopy, which normally yields larger equivalent widths compared with higher resolution data. In this diagram, we have included Classical and Weak-line TTauri stars belonging to Taurus as solid and open circles. Members without classification are included as crosses. Our three targets appear as large asterisks. Overlapping squares and big circles denote those members with forbidden emission lines and near-infrared excesses, respectively. This large W(H$\\alpha$) agree with the fact that we also detect HeI6678 \\AA, another accretor indicator (see Figure 1).  \\setcounter{figure}{5} \\begin{figure} \\centering \\includegraphics[width=8.2cm]{barradoF6.ps} \\caption{H$\\alpha$ equivalent widths for members of the Taurus SFR. Solid  and open circles correspond to Classical and TTauri stars. Objects with no classification are shown as crosses. Large open circles or square  represent objects with mid-IR excesses and forbidden lines, respectively. The three objects studied here are displayed as large asterisks. The dotted, bold curve is the saturation  criterion, whereas two previously proposed criteria (5 and 20 \\AA) to separate CTTS and WTTS are included as long-dashed, thin horizontal segments. The vertical dotted segment denotes the location of the substellar frontier. } \\end{figure}    Active accretion in at least one Taurus member whose mass is close or below the substellar limit is important for several reasons. First, although the sample studied in this paper is very small, it suggests that a significant fraction of the very low mass stars and high mass brown dwarfs might harbor an accretion disk. Second, the new data amass additional evidence for accretion in the substellar domain. To the best of our knowledge, there are only three papers dealing with high resolution spectra in Taurus brown dwarfs (White \\& Basri 2003; Muzerolle et al. 2003;  and Jayawardhana  et al. 2003). The first work found three accretors in  a sample of ten very low mass stars and brown dwarfs, although the less massive, a M6.5 (GM~Tau), has a mass just above the substellar limit. The second paper includes four objects  whose spectral type is M6 or M7. None of them seem to undergo active accretion based on the width of H$\\alpha$, although the authors claim that one of them, namely MHO-5, is accreting based on the detection of forbidden lines of oxygen at 6300 and 6363 \\AA{ } and CaII IRT. Moreover, three out of the four were observed at moderate resolution, including MHO-5 (R$\\sim$8000). Regarding the later study, the four M7-M7.75 brown dwarfs discussed there, with masses down to 0.05 M$_\\odot$ (again, using Baraffe et al. 1998 models), do not show accretion either. Muzerolle et al. (2003) also analyzed seven Taurus members whose spectral types are M4.75--M5.75 (just above or at the substellar borderline). Three out of these seven are accreting, based on the width of H$\\alpha$ and other spectral indicators. Therefore, up to date, our study  presents the only  brown dwarf belonging to Taurus (kpno-tau-03), at the substellar limit, which has been proved to be accreting.  In conjunction with the studies quoted in the previous  paragraph, our results indicate that about 10\\% of the Taurus brown dwarfs (one out of 11) is actively accreting. This fraction is much smaller than the accretion occurrence in low mass stars in the association or the estimate for the substellar domain (about 50\\%) based purely on the strength of H$\\alpha$ measured in low resolution spectra (Barrado y Navascu\\'es \\& Mart\\'{\\i}n 2003). The statistical criterion defined in this last work is based on the saturation of the activity, as measured in several young open clusters (namely IC2391, Alpha Per and the Pleiades, with ages ranging from about 50 to 125 Myr). The discrepancy in the fraction of accreting brown dwarfs  might imply that there is an additional source of flux in H$\\alpha$ line  or that the accreting criterion based on the H$\\alpha$ width at 10\\% of the maximum intensity, as defined by Jayawardhana et al. (2003), i.e. 200 km/s, is too restrictive.  The detection of accretion in substellar objects indicates that they undergo a phase similar to Classical TTauri stars (for a review, see Bertout 1989 or Appenzeller \\& Mundt 1989). A handful of other brown dwarfs belonging to other young stellar association, with ages  ranging from 1 to 10 Myr, have been observed with high resolution spectroscopy and the presence of accretion confirmed. In at least one, LS-RCrA~1 (Fern\\'andez \\& Comer\\'on 2001; Barrado y Navascu\\'es, Mohanty \\& Jayawardhana 2004), outflows have been seen, by means of the detection of intense, narrow forbidden lines. In a previous paper we have named them Classical TTauri substellar analogs (CTTSA). Therefore, we can conclude, at least in the case of high mass brown dwarfs (Mas$\\sim$0.072--0.04 M$_\\odot$), that they are formed as low mass stars, by fragmentation and collapse of the original molecular cloud. Of course, additional studies are needed to confirm this preliminary conclusion, in particular high resolution imaging in the near and mid- infrared, as well as in the optical. Narrow band imaging might show whether these objects present outflows similar to those observed in some Classical TTauri stars."
    },
    "0403/astro-ph0403628_arXiv.txt": {
        "abstract": "{This paper presents the VIMOS VLT Deep Survey around the Chandra Deep Field South (CDFS). We have measured 1599 new redshifts with VIMOS on the European Observatory Very Large Telescope - UT3, in an area $21 \\times 21.6$ arcmin$^2$, including 784 redshifts in the Hubble Space Telescope - Advanced Camera for Surveys GOODS area. 30\\% of all objects with $I_{AB}=24$ have been observed independently of magnitude, indicating that the sample is purely magnitude limited. We have reached an unprecedented  completeness level of 88\\% in terms of the ratio of secure measurements vs. observed objects, while 95\\% of all objects have a redshift measurement. A total of 1452 galaxies, 139 stars, 8 QSOs have a redshift identification, 141 of these being unsecure measurements. The redshift distribution down to $I_{AB}=24$ is peaked at a median redshift z=0.73, with a significant high redshift tail extending up to $\\sim4$. Several high density peaks in the distribution of galaxies are identified. In particular, the strong peak at z=0.735 contains more than 130 galaxies in a velocity range $\\pm2000$km/s distributed all across the transverse $\\sim$20 $h^{-1}$ Mpc of the survey. We are releasing all redshifts to the community, along with the cross identification with HST-ACS GOODS sources on the CENCOS database environment http://cencosw.oamp.fr. ",
        "introduction": "Understanding the major steps in the evolution of galaxies still remains a major challenge to modern astrophysics. While the general theoretical framework of the  hierarchical growth of structures in the universe including the build up of galaxies is well in place (e.g. \\cite{peacock04}), at high redshifts this picture remains largely unconstrained by observations. The detailed properties of the main population of galaxies from large samples  representative of the universe at various epochs remain to be established across most of the life of the universe beyond the large local volumes explored by the 2dFGRS (\\cite{colless} and the SDSS (\\cite{SDSS}), and expanding from smaller exploratory surveys (\\cite{lilly95}, \\cite{lefevre95}, \\cite{steidel03}, \\cite{cimatti03}).  The VIMOS VLT Deep Survey (VVDS) is a deep redshift survey aimed at studying the evolution of galaxies, large scale structures and AGNs over more than 90\\% of the current age of the universe. The unique feature of the VVDS is the simple magnitude selection applied to define a complete magnitude limited sample of distant galaxies, with a goal of more than 100000 objects observed in multi-object spectroscopy. The VVDS rests on the observations of 5 different fields to smooth out the effects of cosmic variance when building the statistical properties of the galaxy population (\\cite{lefevre04}).  This paper presents the redshift survey observations of 1599 objects with $I_{AB}\\leq24$ conducted by the VVDS team around the Chandra Deep Field South, including the HST-GOODS area (\\cite{goods}). The observations have been carried out with the VIsible Multi-Object Spectrograph (VIMOS) on the 8.2m Melipal telescope of the European Southern Observatory Very Large Telescope. We are describing the processing steps and redshift measurements, and the associated quality control we have applied to these data. The content of the final catalog is detailed, as well as the cross identification with the Hubble Space Telescope Advanced Camera for Surveys GOODS images, and we present the main entries available from our interactive database. The main properties of the sample are briefly presented, including the redshift distribution of the sample. ",
        "conclusions": "In the framework of the VIMOS VLT Deep Survey (VVDS), we have observed a large sample of galaxies around the Chandra Deep Field South, and are releasing the redshift data to the community. A total of 1599 objects with $I_{AB} \\leq 24$ have a measured redshift. The completeness in redshift measurement for the targeted objects is high, above 88\\%. We find that the redshift distribution has a median of $z=0.73$, with strong high density peaks observed across the field.  The combination of this redshift survey and the HST-ACS GOODS survey enables detailed studies of the evolution of galaxies in the Chandra Deep Field South."
    },
    "0403/astro-ph0403673_arXiv.txt": {
        "abstract": "Abell~399 and Abell~401 are both rich clusters of galaxies, with global temperatures of 7.2~keV (Abell~399) and 8.5~keV (Abell~401) respectively. They lie at a projected separation of $\\sim$3~Mpc, forming a close pair. We have observed the system with the \\xmm satellite. The data of each cluster show significant departures from our idealised picture of relaxed rich clusters. Neither of the two contains a cooling flow, and we find that their central regions are nearly isothermal, with some small-scale inhomogeneities. The image analysis derives $\\beta$-values that are smaller than the canonical value of 0.65, and the surface brightness distribution is not symmetric around the central cD galaxies: there are irregularities in the central $\\sim$200~kpc, and asymmetries on larger scales, in that the intracluster gas in each cluster is more extended towards the other member of the system. Both clusters host extended radio halos and a plethora of tailed radio galaxies.  The halo in Abell~399 appears to be correlated with a sharp edge apparent in the \\xmm images, and a region of harder X-ray emission. There is also evidence for enhanced X-ray flux in the region between the two clusters, where the temperature is higher than our expectations.  Although tidal or compression effects might affect the large scale structure of the two clusters, we show that these cannot account for the distortions seen in the inner regions. We argue that the reasonably relaxed morphology of the clusters, and the absence of major temperature anomalies, argues against models in which the two have already experienced a close encounter. The properties of the intermediate region suggests that they are at an early stage of merging, and are currently interacting mildly, because their separation is still too large for more dramatic effects. The substructure we find in their inner regions seems to point to their individual merging histories. It seems likely that in the Abell~399/401 system, we are witnessing two merger remnants, just before they merge together to form a single rich cluster of galaxies. This picture is consistent with recent numerical simulations of cluster formation. ",
        "introduction": "It is now common wisdom that clusters of galaxies are formed hierarchically, by the merging of smaller mass units, preferentially along large-scale filaments. During such violent events one expects physical processes to take place that have significant impact on the properties of the constituents of clusters (i.e., gas, galaxies), that would define their subsequent evolution, as well as the energy and entropy budget in the largest structures in the Universe. During mergers, shock waves propagate within the intracluster medium (ICM), thus increasing its thermal energy and entropy. The cluster galaxies may be severely stripped, leaving their metal rich interstellar media (ISM) behind in the cluster, contributing in this way to the enrichment of the ICM (e.g., Acreman et al. 2003). Cooling flows may get displaced or destroyed entirely (e.g., G\\'{o}mez et al. 2002) .  \\begin{figure} \\begin{center} \\leavevmode \\epsfxsize 1.05\\hsize \\epsffile{m1m2_all_500-10000.ps} \\caption{A mosaic image of all 4 pointings in the (0.5-10)~keV energy range. Background subtracted images from the two MOS instruments (MOS1 and MOS2) are superimposed. The pixel size is 8~arcsec.  Abell~401 is to the North-East of the image.}\\label{mosaic} \\end{center} \\end{figure}    Numerical simulations (Schindler \\& M\\\"{u}ller 1993, Roettiger, Burns, \\& Loken 1996; Ricker 1998; Takizawa 1999; Ritchie \\& Thomas 2002; Burns et al. 2004) have advanced to such an extent that they can provide us with detailed temperature maps and synthetic X-ray images of clusters under formation, at each stage of their evolution. All this work has demonstrated that cluster mergers result in dramatic substructure, visible in the X-ray images and spectra of clusters of galaxies. Initially, during the first approach of two clusters a compression and/or shock wave is developed in the interface between them. When the two cores collide, the temperature and luminosity of the system reaches its maximum. As the two gravitational potentials settle down to form a single one, two lens-shaped shock waves propagate outwards, along the direction of merging. Soon afterwards, a single remnant is formed.  Unfortunately, we have not been able yet to test observationally in a systematic manner, that cluster mergers proceed as theory predicts. Markevitch et al. (1998) set the framework for such investigations, and revealed the gross spectral characteristics of disrupted clusters. Additionally, they were able to comment on the statistics of mergers, but their data were lacking the necessary quality to probe the details of the observed structures. But the results of some recent observational work have been encouraging (e.g., Sun et al. 2002; Kempner, Sarazin, \\& Ricker 2002), as a hotter bar, for example, has been found in a few cases separating the two colliding units, providing the means to calculate the relative velocity of the merging subunits. A detailed comparison of a sample with the numerical simulations is still lacking, and it is unclear whether the simulations include all the necessary physics and are able to model the problem correctly. There are a few examples of cluster properties that have been discovered by the observations, and had not been predicted by the simulations.  Recently, for example, X-ray observations have uncovered features (cold fonts; e.g., Markevitch et al. 2000), that have been now explained as signatures of merging clusters. Cold fronts had not been predicted by the simulations, mainly due to the lack of the adequate resolution and the mishandling of the cooling functions. Only recently some numerical simulations have managed to reproduce them and associate their presence to the motion of a galaxy or group through the ICM of another cluster, and explain the observations (Bialek, Evrard, \\& Mohr 2002).  Ideally, one would like to construct a sample of clusters that are at different stages of their evolution, and trace the merger event. Unfortunately, finding clusters at different stages of the merging sequence is not trivial. Projection effects or past inaccurate determinations of their physical properties influence severely their classification and taxonomy in one of the stages. Abell~399 and Abell~401 appear at a first sight as a very good example of two clusters at early stages of merging. They are equally rich, and past X-ray missions have found them to be both at temperatures between 7 and 8~keV. They have been the subject of past investigations and have been observed in different wavelengths, from the optical to X-ray.  Their projected separation is $\\sim$36~arcmin~$\\sim$~2.96~Mpc\\footnote{Throughout this paper we use $H_0=71~{\\rm km \\ s^{-1} \\ Mpc^{-1}}$, $\\Omega_{\\rm M}$=0.3, and $\\Omega_{\\rm \\Lambda}$=0.7.}, approximately 1-2 times their virial radii.  The line-of-sight velocity difference of the two is of the order of $\\sim 700 \\; {\\rm km \\ s^{-2}}$ (Girardi et al. 1997). The velocity dispersions of Abell~399 and 401 are $\\sim 1180$ and $\\sim 1110$~${\\rm km \\ s^{-2}}$ (Oegerle \\& Hill 2001) respectively. Dynamical models of the system based on the dynamics of the galaxies around the Abell~399/401 system reached the conclusion that it consists a bound cluster pair (Oegerle \\& Hill 1994), and that the galaxies of both clusters should move in the total gravitational potential of both clusters. Additionally, since their line-of-sight velocity difference is smaller than their velocity dispersions, it is probable that the encounter between the two clusters is taking place essentially on the plane of the sky.  Basic properties of the clusters are given in Table~1.  \\begin{table} \\begin{center} \\caption{Target Information}\\label{target_info} \\begin{tabular}{ccccccccc} \\hline \\hline  Cluster\t\t& z\t\t& kpc/''\t\t& $D_L$ (Mpc) \\\\  \\hline  A399\t\t& 0.0724\t\t& 1.36 \t\t& 322.4 \\\\  A401\t\t& 0.0737\t\t& 1.38\t\t& 328.5 \\\\  \\hline  \\end{tabular} \\end{center} \\end{table}   \\begin{table*} \\caption{Pointing Information}\\label{obs_info} \\begin{center} \\begin{tabular}{ccccccccc}   \\hline \\hline  (I)\t\t& (II)\t\t& (III)\t\t& (IV)\t\t& (V)\t\t& (VI)\t\t& (VII) \\\\  Rev \t\t& Obs\t\t& $\\alpha$~(2000) \t& $\\delta$~(2000) \t& Instr.\t\t\t& $Exp$\t\t& $Exp_{corr}$\t\t&  \\\\ & & deg \t\t& deg \t\t& & ksec \t\t& ksec\t\t&  \\\\ \\hline  0127\t\t& 0112260101\t& 44.4491255188460\t& 13.0518334856848\t& MOS1\t\t\t& 14.166 \t& 11.429 \t& \\\\  & & & & MOS2\t& 14.170\t& 11.599\t& \\\\  & & & & PN\t& 9.188\t& 5.249 \t& \\\\  0127\t\t& 0112260201\t& 44.5827088537386\t\t& 13.3186112665767\t\t& MOS1\t\t\t\t& 18.231\t\t& 18.133 \t\t& \\\\  & & & & MOS2\t& 18.234\t& 18.152\t& \\\\  & & & & PN\t& 12.505 \t& 12.395 \t& \\\\  0395\t\t& 0112260301& 44.7636255225170\t& 13.5472223803563    & MOS1\t\t\t& 12.953\t& 12.948\t& \\\\  & & & & MOS2\t& 12.959 \t& 12.897\t& \\\\  & & & & PN\t& 8.086 \t& 8.063\t& \\\\    0395\t\t& 0112260401\t& 44.9844588584281\t& 13.7638057162178\t& MOS1\t\t\t& 11.913 \t& 11.882\t& \\\\  & & & & MOS2\t& 11.917\t& 11.829\t& \\\\  & & & & PN\t& 7.143\t& 7.143 \t& \\\\  \\hline  \\end{tabular} \\vspace{0.2cm} \\begin{minipage}{16cm} \\small NOTES-- (I): revolution number; (II): observation number; (III): pointing $Right Ascension$ ($\\alpha$) in degrees; (IV) pointing $Declination$ ($\\delta$) in degrees; (V) EPIC Instrument; (VI) Exposure time (live time for the central CCD); (VII) Reduced Exposure time, after the subtraction of the bright background flares (see text for more details).  \\end{minipage} \\end{center} \\end{table*}  \\begin{figure*} \\begin{center} \\setlength{\\unitlength}{1cm} \\begin{picture}(8,7) \\put(-6.5,7.5){\\special{psfile=A399_n.ps angle=-90 hscale=40 vscale=40}} \\put(3,7.5){\\special{psfile=A401_n.ps angle=-90 hscale=40 vscale=40}} \\end{picture} \\end{center} \\caption{\\xmm images of Abell~399 (left panel) and Abell~401 (right panel). The images have been smoothed with Gaussian kernels of $\\sigma$ = 2~pixels=16~arcsec. The four sectors, centred on the central cDs in each cluster, and used in the subsequent analysis, are shown.}\\label{A399pies} \\end{figure*}  Previous studies have reached contradictory conclusions about the past history of the pair. The first attempts to reveal signs of interactions in the region between the two with the {\\it Einstein} observatory did not lead to any strong conclusions (Ulmer \\& Cruddace 1981).  The \\asca satellite found evidence for some enhancement in the X-ray flux above what is expected from just the superposition of the two clusters (Fujita et al. 1996). A slight temperature increase in-between them was also recorded. Both results led Fujita et al. (1996) to suggest that Abell~399/401 is a pre-merging pair. However, this scenario provides no explanation for the lack of cooling flows in the cores of both clusters. X-ray observations (e.g., Peres et al. 1998) have found that their centres do not host cooling flows as the mass accretion rate is zero. Additionally, the model of Fujita et al. does not tie in well with the past \\rosat observational facts that both clusters appear disrupted : the work of Slezak, Durret, \\& Gerbal (1994), for example, found evidence for `substructure' in Abell~401, and the \\rosat HRI detector (Fabian, Peres, \\& White 1997) revealed a linear structure that emanates from the centre of Abell~399 and points towards Abell~401. This feature lead to the suggestion that the two clusters have already encountered each-other, and are now moving apart.   In order to recover the dynamical state of this system, decide on its past history and future evolution, and derive vital information that would help us to test the results of the numerical simulations of merging clusters, we observed Abell~399/401 with \\xmm. The observations are presented in Section~2. Section 3 is devoted to the presentation of the properties of each cluster individually, as found from the \\xmm data and analysis. In Section~4 the X-ray properties of the region between the two is investigated, while in Section~5 the radio properties of the clusters are presented, and their large scale environment is discussed in Section~6. Finally, in the discussion section (Section~6) we present models for the dynamical states of both clusters that can fit well the \\xmm radio and optical results. ",
        "conclusions": "We summarise here the main results of the \\xmm study of the Abell~399/401 binary cluster system and the conclusions about the its dynamical state that are favoured by them. \\begin{itemize} \\item{The \\xmm data confirm the lack of cooling flows in the cores of both clusters.} \\item{The image analysis gives $\\beta$ values that are lower than the expected canonical value of 0.65 found in rich clusters of galaxies. In neither cluster is the gas azimuthally symmetric around the central cD galaxy.} \\item{A 2-dimensional analysis reveals a lop-sided excess ($<$200~kpc) around the central galaxies.} \\item{There is temperature structure in the inner (200-400)~kpc of both clusters, which argues for the presence of complex structures within the cluster cores, possibly due to shock waves and cold fronts.} \\item{Abell~399 shows a sharp edge to the East of the cluster centre, which is associated with harder X-ray emission. The temperature profile to the North of Abell~401 declines steeper than expected, and there is a lack of flux in the same area.} \\item{Both clusters appear to host central radio halos, although the one seen in the NVSS data of Abell~399 requires confirmation.  This halo in Abell~399 appears to follow an edge apparent in the X-ray images, and is associated with harder X-ray emission.} \\item{There is a plethora of tailed radio galaxies in and around both systems.} \\item{We find that the flux from the space between the two clusters is slightly enhanced, above what is expected from the superposition of the two clusters. This region appears also hotter than the region on the opposite side of the cluster core, and the temperature distributions predicted for relaxed clusters by numerical simulations.} \\item{The large scale environment around the binary system appears surprisingly empty. There is no evidence in the galaxy distribution and RASS data for the presence of other clusters nearby.} \\end{itemize}   The `substructure' \\xmm finds in the core cluster regions cannot be explained by the tidal field generated by the presence of the other cluster.  On larger scales, (10-20)~arcmin$\\simeq$(0.8-1.6)~Mpc from the cluster centres, tidal forces become gradually far more important, in shaping the gas distributions. However, we find that the extensions of the gases towards the other member of the system seen in the \\xmm data cannot again be solely due to their mutual interactions. If the gases were extended by gravitational forces, their temperature should have been lower than what we observe. It is apparent that an additional driving force is required.  We also find that the properties of Abell~399/401 cannot be reproduced by scenarios that either involve off-centre cluster collisions, or that speculate that they have been through each other once. An offset collision provides the only single mechanism which might explain all the properties of the system, but that it doesn't seem consistent with the observed large separation and limited disturbance of the clusters at large radii from the cluster centres.  On the other hand our analysis finds evidence for increased flux and temperature in the region in-between the two. It appears that the two have already stated interacting mildly, and that we are witnessing a compression region between them.  The most possible explanation for the properties we find in the core region of the clusters is that they are due to each ones past merger activity. It seems that each cluster is a merger remnant. A similar evolutionary scenario has been recently proposed by Belsole et al. (2003) to explain the unequal system Abell~1750. Indeed, in a hierarchical structure formation scenario, binary cluster systems might be the late stages of cluster formation, just before the final merger takes place and one rich cluster is formed. The numerical simulations of Frenk et al. (1996) show two clusters forming at a close proximity, by the merging of smaller units, before the two finally merge to form a single rich cluster at present day. Additional support for our proposed model comes from the inspection of simulated X-ray images and temperature maps (from the Simulated X-ray Cluster Data Archive, http://sca.ncsa.uiuc.edu/). These simulations show the X-ray flux and temperature stucture of clusters as they are formed from the continuous accretion of smaller structure. An extensive search though this archive indicates that pairs of large clusters are formed, indeed by the merging of filaments from directions that do not necessarily coinside with the direction to the nearest massive cluster. Additionally, the development of shock waves can be witnessed when the two clusters are at close separations. Additionally, they confirm the presence of features like the indentaion and shock wave to the East of Abell~399, as being produced during the infall of a smaller group. Hence, for this pair of galaxies, we favour a model which combines some early interaction, with a minor merger history."
    },
    "0403/astro-ph0403390_arXiv.txt": {
        "abstract": "Deep spectropolarimetric observations, obtained with the Very Large Telescope (VLT), are presented for two powerful radio galaxies, 0850--206 (z=1.3373) and 1303+091 (z=1.4093). These observations cover the rest-frame wavelength range $\\sim$\\,1450\\,--\\,3750\\,\\AA. New radio observations and continuum images of the same sources are also presented. These galaxies are the first two observed from a complete sample of nine radio sources with redshifts in the range 1.3\\,$\\le$\\,z\\,$\\le$\\,1.5 (selected from the equatorial sample of powerful radio sources of Best, R{\\\"o}ttgering \\& Lehnert\\nocite{best99b}), as part of a project aimed to investigate the multi-component nature of the UV continuum in radio galaxies and, in particular, any variations of the continuum properties with the radio source age.  The larger radio source of the two, 0850--206, presents a high continuum fractional polarization, averaging 17 per cent across the observed wavelength range and reaching 24 per cent at rest-frame wavelengths of $\\lesssim$\\,2000\\,\\AA. The smaller radio source, 1303+091, shows a lower continuum polarization, averaging 8 per cent and rising to 11 per cent for rest-frame wavelengths $\\gtrsim$\\,3000\\,\\AA. For both galaxies, the position angle of the electric vector is generally constant with wavelength and within $\\sim$\\,15$^{\\circ}$ of perpendicular to the radio axis. Both their total flux spectra and polarized flux spectra reveal the 2200\\,\\AA \\ dust feature, and comparison with dust scattering models suggests that the composition of the dust in these galaxies is similar to that of Galactic dust.  In 0850--206, scattered quasar radiation dominates the UV continuum emission, with the nebular continuum accounting for no more than $\\sim$\\,22 per cent and no requirement for any additional emission component such as emission from young stars. By contrast, in 1303+091, unpolarized radiation could be a major constituent of the UV continuum emission, with starlight accounting for up to $\\sim$\\,50 per cent and the nebular continuum accounting for $\\sim$\\,11 per cent.  The emission-line properties of the galaxies are also studied from their total intensity spectra. Comparison of the measured emission-line ratios with both shock- and photo-ionization models shows that the nuclear and extended gas in these galaxies is mainly photoionized by the central active nucleus. ",
        "introduction": "The galaxies associated with powerful extragalactic radio sources are uniquely important for understanding the physics of active galactic nuclei (AGN) and for studying the relationship between radio source activity and the properties of the host galaxy and its environment.  Powerful distant radio galaxies are thought to be the progenitors of present day, giant ellipticals (e.g. \\pcite{best98,mclure2000}). However, compared with normal elliptical galaxies, powerful radio galaxies at z$\\gtrsim$0.6 show enhanced optical/UV continuum and line emission, which is generally elongated and aligned along the radio axis (e.g. \\pcite{mccarthy87}).  The emission-line structures usually extend large distances from the nucleus (5\\,$\\sim$\\,100 kpc; e.g. \\pcite{tadhunter86,baum88}); the properties of these extended emission-line regions (EELR) can be greatly influenced by shocks resulting from interactions between the radio source structures and the interstellar/intergalactic medium (ISM/IGM) (e.g. \\pcite{clark98,carmen01,carmen03}). Interestingly, the emission-line properties of a sample of z$\\sim$1 radio galaxies have been found to evolve strongly as the radio source passes through the host galaxy \\cite{best2000b}: large radio sources ($\\gtrsim$150 kpc) show quiescent `rotation-dominated' velocity profiles and ionization states in agreement with AGN-illumination, while smaller radio sources present highly distorted kinematic profiles and overall ionizations consistent with being shock-dominated.  As for the optical/UV continuum excess observed in radio galaxies, several different mechanisms are known to contribute to it (e.g. \\pcite{tadhunter2002}), but the relative extent of their contributions remains uncertain. A popular early interpretation for this UV excess was recent star formation, induced by the passage of the radio jet through~the~ISM of the host galaxy or by merger events linked to the radio source triggering (e.g. \\pcite{rees89}). Jet-induced star formation has been shown to be feasible in numerical simulations (e.g. \\pcite{mellema2002}). Also, stellar absorption features characteristic of OB stars have been observed in\\,the radio\\,galaxy 4C41.17 at z=3.8 (Dey et al.\\,1997)\\nocite{dey97}.  The detection of polarized continuum and polarized broad permitted emission lines in some radio galaxies (e.g. \\pcite{antonu-mill85,di-serego89,cimatti93}) suggested that light emitted anisotropically by a hidden quasar nucleus and scattered towards us by dust or electrons in the ISM of the radio galaxy \\cite{tadhunter88,fabian89} also makes an important contribution to the UV light. This is favoured by orientation-based unified schemes for radio sources \\cite{barthel89}, according to which radio galaxies and quasars are drawn from the same parent population but viewed from different angles to the line of sight, with the AGN in radio galaxies obscured by a surrounding dusty torus.  Spectropolarimetric studies of a small number of radio galaxies at z$\\sim$1 have shown that the polarized emission is spatially extended, with the electric vector oriented perpendicular to the UV continuum emission at all wavelengths, and that, while the permitted MgII\\,2800 emission line is observed to be broad in polarized light, none of the narrow lines shows significant polarization (e.g. \\pcite{cimatti96}, 1997)\\nocite{cimatti97}. These results clearly show that the origin of the polarization is scattered quasar light.  More recent spectropolarimetric studies of a sample of radio galaxies at z$\\sim$2.5 \\cite{vernet2001} have shed light on the nature of the scattering material. These studies show indications of a continuum upturn beyond 2200\\,\\AA \\ in the composite spectrum of the galaxies in the sample, which is interpreted as a possible detection of the 2200\\,\\AA \\ dust feature, indicating that the scattering medium is dust. However, caution must be taken because, due to the high redshift of their sample, those wavelengths are redshifted into a region of strong sky lines.  Together with the young stellar and scattered AGN components, other processes known to contribute to the UV continuum excess in some sources are nebular continuum emitted by the extended emission-line gas \\cite{dickson95} and direct AGN light \\cite{shaw95}.  As with the emission-line gas, some properties of the continuum emission in radio galaxies are observed to vary with radio source size: smaller sources show strings of bright knots aligned along the radio axis, and larger sources present more compact nebulae with fewer bright components \\cite{best96}.  Do the contributions of the different components to the UV excess also evolve over the radio source lifetime?  For instance, a large variation is seen in the strengths, and to a lesser extent colours, of the polarized emission in radio galaxies. Is this linked to the radio source size?  Such questions cannot be answered by existing spectropolarimetric data, because targets were generally selected on the basis of previously detected polarization in imaging observations, an interesting UV morphology or an ultra-steep radio spectrum, making the samples studied biased or incomplete.  To remedy this and address many of the issues raised above, such as the nature of the scatterers at high redshift and the variation with radio source size of the different contributions to the UV excess, we have begun a programme to make deep spectropolarimetric observations of a {\\it complete} sample of nine radio galaxies with redshifts in the range 1.3\\,$\\le$\\,z\\,$\\le$\\,1.5 and a wide range in radio sizes.  In this paper we present the observations of the first two galaxies from the sample.  The paper is organized as follows. In Section 2, the selection criteria of our sample are described. Section 3 contains the details of the observations, data reduction and analysis. The results are presented in Sections 4 and 5. Discussion of the results follows in Section 6. Summary and conclusions are provided in Section 7.  Throughout this paper values of the cosmological parameters of H$_{0}$\\,=\\,65\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{\\rm M}$\\,=\\,0.3 and $\\Omega_{\\Lambda}$\\,=\\,0.7 are assumed. ",
        "conclusions": ""
    },
    "0403/astro-ph0403359_arXiv.txt": {
        "abstract": "Solid observational evidences   indicate  a strong dependence  of  the galaxy formation and evolution on the environment. In order to study in particular the  interaction between  the intracluster medium   and the evolution  of cluster galaxies, we  have  created a large database of clusters of galaxies based on the largest available X-ray and optical surveys: the ROSAT  All Sky Survey (RASS),  and the Sloan  Digital Sky Survey (SDSS). We  analyzed the correlation  between the total optical and  the X-ray cluster luminosity.  The resulting correlation of $L_X$ and $L_{op}$  shows a logarithmic slope  of 0.6, a  value close to the self-similar correlation.  We analysed also the  cluster mass to light ratio, by finding  a significant  dependence of  $M/L$ on the  cluster mass with a logarithmic slope ranging from  0.27 in the  i and r bands to 0.22 in the z band. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403445_arXiv.txt": {
        "abstract": "Deep surveys indicate a bubbly structure of cosmological large scale which should be the result of evolution of primordial density perturbations. Several models have been proposed to explain origin and dynamics of such features but, till now, no exhaustive and fully consistent theory has been found. We discuss a model where cosmological black holes, deriving from primordial perturbations, are the seeds for large-scale-structure voids. We give details of dynamics and accretion of the system voids-cosmological black holes from the epochs $(z\\simeq10^{3})$ till now finding that void of $40h^{-1}Mpc$ of diameter and under-density of $-0.9$ will fits the observations without conflicting with the homogeneity and isotropy of cosmic microwave background radiation. ",
        "introduction": "The existence of voids has been  evident after the discovery by Kirshner et al. of a large void with diameter of $60$ Mpc in B\\\"{o}otes (\\cite{kirsh}). Systematic surveys have shown the existence of many regions with similar characteristics. Computer analysis of galaxy distribution gives evidence that voids occupy about 50\\% of the volume of the universe (e.g. see \\cite{El-Ad:1997af}) or, according to a more recent paper (\\cite{hoyle}), about $40$\\% of the volume of the universe.  Today, there is a large agreement on the issue that voids are not just empty regions of space, but that they are regions with a very low density of luminous matter.  As observed by Peebles (\\cite{peebvoid}), the low dispersion of velocities of galaxies indicates that most of matter must be inside the voids, not only if the density parameter (for the matter component) is $\\Omega_{m}=1$ but also in the case $\\Omega_{m}\\ll 1$. In any case, recent observations (\\cite{boomerang}; \\cite{supernovae1}; \\cite{supernovae2}) indicate that the total value of density parameter is $\\Omega= \\Omega_{m}+ \\Omega_{\\Lambda}=1$ where $\\Omega_{m}\\simeq 0.3$ and $\\Omega_{\\Lambda}\\simeq 0.7$. In this case $\\Omega_{\\Lambda}$ is the contribution due to the whole content of unclustered matter which can be cosmological constant, some kind of scalar field (\\cite{quintessence1}; \\cite{quintessence2}; \\cite{quintessence3}; \\cite{quintessence4}) or, in general, \"dark energy\".  It is worthwhile to stress that the visual inspection of galaxy distribution suggests nothing else but the absence of large amount of luminous matter in wide regions. Furthermore, it is not clear whether the voids are spherical regions approximately empty or under-dense regions with arbitrary shapes. Several definitions of voids have been proposed, but a general agreement on their real nature has not been reached yet (\\cite{sch}).  The Swiss-Cheese cosmological model,  initially proposed  by Einstein and Straus (\\cite{einstraus1}; \\cite{einstraus2}), appears suitable for the description of the cosmological voids. In a recent paper (\\cite{cosimo}), it was proposed an approach for the formation of the cosmological voids in the framework of such model. It was shown that voids are the consequence of the collapse of extremely large wavelength perturbations into low-density black holes and of the comoving expansion of matter surrounding the collapsed perturbations.  As a result, it was claimed that in the center of each void there is a black hole whose mass $M$ compensates the mass which the void would have if it were completely filled with matter having a cosmological density.  In \\cite{cosimo}, the voids are empty regions of the universe which grow comovingly with the cosmological expansion. In that case, the presence of cosmic background radiation was neglected.  In this paper, we analyze the physical mechanism capable of explaining the structure of voids in presence of baryonic matter, cosmic background radiation (CBR) with central black holes acting as seeds. The layout of the paper is the following: in Sect.2, we will present the cosmological black hole (CBH) model in the framework of the Friedmann-Lema\\^{i}tre-Robertson-Walker (FLRW) cosmology. Sect.3 is devoted to the discussion of the effects of interaction between the CBR and CBH. A mechanism for the formation of an under-density regime void is analyzed in Sect.4, while the matching with observations, which allows to determine the initial time of voids formation and the mass function of CBHs, is studied in Sect.5. The discussion of results and conclusions are given in Sect.6. ",
        "conclusions": "In this paper, we have developed a model where cosmological black holes are seeds for large scale structure voids. Such systems come out from the evolution of primordial perturbations and result as stable structures from $(z\\simeq10^{3})$ up to now. They enlarge till diameters of about $40h^{-1}Mpc$ and the under-density of voids is of the order $-0.9$ with respect to the background. The whole structure is a sort of honeycombs where most of galaxies (i.e. luminous matter) are located on the edge of voids while most of dynamical mass is sited in the central black hole. The edge is defined by a natural equilibrium condition on the energy due to the balance of gravitational pull of the central black hole and the cosmic expansion. The cosmic background radiation contributes to the accretion of black hole flowing inside the void but its homogeneity and isotropy is not affected in agreement with data. The picture which emerge agrees with optical and IRAS observation (\\cite{El-Ad:1997af}) giving a $\\sim50\\%$ of the volume filled by voids with the above characteristics. The presence of central black holes seems to confirm the gravitational origin of the voids and stabilizes the system against cosmic expansion preventing its evaporation.  It is interesting to note that the order of magnitude observed for the masses of CBH concides with the one of the Great Attractor (\\cite{Fairall}). It is very tempting for us to identify the Great Attractor as a CBH and to use the model described in this paper for explaining the large scale motions observed for the galaxies surrounding it.  However, the main problem with observations of cosmic velocity fields (\\cite{Faber}) is that the voids are, in general, the contrary to Great Attractor, and large scale structure around the voids do not show velocity fields converging toward the voids, but toward the visible clusters and superclusters around the voids. This apparent shortcoming, in the framework of our model (see also \\cite{cosimo}), can be overcome by the Birkoff theorem which states that the stationary solutions are also static if the spherical symmetry is restored. So, a fraction of galaxies is attracted by clusters and superclusters \"outside\" the void while another fraction has no dynamics since it has been already attracted \"inside\" the void. This fact could be interpreted as an early selection due to a competitive mechanism between CBHs and external matter contained into clusters and superclusters.  However, if the Swiss-Cheese model were always valid such a selection would have never been achieved; instead, in a more realistic situation, the model holds only approximately so then we have to expect galaxies inside and outside the void due to the deviations from sphericity and to the perturbations of the CBH mass.  Furthermore, as observed in \\cite{Davis} and \\cite{peebvoid}, the small relative velocity dispersion shows that, if $\\Omega_{M}=1$, then most of the mass has to be contained into the voids. The same authors conjecture that this must be true even when $\\Omega_{M}$ is smaller than 1 as predicted by several CDM simulations."
    },
    "0403/astro-ph0403029_arXiv.txt": {
        "abstract": "We use the cosmic microwave background angular power spectra to place upper limits on the degree to which global defects may have aided cosmic structure formation. We explore this under the inflationary paradigm, but with the addition of textures resulting from the breaking of a global $O(4)$ symmetry during the early stages of the Universe. As a measure of their contribution, we use the fraction of the temperature power spectrum that is attributed to the defects at a multipole of 10. However, we find a parameter degeneracy enabling a fit to the first-year WMAP data to be made even with a significant defect fraction. This degeneracy involves the baryon fraction and the Hubble constant, plus the normalization and tilt of the primordial power spectrum. Hence, constraints on these cosmological parameters are weakened. Combining the WMAP data with a constraint on the physical baryon fraction from big bang nucleosynthesis calculations and high-redshift deuterium abundance limits the extent of the degeneracy and gives an upper bound on the defect fraction of 0.13 (95\\% confidence). ",
        "introduction": "\\label{sec:intro}  Recent measurements of the cosmic microwave background (CMB), notably the first year WMAP data \\cite{Hinshaw2003, Kogut2003}, have proved highly successful in probing the early stages of cosmic structure formation. The observed CMB anisotropies may be produced by taking adiabatic primordial perturbations of roughly the Harrison-Zel'dovich form and evolving these using well-understood, linear physics. Further, the parameter values that are required for this process to a give a match to the data are consistent with those measured using other astronomical techniques. That the primordial power spectrum predicted by many models of inflation is of the required form has become an important success of the inflationary paradigm.  On the other hand, a less attractive property of the paradigm is that successful inflationary models may involve quite different fields, interactions and levels of physical motivation. Here we address the issue using CMB power spectra to constrain models of hybrid inflation \\cite{Linde1994,Lyth1999} that involve the formation of topological defects as inflation ends \\cite{Copeland1994}. Such models, however, do not fit exactly into the above regime. With the existence of topological defects, the seeding of cosmic structure continues after inflation ends, for the defects further perturb the cosmic fluid as long as they continue to be present. For a detailed review of structure formation with defects see Ref. \\cite{Durrer2002}, but generally, a defect-dominated temperature power spectrum does not have pronounced acoustic peaks \\cite{Albrecht1996}. Hence, if defects are added to a passive-evolution case and the normalization reduced to maintain the fit to data on large scales, then the acoustic peaks are slightly suppressed. That is, however, assuming that the other cosmological parameters are not also changed.  Earlier work along these lines \\cite{Battye2000, Bouchet2001, Contaldi1999, Jeannerot1997, Pogosian2003} has tended to be of quite a different vein to that described here. In particular we employ a full likelihood analysis for the fit to data in which the cosmological parameters are free to vary. This freedom has important consequences for the independent use of the temperature power spectrum to constrain such hybrid models --- at least given the presently available data. Previous studies have also tended to focus on local defects in the form of cosmic strings, a class of models that we shall not look at here in any detail, and of such work only \\cite{Pogosian2003} has involved data from the WMAP project.  In our work, the cosmological perturbations are the uncorrelated sum of those from (i) an inflationary adiabatic model, including both scalar and tensor perturbations, and (ii) a global defect model, with an $O(4)$ symmetry breaking at (or after) the end of inflation producing textures (see e.g. \\cite{Durrer2002}). In parametrizing the primordial perturbations, we take there to be negligible running of the scalar spectral index and that the tensor index obeys the single-field consistency condition (see Sec. \\ref{sec:MCMC}). We also assume negligible neutrino masses, as would result from a hierarchical mass variation with neutrino flavor.  The contribution to the CMB power spectra from inflation is found using a variant of the CMBFAST \\cite{Seljak1996} approach, in the form of CAMB \\cite{Lewis2000}. The defect contribution is found by applying the unequal-time correlator (UETC) approach \\cite{Pen1997} to numerical field evolution simulations, as will be discussed in the next section. The UETC method addresses the problem that, in order to calculate CMB power spectra to sub-degree scales, a simulation conventionally requires a dynamic range that is far in excess of that which is feasible with current technology. While it is possible to make full-sky CMB maps \\cite{Pen1994,Landriau2004}, this direct approach is currently limited to relatively low multipoles: $\\ell<20$, although these potentially contain more information than the power spectra alone. The UETC method boosts the dynamic range via a series of theoretical simplifications, for example causality, such that high resolution power spectra calculations may be performed from the simulation data. To do so also involves calculations of the CMBFAST-type, which we carried out using a version of CMBEASY \\cite{Doran2003} modified so as to deal with the defect scenario.  Despite the benefits of the UETC approach, the computational requirements of CMB calculations that include defects far exceed those that do not. As a result, the calculation of CMB spectra for a vast number of different cosmological parameter values is not attainable. Hence, the popular Markov chain Monte Carlo (MCMC) approach \\cite{Verde2003, Lewis2002, Gelman1992}, which involves many thousand such calculations, cannot be fully applied to the defect case. However, the non-defect case fits the WMAP data well and defect-dominated structure formation does not give the required acoustic peaks, suggesting that the defect contribution is small.  If this is the case, then the result of a small change in the cosmological parameters used for the defect calculation is a second order effect. Therefore, the defect contribution needs only to be calculated once, using currently favored values of the cosmological parameters (see Sec. \\ref{sec:CMBcalc}). The defect contribution is then fixed, except for a normalization factor, which is free since it is not known at which energy scale the defects formed. Hence the approach used here, which is described more fully in Sec. \\ref{sec:MCMC}, is to apply the standard MCMC procedure to the primordial contribution and add in the defect component with its normalization varied as an MCMC parameter. This has been achieved using a slightly modified version of CosmoMC \\cite{Lewis2002}, which is directly linked to CAMB. As this extra parameter controls the degree to which the CMB power spectra differ from the usual non-defect spectra, it shall be the main focus of this paper.  However, we shall not present our results in terms of this parameter directly. Rather we shall use the fractional defect contribution to the temperature power spectrum at a particular multipole, $\\ell=10$. The correspondence between the two is roughly linear for low fractions but with a slight spread due to the variation of the non-defect contribution to the chosen multipole. The fractional quantity is, however, more directly understandable.  The data that we have used here are principally that from the first year WMAP release \\cite{Hinshaw2003, Kogut2003}: the temperature power spectrum and the temperature-polarization (TE) cross-correlation spectrum. Other CMB projects, such as ACBAR \\cite{Kuo2004}, CBI \\cite{Pearson2003, Readhead2004} and VSA \\cite{Dickinson2004}, which give data out to higher multipoles than WMAP do not provide much in the way of additional constraints on our model. Applying data on cosmological parameters from, for example, work on big bang nucleosynthesis (BBN) \\cite{Kirkman2003} and measurements of the Hubble parameter by the Hubble Key Project (HKP) \\cite{Freedman2001} has proved more important, as will be detailed in section \\ref{sec:MCMC}.  The changes to the matter power spectrum that the inclusion of defects causes may be found in an entirely analogous manner to the CMB calculation. However, while there have been recent steps forward in measurements of this from galaxy redshift surveys, such as 2DFGRS \\cite{Percival2001} or SDSS \\cite{Tegmark:2003ud}, and from the Lyman-$\\alpha$ forest \\cite{Croft2002}, we have chosen not to use such data. Galaxy formation in the presence of defect-induced density perturbations is not understood, and even in pure inflation scenarios, inferences from the Lyman-$\\alpha$ forest must be drawn with care \\cite{Seljak:2003jg}. Further, we do not believe that the use of such data would significantly change our results, as we shall discuss in Sec. \\ref{sec:MCMC}. This is a conservative position, driven by our desire to make reliable and statistically meaningful statements about the relative importance of global defects. ",
        "conclusions": "\\label{sec:con}  We have found that in order to constrain the extent to which global defects assisted the seeding of cosmic structure, extra data in addition to the CMB temperature power spectrum must be used, or we require a greater knowledge of the spectrum than we have at present. This is despite the uncertainties in the WMAP data set used being dominated by cosmic-variance over the range where the defect contribution to the power spectrum is greatest. The freedom in the cosmological parameters and the primordial power spectrum are sufficient to allow a significant defect contribution while still fitting the data. It is only when parameters such as the physical baryon density and the Hubble parameter are restricted, via alternative astronomical techniques, that our model is significantly constrained. Then we find that the temperature power spectrum at a multipole of $\\ell=10$ may have a fraction of 0.13 attributed to defects (95\\% confidence). The uncertainty in this upper bound that comes from the MCMC approach is 3\\%. More importantly, the numerical errors in the defect calculation are believed to be of order 10\\%, in which case there may be of order 10\\% change in this result. Also, if the primordial tensor component was removed then this result would increase by about 15\\%. In addition, this bound is quite sensitive to BBN result for the value for the physical baryon fraction $\\Obhh$ used and there has yet to be full agreement in this value among authors.  A further result of the degeneracy found is that, if defects were to contribute to cosmic structure formation, then there would a change in the values of the cosmological parameters estimated from the current CMB data. Most notably this affects $\\Obhh$, $h$ and $\\ns$ all of which are subject to an increase upon the addition of defects. This acts to re-enforce the more general point that any inferences made about the cosmological parameters from the WMAP data are model dependent and should be treated with caution.  We note that only a single defect type and model has been used in this investigation: textures resulting from the breaking of an $O(4)$ symmetry. However, the contributions from other global defect models are broadly the same as that considered here, although their spectra are by no means identical. Considering the $O(N)$ class of models for $N=2$ (strings), 3 (monopoles), 4 (textures) and 5 (non-topological textures) there is a gradual variation with $N$ of the relative contributions at low multipoles ($\\ell\\sim10$) and high multipoles ($\\ell\\sim300$), with strings giving preference to the latter \\cite{Pen1997}. This is likely to reduce the fractional contribution allowed from strings by perhaps 30\\% at $\\ell=10$. Local defects in the form of cosmic strings may give a broad peak at $\\ell\\sim400$ \\cite{Battye2000, Contaldi1999}, beyond the first acoustic peak, and hence the results may be quite different.  Further, we wish to point out that the WMAP project is ongoing and new data with reduced uncertainties, as well as the addition of the EE polarization power spectrum, will be released. While the defect contribution to the EE spectrum is likely to be significant at only at large scales, the freedom for defects when using the CMB power spectra alone will be lessened by these data. It may be, however, that the Planck satellite \\cite{planck} is required to really explore sub-dominant defect contributions using these power spectra independently of other data. This mission will give precise temperature power spectrum measurement at sub-WMAP resolutions. It will also detect the B-mode polarization, which would be produced by the vector and tensor components of the defect perturbations. Finally, the non-Gaussianity of defect-induced perturbations has not been satisfactorily addressed, either in the matter power spectrum or the CMB. This may lead to more sensitive tests of defect scenarios."
    },
    "0403/astro-ph0403265_arXiv.txt": {
        "abstract": "We present an analysis of high resolution optical spectra for a sample of very young, mid- to late M, low-mass stellar and substellar objects: 11 in the Upper Scorpius association, and 2 (GG Tau Ba and Bb) in the Taurus star-forming region.  Effective temperatures and surface gravities are derived from a multi-feature spectral analysis using TiO, \\na and \\pot, through comparison with the latest synthetic spectra. We show that these spectral diagnostics complement each other, removing degeneracies with temperature and gravity in the behavior of each.  In combination, they allow us to determine temperature to within 50K and gravity to within 0.25 dex, in very cool young objects.  Our high-resolution spectral analysis does not require extinction estimates.  Moreover, it yields temperatures and gravities {\\it independent} of theoretical evolutionary models (though our estimates do depend on the synthetic spectral modeling).   We find that our gravities for most of the sample agree remarkably well with the isochrone predictions for the likely cluster ages.  However, discrepancies appear in our coolest targets: these appear to have significantly lower gravity (by upto 0.75 dex) than our hotter objects, even though our entire sample covers a relatively narrow range in effective temperature ($\\sim$ 300K).  This drop in gravity is also implied by inter-comparisons of the data alone, without recourse to synthetic spectra.  We consider, and argue against, dust opacity, cool stellar spots or metallicity differences leading to the observed spectral effects; a real decline in gravity is strongly indicated.  Such gravity variations are contrary to the predictions of the evolutionary tracks, causing improbably low ages to be inferred from the tracks for our coolest targets.  Through a simple consideration of contraction timescales, we quantify the age errors introduced into the tracks through the particular choice of intial conditions, and demonstrate that they can be significant for low-mass objects that are only a few Myr old.  However, we also find that these errors appear insufficient to explain the magnitude of the age offsets in our lowest gravity targets.  We venture that our results may arise from evolutionary model uncertainties related to accretion, deuterium-burning and/or convection effects.  Finally, when combined with photometry and distance information, our technique for deriving surface gravities and effective temperatures provides a way of obtaining masses and radii for susbtellar objects independent of evolutionary models; radius and mass determinations are presented in Paper II. ",
        "introduction": "In the study of stars, there is perhaps no parameter more fundamental than stellar mass, which is pivotal in determining the entire evolutionary path traced by a star.  With the discovery of brown dwarfs, mass determination has become particularly important at the bottom of the Main Sequence.  After all, the very notion of brown dwarfs is predicated on mass:  these are substellar objects, which is to say they are less massive than the hydrogen-burning limit of $\\sim$ 80 \\mj. The derivation of ultra-low masses has become particularly crucial in the light of new claims of planetary mass objects (which we henceforth contract to ``planemos'') occuring in isolation in star-forming regions (e.g., \\cite{Zapa00, Lucas00}).  While both planemos and brown dwarfs are substellar, the distinction between the two is drawn at the fusion boundary of $\\sim$ 13 \\mj: brown dwarfs undergo an initial phase of deuterium fusion, while planemos never harbor any fusion at all.  The existence of isolated planemos with a few Jupiter masses, if proven, has significant consequences for both star and planet formation (\\cite{Padoan02, Bate02, Boss01}).  However, the key issue of whether the newly discovered free-floating objects really have planetary masses, or are simply misidentified brown dwarfs, remains unsettled for reasons we touch on below.  Only a precise mass derivation can unequivocally resolve this question.  The empirical determination of substellar masses, though, has so far proven rather difficult.  The most direct approach is to apply Kepler's laws to binary (or higher-order) systems with known orbital parameters, in order to obtain a `dynamical mass'.  If the components can be directly observed, then the mass can be related to their other properties, such as temperature and luminosity, and the dynamical mass results extended to other directly detected bodies for which a dynamical mass cannot be acquired (e.g., free-floating objects).  However, suitable systems with brown dwarf components remain elusive (though a handful are now coming to light; e.g, \\cite{Close02, Potter02, Lane01}).  Dynamical masses for planemos in circumstellar planetary systems (i.e., extrasolar planets) {\\it have} been derived, but these objects have not been directly observed, so it is impossible to relate their masses to other physical properties\\footnote{With the exception of the transiting planets HD 209458 and OGLE-TR056 (\\cite{Charbonneau00} and \\cite{Konacki03}, respectively), which have both mass and radius determinations.  Also, except in transiting cases, masses from RV surveys are only lower limits, due to the unknown inclination of the system.}.  To date, with the exception of HD 209458, no substellar object - whether brown dwarf or planemo - whose intrinsic spectrum has been observed has also been proven to be substellar by a direct, dynamical measurement of its mass.  Conversely, no object outside the Solar System with a dynamical substellar mass has actually been directly detected (again with the sole exception of HD 209458)\\footnote{Some of the atmospheric features of the transiting planet HD 209458 have now been directly detected (\\cite{Charbonneau02}; \\cite{Vidal03}).  A comparison with the theoretical tracks reveals serious discrepancies between observed and predicted properties, probably due to stellar irradiation effects.  A total mass has now also been obtained for one binary system with a spectrum, GJ 569B, indicating at least one brown dwarf component; individual masses should be available in the near future; \\cite{Lane01}.}.  Currently, the substellarity of all directly observed bodies is certified either through the lithium test, or by the fact that their effective temperatures are below the minimum Main Sequence temperature (\\cite{Basri00}).  However, both tests can only indicate an upper limit to the mass, and not its precise value.  Moreover, they are not useful at very early ages, when lithium is undepleted even in low mass stars and substellar objects have not yet cooled sufficiently.  For these reasons, precise mass estimates for all directly observed substellar bodies have so far been obtained solely through comparison with theoretical evolutionary tracks.  However, the evolutionary models remain largely untested in an absolute sense, and different groups have generated somewhat different tracks (e.g., \\cite{Baraffe98, Burrows97, Dantona94}).  Evolutionary modeling uncertainties arise from a variety of sources, such as the treatment of convection, choice of initial conditions and the modifying effects of any initial accretion phase.  Additionally, these uncertainties are exacerbated at the earliest ages and for very low masses (\\cite{Baraffe02}).  It is precisely for these ages and masses, however, that the evolutionary models are currently most widely used: to distinguish between stellar and substellar objects in young clusters (since the lithium and temperature tests are not robust at these ages), as well as to infer the actual substellar masses.  For example, the planemo status of newly identified isolated bodies in young star-forming regions, mentioned earlier, depends entirely on the accuracy of the theoretical tracks.  Therefore, given the burgeoning uncertainty in the tracks with decreasing mass and age, it would be extremely useful to have an independent method for determining mass, in order to check the theoretical predictions.  We present such a method here, which relies on an accurate measurement of surface gravity from high-resolution spectra.  In a nutshell, we first derive surface gravities and effective temperatures (\\teff), by comparing our high-resolution optical spectra to the latest synthetic spectra.  We then calculate radii and masses by combining our inferred gravity and \\teff with photometry and distance measurements.  In this paper, we present the gravity and temperature spectral analysis.  The radius and mass calculations are presented in a companion paper (\\citet{Mohanty03}; hereafter, Paper II).  For cool, very low mass objects, the derivation of surface gravities from observed spectra is not trivial.  A vast multitude of molecular and atomic spectral lines arise in their low-temperature photospheres; an adequate modeling of these is a prerequite for inferring gravity (and \\teff).  This has only become possible in the last several years, due to large advances in computing power and the compilation of increasingly comprehensive linelists.  The resulting synthetic spectra are now routinely used to infer the \\teff and model the colors of low-mass stellar and substellar objects (e.g., \\cite{Allard01, Schweitzer01, Leggett01, Leggett00, Burrows00, Basrimoh00, Kirk99}).  Here we take the next logical step, of using the synthetic spectra to derive surface gravities as well.  \\teff and gravity, together with empirical distance and photometry information, then allow us to calculate masses and radii without resorting to evolutionary model predictions.  Clearly, the validity of our analysis depends on that of the synthetic spectra; we will assess the veracity of these spectra at appropriate junctures.  It must be pointed out, however, that synthetic spectra are essential even when masses are derived through comparisons with theoretical evolutionary models - either for translating observed spectral types to temperatures (when comparing the data to theoretical HR diagrams), or for converting theoretical temperatures and luminosities to expected photometry (when comparing the data to theoretical color-magnitude diagrams).  In other words, both our technique, as well as analyses using evolutionary models, are limited by any shortcomings in the synthetic spectra.  In this paper we will derive, through spectral analysis, \\teff and gravities for a sample of young, late-type (i.e., low mass) Pre-Main Sequence objects belonging to the Upper Scorpius and Taurus regions.  Since we observe them prior to their main contraction phase, we expect their surface gravities to be lower than in field objects of similar mass.  By comparing our results to the predictions of evolutionary models, we will also test these models in the low mass, young age regime where they are most uncertain. ",
        "conclusions": "The primary conclusion of this paper is that it is quite feasible to derive fairly precise (0.25 dex) gravities in low mass stellar and substellar objects. This can be accomplished with spectra of moderate S/N and high resolution (which generally require 8-m class telescopes), along with the most advanced model atmosphere and spectral calculations. If one performs this for objects whose radii can be found (this generally requires reasonably precise distances and photometry), then masses can be found well enough to distinguish planemos, brown dwarfs, and stars (see Paper II). Along with the gravity, we obtain a good effective temperature ($\\pm$50K), which facilitates the conversion of a luminosity to a radius.  The basic reason for our success is that TiO is a spectral diagnostic that is strongly dependent on temperature, and weakly dependent on gravity (and with the opposite gravity dependence of our atomic line diagnostics). Our gravity diagnostics are both subordinate and resonance neutral alkali lines (of potassium and sodium). They have both a temperature and gravity dependence, and yield degenerate solutions in the two parameters. We therefore rely on TiO to break the temperature degeneracy, and the alkali lines then provide a gravity. We note that there are a number of other alkali lines that we have not yet employed (other lines of potassium and sodium, plus lines of rubidium, cesium, and lithium), which cover a range of wavelengths and have different strengths for different stellar temperatures. In the L dwarfs, a substitute for the TiO temperature diagnostic must be found (since TiO has condensed into dust). The metal hydride molecules (FeH and CrH) are the obvious candidates. Properly developed, our methodology should work well over the range of substellar objects from mid-M through L (which encompass a range of masses from the planetary to stellar domains, depending on age). Indeed, once theoretical tracks have been properly computed and calibrated, one might be able to directly infer age as well (in principle).  Having applied our methodology to very young, low-mass objects in the Upper Scorpius and Taurus star-forming regions, we find {\\it (1)} a good agreement between our gravities and the theoretical ones for most of our sample, but {\\it (2)} a significant discrepancy between the two for our coolest targets.  We show that even without detailed comparison against model spectra, the data themselves suggest a sharp fall-off in gravity in the coolest objects.  We have examined various processes and synthetic spectral uncertainties which might lead us to infer erroneous gravities (e.g., dust, cool spots, metallicity variations, and inadequacies in the model treatment of collisional broadening), and found that they are unlikely to give rise to the spectral effects we see: real gravity variations are indicated. Theoretical tracks would interpret this range of gravities as a range of ages, and we provide several arguments against that interpretation. The alternative is that the theoretical tracks are allowing the very low-mass objects to contract too quickly; the gravities we measure for them are lower than predicted. We discuss various ways in which this could happen, and suggest remaining uncertainties in the model treatment of accretion, deuterium-burning and/or convection are most likely responsible.  This means that the interpretation of masses and ages from isochrone analysis in young clusters is problematic at the moment for very low-mass bodies. Nonetheless, we show in Paper II that the interpretation of the faintest of these objects as low-mass brown dwarfs or planetary mass objects looks correct.  Finally, we point out that our spectral analysis is carried out in the optical, but it could potentially be accomplished at other wavelengths as well.  In particular, the near-IR would be very interesting to explore.  Such investigations have already begun: \\cite{Doppmann03a} and \\cite{Doppmann03b} present a similar high-resolution study in the near-IR for higher mass, hotter PMS objects, while \\cite{Gorlova03} examine near-IR PMS substellar spectra, albeit at low-resolution.  A high-resolution analysis of cool, very low-mass PMS objects  would be extremely fruitful, especially for probing the youngest, most extincted bodies; we have recently embarked on a project to accomplish this."
    },
    "0403/gr-qc0403014_arXiv.txt": {
        "abstract": "We study the anisotropies of the Galactic confusion noise background and its effects on LISA data analysis.  LISA has two  data streams of the gravitational waves signals relevant for low frequency regime. Due to the anisotropies of the background, the matrix for their confusion noises has off-diagonal components and depends strongly on the orientation of the detector plane. We find that the sky-averaged confusion noise level $\\sqrt {S(f)}$ could change by a factor of 2 in three months, and would be minimum when the orbital position of LISA is either around the spring or autumn equinox. ",
        "introduction": "The Laser Interferometer Space Antenna (LISA) is planned to  be launched around 2011 and is expected to establish  a new window in the low frequency gravitational wave astronomy from $0.1{\\rm mHz}$ to $ 100$mHz \\cite{lisa,Cutler:2002me}. Its main astrophysical targets are  Galactic binaries and  cosmological massive black holes (MBHs) that merge with other MBHs or capture compact objects.  As  for the Galactic binaries, there would be  a lot of sources in the LISA band. For example we will  be able to resolve several thousand  close white dwarf binaries at $f\\gsim 3$mHz \\cite{bcn1,bcn2}. At lower frequency regime $f\\lsim  3$mHz they are highly overlapped in the frequency bins, and it would be difficult to  resolve them individually. As a result, they form a confusion noise background whose magnitude could be larger than the detector noise at $0.1{\\rm mHz}\\lsim f \\lsim 3$mHz \\cite{lisa}. The coalescence frequency  $f_c$ of a MBH system is given by its redshifted total mass $M_z$ as $f_c\\sim 2(M_z/2\\times 10^6M_\\odot)^{-1}$mHz. The mass function of MBHs is highly uncertain at the lower end $\\sim 10^5M_\\odot$ \\cite{Ferrarese:2000se}, and we might have to  search treasurable signals  from cosmological  MBHs in the Galactic confusion noise background.   The spatial distribution of the Galactic binaries would trace the Galactic structure well,  and the confusion noise background would be strongly anisotropic. The background itself can be regarded as a signal that would provide some information on our Galaxy \\cite{gia,Cornish:2001hg,Ungarelli:2001xu}. But we should notice that there exists a more straightforward approach to probe the Galactic structure using  thousands of resolved binaries whose three dimensional positions are estimated in some error boxes \\cite{gia}.  In this paper we will focus on the role of the background as a noise, and study the effects of its anisotropies on the signal  analysis of LISA. We do not pay attention to the normalization of the background that has been studied by other papers using an isotropic approximation of the source distribution \\cite{bcn1}.   This paper is organized as follows. In Section 2 we study the response of interferometers to the anisotropic gravitational wave background with using the long wave approximation, and formulate the noise matrix for the Galactic confusion background. In Section 3 we numerically evaluate the bakground noise as a function of the orientation of the detector,  and estimate the annual modulation of the noise level  for LISA. In Section 4 we discuss the anisotropies of the signal to noise ratio (SNR) of sources at a fixed distance. Section 5 is devoted to a brief summary of this paper. In the Appendix the confusion noise matrix is analyzed in a  different manner from Section 2. ",
        "conclusions": "In this paper we  have studied the anisotropies of the Galactic confusion noise background and its effect on the data  analysis of LISA at the low frequency regime  $f\\lsim 3$mHz. In contrast to the traditional monopole approximation the anisotropies induce the correlation of confusion noises between the $A$ and $E$ modes of LISA, and the effective noise level depends strongly on the time. In the followings we briefly summarize our results.   The effective noise level of LISA becomes smallest around the autumn equinox, and largest in summer for the configuration in figure 2. The difference of the noise $\\sqrt{S^B_{eff}}$ at these two epochs is a factor of 2.  The detector plane of LISA at the autumn equinox is almost normal to  the Galactic poles and  the effective noise level is very close to  the global minimum  of the noise map given for the all possible orientations of the detector. The    time average of the effective noise is different from the traditional  monopole approximation by less than $50\\%$.   We have also analyzed the dependence of the SNR on the sky position of sources with a fixed  distance and averaged orientation. For a long lived source the dependence   after 1yr integration is weak, and fluctuations  on the sky  are   $\\sim 30\\%$ level. This is mainly due to the angular average of the response function by annual rotation of LISA. For a short lived source the SNR at a given direction changes strongly with time. For example the SNR of a source  at the direction of the Galactic pole changes by a factor of $\\sim4 $ in three months. A factor of  2 comes from the modulation of the confusion noise level and another factor of  2 is from that of  the angular response function of the detector."
    },
    "0403/astro-ph0403579_arXiv.txt": {
        "abstract": " ",
        "introduction": "In this contribution, I review the basic ingredients of the so-called Randall-Sundrum brane cosmology. This new cosmological scenario is based on the  assumption that our universe is a {\\it brane}: a sub-space embedded in a  bulk spacetime, with a single extra  dimension. In contrast with other braneworld models, the self-gravity of the brane is  taken into account.  As will be recalled in this contribution, the cosmology for such a brane-universe is modified in two respects: \\begin{itemize} \\item  the Friedmann equation is modified at high energy; \\item the bulk influences the cosmological evolution via an addititional term, usually called  dark radiation or Weyl radiation. \\end{itemize} Whereas the first modification has an impact only during the very early universe, since it is significant only at high energy, the second effect could have observable consequences today, as discussed below.  In this contribution, I do not discuss other important topics in the context of brane cosmology, such as the issue of cosmological perturbations, since this topic will be discussed by Roy Maartens. Finally, for the reader who wants to learn more on this subject, he/she can find in the literature several detailed reviews \\cite{reviews} that cover much more  than the present contribution. ",
        "conclusions": "The various models of extra dimensions with branes have raised a considerable interest in the last few years, motivated by their more or less direct connections with the recent developments in string/M theory. To make further progress in this direction, it is important to see how these models can be tested by experiments.  Roughly speaking, the tests can be classified into three broad categories: modification of Newton's law; signatures in colliders; cosmology. As usual in high energy physics, if the scale characterizing new physics is too high then it cannot be  reached directly in collider experiments. In this case cosmology is the only place where the effects of  new physics can be, indirectly, observed.  As illustrated by numerous works in the last few years, Randall-Sundrum type cosmology is a very rich playground to study the very peculiar consequences of the braneworld idea in cosmology. As recalled  in this contribution, its essential new features are:  a $\\rho^2$ term in the generalized Friedmann equation, which dominates at high energies; ``dark radiation'',  {\\it produced} during the high energy phase, and of potential relevance for observations, via the nucleosynthesis constraints on the number of extra relativistic degrees of freedom."
    },
    "0403/astro-ph0403053_arXiv.txt": {
        "abstract": "We present the results of qualitative consideration of possible \\q changes occurring during the transition from the hot accretion disc to the cool one. We argue the possible existence of one more type of spiral density waves in the inner part of the disc where gasdynamical perturbations are negligible. The mechanism of formation of such a wave as well as its parameters are considered.  We also present the results of 3D gasdynamical simulation of cool accretion discs. These results confirm the hypothesis of possible formation of the spiral wave of a new, ``precessional\" type in the inner regions of the disc. Possible observational manifestations of this wave are discussed. ",
        "introduction": "The analysis of principal processes of matter heating and cooling in accretion discs presented in the work [\\ref{di12}] has shown that for realistic parameters of accretion discs in semidetached binaries ($\\mdot\\simeq10^{-12}\\div10^{-7}M_\\odot/\\mbox{year}$ and $\\alpha\\simeq10^{-1}\\div10^{-2}$)\\footnote{$\\mdot$ -- mass transfer rate; $\\alpha$ -- dimensionless parameter introduced by Shakura and Sunyaev [\\ref{shakura}, \\ref{shak2}] for expression of viscosity coefficient $\\nu=\\alpha c_s H$ ($H$ -- disc semithickness, $c_s$ -- sound speed).} the gas temperature in outer part of the disc lies in the range from $\\sim 10^4$~K to $\\sim 10^6$~K. Earlier we have conducted 3D gasdynamical simulation of accretion discs both for a `hot' case (the gas temperature in the outer part of the disc was 200--500 thousands~K [\\ref{di1}--\\ref{di11}]), and for the `cool' case (the temperature of gas in the outer part of the disc didn't exceed $\\sim 2 \\times 10^4$~K [\\ref{di12}]). The analysis of these results has shown that for both cases (i.e. independently on the disc temperature) the self-consistent solution didn't involve the shock interaction between the stream of matter from the inner Lagrangian point $L_1$ and formed accretion disc (``hot spot\"). Energy release zone (``hot line\") is located outside the disc and is formed as the result of the interaction of the circumdisc halo and circumbinary envelope with the stream.  The ``hot line\" model is found to be in a good agreement with observations [\\ref{Tanya98}--\\ref{Tanya2003b}].   For `hot' solutions we have investigated both general morphology of gaseous flows in semidetached binaries and the structure of formed hot accretion discs (see, e.g., [\\ref{di11},\\ref{ippeg}]). In particular, we have found only one arm of spiral shock wave generated by the tidal influence of the mass-losing star. The two-armed spiral shocks were discovered by Matsuda {\\it et al.} in [\\ref{Sawada86}--\\ref{spiral2}]. Nevertheless, our 3D gasdynamical simulations for the `hot' case have shown the presence of only one-armed spiral shock while in the place where the second arm should be the stream from $L_1$ dominates and presumably prevents the formation of the second arm of tidally induced spiral shock. Besides we have found out that for the `hot' case the variation of mass transfer rate leads to the disc perturbations and to the formation of spiral-vortex structure in the disc [\\ref{blob1}--\\ref{fridman}].   Even a glance at the morphology of gaseous flows for the `cool' case (see, e.g., [\\ref{di12},\\ref{rujichka}]) discovers that the accretion disc for this case is characterized by principally different parameters as compared to the `hot' case. In particular, the disc has more circular form and the second arm of the tidal spiral wave is present. Our analysis [\\ref{di12}] shows that opposite to the `hot' case tidal spiral waves don't propagate to the inner part of the cool accretion disc and are located in the outer part of the disc only.  The aim of this work is the investigation of the structure of cool accretion discs in semidetached binaries. Section~2 contains the qualitative analysis of changes occurring in transition from a hot accretion disc to the cool one. In particular, here we suggest the possible existence of one more type of spiral density waves in the inner part of the disc where gasdynamical perturbations are negligible. We also consider the mechanism of formation of such a wave and its parameters. Section~3 contains the results of 3D gasdynamical simulation of flow structure for the case when radiative cooling is effective and the gas temperature drops down to $\\sim 10^4$~K in the whole region. These results confirm the hypothesis of possible formation of a spiral wave of a new type in the inner regions of the disc.  Possible observational manifestations of the discovered spiral wave of the new, ``precessional\" type as well as main conclusions~are~drawn~in~Section~4.  \\renewcommand{\\thefigure}{1} \\begin{figure}[t] \\centerline{\\hbox{\\psfig{figure=ris1.eps,width=\\textwidth}}} \\caption{\\small The sketch of main peculiarities of the morphology of gaseous flows in semidetached binaries for the case of high gas temperature.} \\end{figure}  \\renewcommand{\\thefigure}{2} \\begin{figure}[t] \\centerline{\\hbox{\\psfig{figure=ris2.eps,width=\\textwidth}}} \\caption{\\small The sketch of main peculiarities of the morphology of gaseous flows in semidetached binaries for the case of low gas temperature.} \\end{figure} ",
        "conclusions": "The qualitative analysis of possible changes occurring in transition from a hot accretion disc to the cool one shows the possible generation of spiral wave of a new type in the accretion disc. The appearance of this spiral wave in the inner part of the disc where gasdynamical perturbations are negligible is due to retrograde precession of flowlines in the binary system.  The analysis of presented results of 3D gasdynamical simulation fully confirms our hypothesis on the possible generation of spiral wave in the inner part of the cool disc. The correspondence between the qualitative analysis and the computational results permits us to argue the precessional mechanism of the wave formation. Increasing of the radial flux of matter after passing the density wave results in growth of accretion rate and formation of a compact zone of energy release on the accretor surface. This zone can be seen as a periodic increase of brightness on light curves of semidetached binaries. Observation of this zone will permit to determine the precession rate of the wave so these observations will both give the proof of the existence of ``precessional\" wave in the inner part of the cool accretion disc and provide the information on characteristics of the inner parts of the disc."
    },
    "0403/astro-ph0403609_arXiv.txt": {
        "abstract": "Infrared spectroscopy provides a direct handle on the composition and structure of interstellar dust.  We have studied the dust along the line-of-sight towards the Galactic Center using Short Wavelength Spectrometer (SWS) data obtained with the Infrared Space Observatory (ISO). We focussed on the wavelength region from 8--13 $\\mu$m which is dominated by the strong silicate absorption feature. Using the absorption profiles observed towards Galactic Center Sources (GCS) 3 and 4, which are C-rich Wolf-Rayet Stars, { as reference objects,} we are able to disentangle the interstellar silicate absorption and the silicate emission intrinsic to the source, toward Sgr A$^*$ and derive a very accurate profile for the intrinsic 9.7 $\\mu$m band. The interstellar absorption band is smooth and featureless and is well reproduced using a mixture of 15.1\\% amorphous pyroxene and 84.9\\% of amorphous olivine { by mass}, all in spherical sub-micron-sized grains. There is no direct evidence for substructure due to interstellar crystalline silicates. By { minimizing} $\\chi^2$ { of spectral fits to the absorption feature} we are able to determine an upper limit to { the degree of crystallinity of silicates} in the diffuse interstellar medium (ISM), and conclude that the { crystalline fraction of the interstellar silicates} is 0.2\\% $\\pm$ 0.2\\% by mass.  { This} is much lower than the degree of crystallinity observed in silicates { in the circumstellar environment of} evolved stars, { the main contributors of dust to the ISM}. { There are two possible explanations for this discrepancy.} First, an amorphization process  occurs in the { ISM} on a time scale significantly shorter than the destruction time scale, possibly caused by { particle bombardment by} heavyweight ions. Second, we consider the possibility that the crystalline silicates in stellar ejecta are diluted by an additional source of amorphous silicates, in particular supernovae.  { We also compare our results with a study on silicate pre-solar grains found in interplanetary dust particles.} ",
        "introduction": "In the last decade, a multitude of evidence for the presence of crystalline silicates in various astrophysical environments has emerged. In particular, infrared spectra have revealed that silicates in circumstellar environments often contain a significant crystalline fraction, both around post-main-sequence stars \\citep[e.g.][]{WMJ_96_mineralogy,MWT_02_xsilI} and pre-main-sequence stars \\citep[e.g.][]{WWD_96_xsilsyoung,MWB_01_haebe}. In addition, crystalline silicates are ubiquitous in the Solar System, not only in the more evolved bodies such as planets, but also in primitive objects like comets \\citep[e.g.][]{W_02_cometgrains}. Because crystallization is inhibited by high energy barriers, the origin and evolution of the crystalline silicate fraction in interstellar and circumstellar media has the potential to provide direct evidence of the energetic processing of grains.  The life cycle of dust starts in the outflow of evolved stars, continues upon ejection in the interstellar medium and eventually ends in the planet forming disk around a young star.  It is surprising that, whereas at the beginning and end of dust grains' lifes crystallinity is prevalent, no crystallinity is found in the intermediate phase (i.e.~in the diffuse interstellar medium). In fact, a relatively high upper limit to the degree of crystallinity in the diffuse ISM has been determined recently \\citep{LD_01_silicate}. Less than 5\\% by number of the interstellar Si-atoms were found to be incorporated in crystalline silicate grains of $<1$ $\\mu$m in size, which is roughly equivalent to a mass fraction of $< 5$\\%. On the other hand, \\citet{BA_02_mineralogy} have suggested that, while the broad and structureless interstellar 10 $\\mu$m absorption feature is commonly ascribed to amorphous silicates, the crystalline spectral detail may be washed out in a very complex mixture of crystalline silicates. Furthermore, some studies on silicates in the dense ISM have reported the (controversial) detection of crystalline silicate features \\citep{CJL_00_xsil_in_orion,OO_03_carbononions}, but one has to bear in mind that this environment has very different physical properties than the diffuse ISM.  In this work, we will re-address the issue of crystallinity in the diffuse ISM by studying the line-of-sight towards the Galactic Center. Because of its large amount of extinction and its high infrared flux, the sightline towards the Galactic Center has often been used to characterize the properties of interstellar dust \\citep{RA_85_extinction,RRP_89_gc,PSA_94_GC,TWA_96_GC,LFG_96_gc,CPG_98_hydrocarbon}. We will study the 10 $\\mu$m silicate absorption feature in order to determine the degree of crystallinity in this line-of-sight. In Sect.~\\ref{sec:obs} we will discuss the ISO SWS observations and data reduction, as well as the characteristics of the region around Sgr A$^{*}$ and the correction method applied for emission intrinsic to the Galactic Center (GC) region. The method used to determine the dust composition in the diffuse ISM is described in Sect.~\\ref{sec:determination} along with the results. In Sect.~\\ref{sec:processing} we discuss two mechanisms to explain the discrepancy in crystallinity observed between stellar ejecta and the diffuse ISM. A comparison with silicates in the solar system and planet forming disks around other stars is given in Sect.~\\ref{sec:comparison}. Sect.~\\ref{sec:conclusions} contains the conclusions. ",
        "conclusions": "\\label{sec:conclusions}  In this work, the results of a study on the crystallinity of silicates in the diffuse ISM have been presented. It is known that crystalline silicates are common in the circumstellar environment of both pre- and post-main-sequence stars. It is generally accepted that the dust particles found in the circumstellar environment of young stars originate from evolved stars, and have arrived in their current location after a long (several Gyr) residence time in the interstellar medium. The puzzling lack of evidence of crystalline silicates in the ISM, prompted us to set a firm upper limit on the crystallinity.  Studying the 10 $\\mu$m silicate feature in absorption towards Sgr A$^*$ we are able to determine that at most 0.4\\% of the silicates in the interstellar medium have a crystalline structure. The data are best fitted with a degree of crystallinity of 0.2\\%. In addition, we have determined the composition of the amorphous silicate component in the diffuse interstellar medium. We found that 84.9\\% of the amorphous grains are olivine (Mg$_{2x}$Fe$_{2(1-x)}$SiO$_4$) and 15.1\\% are pyroxene (Mg$_{x}$Fe$_{(1-x)}$SiO$_3$). The amorphous grains were found to be spherical and smaller than $\\sim 0.1$ $\\mu$m in radius. { Detailed analysis of the 10 $\\mu$m feature indicates that the pyroxenes are probably slightly Mg-rich (with $0.5 < x < 0.6$) and the interstellar olivines may be Fe-rich (with $0.4 < x < 0.5$). }   By comparison of these results with the crystallinity of the silicates produced by mass-losing stars, and the interstellar grain destruction rate, we have determined the interstellar amorphization rate, and found that crystalline silicates are effectively amorphitized in 5 Myr to achieve a final crystallinity of 0.2\\%, while an amorphization time scale of 9 Myr is consistent with the determined upper limit of 0.4\\%. These numbers are only estimates, as the crystallinity of stellar ejecta is not determined very accurately yet.  Amorphization by low energy { ion bombardment} has been explored as an explanation for the low degree of crystallinity of silicates in the ISM, but the results { are not yet} consistent with the observed low degree of crystallinity.  In addition, dilution by other sources of amorphous silicates, such as supernovae, are not sufficient to explain the observed lack of crystalline silicates in the ISM.  We conclude that effectively all silicates in the diffuse interstellar medium become amorphous on a very short time scale, compared to the total residence time in the diffuse ISM. Hence, the crystalline silicates found in the circumstellar environments of young stars, in the solar system and in star formation regions have not survived the ISM unaltered, but are probably crystallized locally. Both annealing as well as evaporation and subsequent condensation seem to be significant crystallization processes."
    },
    "0403/astro-ph0403323_arXiv.txt": {
        "abstract": "We present here a scenario, based on a low reheating temperature $ T_R \\ll 100$~MeV at the end of (the last episode of) inflation, in which the coupling of sterile neutrinos to active neutrinos can be as large as experimental bounds permit (thus making this neutrino ``visible'' in future experiments). In previous models this coupling was forced to be very small to prevent a cosmological overabundance of sterile neutrinos. Here the abundance depends on how low the reheating temperature is. For example, the sterile neutrino required by the LSND result does not have any cosmological problem within our scenario. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/nucl-th0403084_arXiv.txt": {
        "abstract": "Cross sections for the photon-induced particle-emission reactions ($\\gamma$,n), ($\\gamma$,p), and ($\\gamma$,$\\alpha$) are given for all natural isotopes from Ti to Bi. The target nuclei are assumed to be in their ground states, except for $^{180}$Ta which is naturally occurring as the isomer $^{180\\mathrm{m}}$Ta. The cross sections are calculated in a statistical model (Hauser-Feshbach) approach and covering an energy range from threshold up to 7.35 MeV above the threshold (14.7 MeV above threshold for ($\\gamma$,$\\alpha$) reactions). The results are intended to aid conception and analysis of experiments which can also be used to test the methods involved in predicting astrophysical reaction rates for nucleosynthesis. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/nucl-th0403059_arXiv.txt": {
        "abstract": "The possibility of appearance of spin polarized states in symmetric and strongly asymmetric nuclear matter is analyzed within the framework of a Fermi liquid theory with the Skyrme effective interaction. The zero temperature dependence of the neutron and proton spin polarization parameters as functions of density is found for SkM$^*$, SGII (symmetric case) and SLy4, SLy5 (strongly asymmetric case) effective forces. By comparing free energy densities, it is shown that in symmetric nuclear matter ferromagnetic spin state (parallel orientation of neutron and proton spins) is more preferable than antiferromagnetic one (antiparallel orientation of spins). Strongly asymmetric nuclear matter undergoes at some critical density a phase transition to the state with the oppositely directed spins of neutrons and protons while the state with the same direction of spins does not appear. In comparison with neutron matter, even small admixture of protons strongly decreases the threshold density of spin instability. It is clarified that protons become totally polarized within a very narrow density domain while  the density profile of the neutron spin polarization parameter is characterized by the appearance of long tails near the transition density. ",
        "introduction": "The spontaneous appearance of  spin polarized states in nuclear matter is the topic of a great current interest due to relevance in astrophysics. In particular, the effects of spin correlations in the medium strongly influence the neutrino cross section and neutrino luminosity. Hence, depending on whether nuclear matter is spin polarized or not, drastically different scenarios of supernova explosion and cooling of neutron stars can be realized. Another aspect relates to pulsars, which are considered to be rapidly rotating neutron stars, surrounded by strong magnetic field. There is still no general consensus regarding the mechanism to generate such strong magnetic field of a neutron star. One of the hypotheses is that magnetic field can be produced by a spontaneous ordering of spins in the dense stellar core.  The possibility of a phase transition of normal nuclear matter to the ferromagnetic state was studied by many authors. In  the gas model of hard spheres,  neutron matter becomes ferromagnetic at $\\varrho\\approx0.41\\,\\mbox{fm}^{-3}$~\\cite{R}. It was found in Refs.~\\cite{S,O} that the inclusion of long--range attraction significantly increases the ferromagnetic transition density (e.g., up to $\\varrho\\approx2.3\\,\\mbox{fm}^{-3}$ in the Brueckner theory with a simple central potential and hard core only for singlet spin states~\\cite{O}). By determining magnetic susceptibility with Skyrme effective forces, it was shown in Ref.~\\cite{VNB} that the ferromagnetic transition occurs at $\\varrho\\approx0.18$--$0.26\\,\\mbox{fm}^{-3}$. The Fermi liquid criterion for the ferromagnetic instability in neutron matter with the Skyrme interaction is reached at $\\varrho\\approx2$--$4\\varrho_0$~\\cite{RPLP}, where $\\varrho_0=0.16\\,\\mbox{fm}^{-3}$ is the nuclear matter saturation density.    The general conditions on the parameters of neutron--neutron interaction, which  result in a magnetically ordered state of neutron matter, were formulated in Ref.~\\cite{ALP}. Spin correlations in dense neutron matter were studied within the relativistic Dirac--Hartree--Fock approach with the effective nucleon--meson Lagrangian in Ref.~\\cite{MNQN}, predicting the ferromagnetic transition at several times nuclear matter saturation density. The importance of the Fock exchange term in the relativistic mean--field approach for the occurrence of ferromagnetism in nuclear matter was established in Ref.~\\cite{TT}. The stability of strongly asymmetric nuclear matter with respect to spin fluctuations was investigated in Ref.~\\cite{KW}, where it was shown that  the system with localized protons can develop a spontaneous polarization, if  the neutron--proton spin interaction exceeds some threshold value. This conclusion was confirmed also by calculations within the relativistic Dirac--Hartree--Fock approach to strongly asymmetric nuclear matter~\\cite{BMNQ}.  For the models with realistic nucleon--nucleon (NN) interaction, the ferromagnetic phase transition seems to be suppressed up to densities well above $\\varrho_0$~\\cite{PGS}--\\cite{H}. In particular, no evidence of ferromagnetic instability has been found in recent studies of neutron matter~\\cite{VPR} and asymmetric nuclear matter~\\cite{VB} within the Brueckner--Hartree--Fock approximation with realistic Nijmegen II, Reid93 and Nijmegen NSC97e NN interactions. The same conclusion was obtained in Ref.~\\cite{FSS}, where magnetic susceptibility of neutron matter was calculated with the use of the Argonne $v_{18}$ two--body potential and Urbana IX three--body potential.  Here we  continue the study of spin polarizability of nuclear matter with the use of an effective NN interaction. As a framework of consideration, a Fermi liquid (FL) description of nuclear matter is chosen~\\cite{AKP,AIP}. As a potential of NN interaction, we use the Skyrme effective interaction, utilized earlier in a number of contexts for nuclear matter calculations~\\cite{SYK}--\\cite{AAI}.  We explore the possibility of FM and AFM phase transitions in nuclear matter, when the spins of protons and neutrons are aligned in the same direction or in the opposite direction, respectively. In contrast to the approach, based on the calculation of magnetic susceptibility, we obtain the self--consistent equations for the FM and  AFM spin order parameters and find their solutions at zero temperature. This allows us to determine not only the critical density of instability with respect to spin fluctuations, but also to establish the density dependence of the order parameters and to clarify the question of thermodynamic stability of various phases. The main emphasis in our study will be laid on the region of zero isospin asymmetry (symmetric nuclear matter) and large isospin asymmetry  (strongly asymmetric nuclear matter and neutron matter).  Note that we consider  the thermodynamic properties of spin polarized states in nuclear matter up to the high density region relevant for astrophysics. Nevertheless, we take into account the nucleon degrees of freedom only,  although other degrees of freedom, such as pions, hyperons, kaons,  or quarks could be important at such high densities. ",
        "conclusions": "Spin instability is a common feature, associated with a large class of Skyrme models, but is not realized in more microscopic calculations. The Skyrme interaction has been successful in describing nuclei and their excited states. In addition, various authors have exploited its applicability to describe bulk matter at densities of relevance to neutron stars \\cite{SMK}. The force parameters are determined empirically by calculating the ground state in the Hartree--Fock approximation and by fitting the observed ground state properties of nuclei and nuclear matter at the saturation density. In particular, the interaction parameters, describing spin--spin and spin--isospin correlations, are constrained from the data on isoscalar \\cite{T,LS} and isovector (giant Gamow--Teller)~\\cite{HHR,SGE,BDE} spin--flip resonances.  In a microscopic approach, one starts with the bare interaction and obtains an effective particle--hole interaction by solving iteratively the Bethe--Goldstone equation. In contrast to the Skyrme models, calculations with realistic NN potentials predict more repulsive total energy per particle for a polarized state comparing to the  nonpolarized one for all relevant densities, and, hence, give no  indication of a phase transition to a spin ordered state. It must be emphasized that the interaction in the spin-- and isospin--dependent channels is a crucial ingredient in calculating spin properties of isospin symmetric and asymmetric nuclear matter and different behavior at high densities of the interaction amplitudes, describing spin--spin and spin--isospin correlations, lays behind this divergence in calculations with the effective and realistic potentials.  In this study as a potential of NN interaction we chose SkM$^*$ and SGII (symmetric nuclear matter) as well as  SLy4 and SLy5 (strongly asymmetric nuclear matter) Skyrme effective forces. The models SkM$^*$ and SGII~\\cite{SG} have been constrained by fitting the properties of nucleon systems with very small isospin asymmetries, while the models  SLy4 and SLy5 were further constrained to reproduce the results of microscopic neutron matter calculations (pressure versus density curve)~\\cite{CBH}. Besides, in a recent publication~\\cite{SMK} it was shown that the density dependence of the nuclear symmetry energy, calculated up to densities $\\varrho\\lesssim3\\varrho_0$ with SLy4 and SLy5 effective forces (together with some other sets of parameters among the total 87 Skyrme force parametrizations checked) gives the neutron star models in a broad agreement with the observables, such as the minimum rotation period, gravitational mass--radius relation, the binding energy, released in supernova collapse, etc. This is an important check for using these parametrizations in the high density region of strongly asymmetric nuclear matter. However, it is necessary to note, that the spin--dependent part of the Skyrme interaction at densities of relevance to neutron stars still remains  to be constrained. Probably, these constraints will be obtained from the data on the time decay of magnetic field of isolated neutron stars~\\cite{PP}. In spite of this shortcoming, SLy4 and SLy5 effective forces hold one of the most competing Skyrme parametrizations at present time for description of isospin asymmetric nuclear matter at high density (together with SkM$^*$ and SGII forces for description of symmetric nuclear matter) while a Fermi liquid approach with Skyrme effective forces provides a consistent and transparent framework for studying spin instabilities in a nucleon system.  In summary, we have considered the possibility of phase transitions into spin ordered states of symmetric and strongly asymmetric nuclear matter within the Fermi liquid formalism, where NN interaction is described by the Skyrme effective forces (SkM$^*$, SGII and SLy4, SLy5 potentials for the regions of vanishing and strong isospin asymmetry, respectively). In contrast to the previous considerations, where the possibility of formation of FM spin polarized states was studied on the base of calculation of magnetic susceptibility,  we obtain the self--consistent equations for the FM and AFM spin order parameters and solve them in the case of zero temperature. It has been found that nuclear matter demonstrates different behavior at high densities with respect to spin fluctuations in isospin symmetric and strongly isospin asymmetric cases. In the model with SkM$^*$ and SGII effective forces symmetric nuclear matter undergoes a FM phase transition, when the spins of protons and neutrons are aligned along the same direction. In the model with SLy4 and SLy5 effective forces strongly asymmetric nuclear matter is subjected to a phase transition into the spin polarized state with the oppositely directed spins of neutrons and protons, while the state with the same direction of the neutron and proton spins does not appear. In the last case, an important peculiarity of the corresponding phase transition is the existence of long tails in the density profile of the neutron spin polarization parameter near the transition point. This means that even small admixture of protons to neutron matter leads to a considerable shift of the critical density of spin instability in the direction of low densities. In the model with SLy4 effective interaction this displacement is from the critical density $\\varrho\\approx3.7\\varrho_0$ for neutron matter to $\\varrho\\approx2.4\\varrho_0$ for asymmetric nuclear matter at the isospin asymmetry $\\alpha=0.95$, i.e. for $2.5\\%$ of protons only. As a result,  the state with the oppositely directed spins of neutrons and protons appears, where protons become totally polarized in a very narrow density domain.  As a consequence of this study,  important questions appear, what is the value of the threshold asymmetry, at which the parallel spin ordering at small isospin asymmetry is changed to the antiparallel spin ordering at large isospin asymmetry, and do the obtained results survive for another type of an effective interaction, e.g., for Gogny effective force~\\cite{BGG,FVS} or monopole  effective interaction~\\cite{RSM}? This research is in progress and will be reported elsewhere."
    },
    "0403/astro-ph0403101_arXiv.txt": {
        "abstract": "Although atmospheric transmission spectroscopy of HD209458b with the Hubble Space Telescope has been very successful, attempts to detect its atmospheric absorption features using ground-based telescopes have so far been fruitless. Here we present a new method for probing the atmospheres of transiting exoplanets which may be more suitable for ground-based observations, making use of the Rossiter effect. During a transit, an exoplanet sequentially blocks off light from the approaching and receding parts of the rotating star, causing an artificial radial velocity wobble. The amplitude of this signal is directly proportional to the effective size of the transiting object, and the wavelength dependence of this effect can reveal atmospheric absorption features, in a similar way as with transmission spectroscopy. The advantage of this method over conventional atmospheric transmission spectroscopy is that it does not rely on accurate photometric comparisons of observations on and off transit, but instead depends on the relative velocity shifts of individual stellar absorption lines within the same on-transit spectra. We used an archival VLT/UVES data set to apply this method to HD209458. The amplitude of the Rossiter effect is shown to be $1.7^{+1.1}_{-1.2}$ m/sec higher in the Sodium D lines than in the weighted average of all other absorption lines in the observed wavelength range, corresponding to an increment of 4.3$\\pm$3\\% (1.4$\\sigma$). The uncertainty in this measurement compares to a photometric accuracy of 5$\\times 10^{-4}$ for conventional atmospheric transmission spectroscopy, more than an order of magnitude higher than previous attempts using ground-based telescopes. Observations specifically designed for this method could increase the accuracy further by a factor 2$-$3. ",
        "introduction": "\\begin{table*} \\caption{ \\label{log} The log of the observations, with in column 1 the date, in column 2 the total the exposure time, in column 3 the typical signal to noise ratio per exposure, in columns 4 and 5 the range in airmass and seeing, and in column 6 the orbital phase of HD209458b.} \\begin{tabular}{ccccccrl} Data&Observing      & Exposure        & S/N & Airmass & Seeing  & Phase   & Comments \\\\ Set&Date           & Times (sec)     &     &         & ($''$)  &         &          \\\\ 1&05/08/2002& 15$\\times$400   & 400 & 1.407$-$1.558&  $\\sim$2& $-$0.007,$+$0.021&on transit\\\\ 2&11/08/2002& 15$\\times$400   & 535 & 2.151$-$1.394&0.9$-$1.5& $+$0.663,$+$0.693&off transit\\\\ 3&12/08/2002& 15$\\times$400   & 515 & 1.391$-$2.085&0.6$-$0.9& $-$0.012,$+$0.017&on transit\\\\ 4&15/08/2002& 15$\\times$400   & 475 & 1.388$-$1.704&0.7$-$1.7& $+$0.826,$+$0.857&off transit\\\\ 5&13/09/2002& 15$\\times$400   & 485 & 1.958$-$1.385&0.8$-$1.8& $+$0.003,$+$0.026&on transit\\\\ 6&20/09/2002& 15$\\times$400   & 515 & 1.695$-$1.384&0.7$-$1.0& $-$0.009,$+$0.020&on transit\\\\ 7&22/09/2002& 15$\\times$400   & 590 & 1.446$-$1.493&0.5$-$1.0& $+$0.569,$+$0.599&off transit\\\\ \\end{tabular} \\end{table*}   Since the discovery that the extra-solar planet HD209458b transits its host star (Charbonneau et al. 2000; Henry et al. 2000; Mazeh et al. 2000), many attempts have been made to detect its atmosphere using transmission spectroscopy. During a transit, light that passes from the star through the outer parts of the planet's atmosphere has impressed on it a spectrographic signature of the atmospheric constituents, which can be observed as an extra absorption on top of the stellar spectrum (Seager \\& Sasselov 2000; Brown 2001; Hubbard et al. 2001). Observations with the Hubble Space Telescope (HST) using this technique have been a great success. First the Sodium D feature was discovered by Charbonneau et al. (2002) with a relative depth of 0.02\\%, followed by the detection of very strong absorption features at a 5$-$10\\% level in the ultra-violet from Hydrogen, Oxygen and Carbon (Vidal-Madjar et al. 2003; 2004). The latter most likely come from the `exosphere' of HD209458b, an extended cometary tail of gas evaporating from the planet caused by its migration close to the star (eg. Schneider et al. 1998).  In strong contrast, attempts to detect the atmosphere of HD209458b from the ground have so far been unsuccessful. So far, only upper limits of typically 1-2\\% have been reached, for the Sodium D feature in the optical (Brown et al. 2000; Bundy \\& Marcy 2000; Moutou et al. 2001), and  for He, CO, H$_2$O and CH$_4$ in the near-infrared (Brown et al. 2002; Harrington et al. 2002; Moutou et al. 2003).  Ground-based observations suffer greatly from the fact that transmission spectroscopy relies on the comparison of spectra on and off transit, which are necessarily taken on different nights. This limits the accuracy due to varying weather conditions and instability of the instruments. In this paper a new method for probing the atmospheres of transiting exoplanets is presented, making use of the Rossiter effect. Due to the rotation of the host star, a transiting planet will first block off light from the approaching and then from the receding parts of the stellar surface. This results in an artificial wobble in radial velocity, an effect first observed by Rossiter (1924) for the eclipsing binary $\\beta$ Lyrae. The amplitude of this signal is directly proportional to the effective size of the transiting planet, and the wavelength dependence of this effect can reveal atmospheric absorption features, in the same way as with transmission spectroscopy. The advantage of this method is that it does not rely on accurate photometric comparison of spectra in and out of transit, but instead depends on relative velocity shifts of stellar absorption lines in the same in-transit spectra. ",
        "conclusions": "The analysis of the UVES archive data shows that the Rossiter effect provides a powerful tool to probe the atmospheres of transiting exoplanets. Although no significant signal has been detected for Sodium D in HD209458b, it is the first time that ground-based observations result in an accuracy of $<$0.001. Previous ground-based observations have reached accuracies of only 1$-$2\\%, corresponding to an excess in the Rossiter effect of 20$-$50 m/sec, clearly inferior to the precision achieved here.  The HST Transmission spectrum of HD209458b obtained by Charbonneau et al. (2002) revealed an excess in absorption in Sodium D of 0.00023 compared to the surrounding continuum, implying a relative increase of $\\sim$1.8\\%. This indicates that for this particular absorption feature an excess in the Rossiter effect of $\\sim$0.7 m/sec can be expected, making its detection very challenging for the current instrumentation. However, it is important to realise that the excess in absorption as determined with the HST was measured over a bandwidth of 12\\AA. In contrast, line-to-line variations in the Rossiter effect will be most sensitive to planetary absorption on scales of the half-power width of the stellar lines (a fraction of an \\AA ngstrom). Most models of transmission spectra (eg. Brown 2001) indicate that on this scale the absorption signal for Sodium D could be 2$-$5 times higher (1$-$3 m/sec).   The archival data sets used in this paper are not optimal for measuring the Rossiter effect. First of all, since a significant fraction of the time was devoted to calibrators, typically only 60\\% of the on-transit time was spent on target, and on several occasions the periods that the amplitude of the Rossiter effect is expected to be the strongest were missed. In addition, the current analysis is significantly hampered by the fact that no or hardly any time was spent on target directly before or after each transit. This was often dictated by the timing of the transit and visibility of HD209458 from Paranal. For the most optimal transits, HD209458 can be observed from 1 hour before to 1 hour after the overall transit with the VLT. In those cases the zero-point in velocity can be determined independently from the on-transit data, and one could also correct for any possible residual drifts in radial velocity over the transit, eg. due to residual telluric absorption line features or inaccuracies in flat fielding. Overall, this should result in an increase in precision of a factor $2-3$. It means that, if the current value is believed, a 3$-$5 $\\sigma$ detection of Sodium D should be possible over a couple of transits.  \\subsection*{Wavelength variations in limb-darkening}  The exact shape and amplitude of the Rossiter effect are dependent on the stellar limb darkening. So far we have not taken into account that limb darkening decreases as function of wavelength. Furthermore, the limb darkening effect also changes differently over the profile of each absorption line. Here we assess the possible influence of both effects. The wavelength dependence of the limb darkening of HD209458 has been observed by Deeg et al. (2001), with $\\epsilon$ changing from 0.62 at 5900 \\AA$ $ to 0.58 at 6180 \\AA. To estimate the influence of this effect we assumed a linear wavelength dependence of $\\epsilon(\\lambda) = 1.46-\\lambda/7000\\AA$, and calculated the change in radial velocity for each line relative to that of Sodium using the model described above. By incorporating this wavelength dependent limb-darkening, the excess of the Rossiter effect in the Sodium D lines is increased by $\\sim$0.2 m/sec.  The variations in limb darkening across solar absorption lines is studied in great detail by Pierce \\& Slaughter (1982) for Sodium D, and by Balthasar (1988) for a whole range of bright absorption lines, by measuring changes in the absorption line depths across the solar disk. The strength of this effect is dependent on the mechanism of formation of the line. Most absorption lines weaken towards the limb of the Sun, meaning that the limb darkening in the centre of the lines is less than in the surrounding continuum. Assuming that this effect is similar for HD209458 as for the Sun, our modeling indicates that these variation contribute at a level of again $\\sim$0.2 cm/sec. Although this is insignificant for these data, it means that future observation aimed at measuring variations in the Rossiter effect at much higher precision should take these effects into account.  \\subsection*{Future applications for this method}  We believe that the new method presented here to probe atmospheres of transiting exoplanets using the Rossiter effect, has great potential. Specifically targeted observations, rather than the archival data used here, can increase the accuracy further by a factor 2$-$3. Furthermore, in the near-infrared, where the atmospheric absorption features due to H$_2$O, CO, and CH$_4$ cover many stellar lines, a great improvement over current observations can be reached. Although confusion with water-vapor and methane in the earth atmosphere will complicate the analysis. In addition, anticipating the discovery of exoplanets transiting bright stars of either later spectral type (with many more stellar lines to average), or with faster stellar rotations (resulting in a Rossiter effect with a higher amplitude), a further significant increase in accuracy can be expected."
    },
    "0403/astro-ph0403337_arXiv.txt": {
        "abstract": "The structure and evolution of low-mass W UMa type contact binaries are discussed by employing Eggleton's stellar evolution code \\citep{egg71,egg72,egg73}. Assuming that these systems completely satisfy Roche geometry, for contact binaries with every kind of mass ratios (0.02$\\sim$1.0), we calculate the relative radii ($R_{1,2}/A$, where $R_{1,2}$ are the radii of both stars, and $A$ the orbital separation) of both components of contact binaries in different contact depth between inner and outer Roche lobes. We obtain a radius grid of contact binaries, and can ensure the surfaces of two components lying on an equipotential surface by interpolation using this radius grid when we follow the evolution of the contact binaries. Serious uncertainties concern mainly the transfer of energy in these systems, i.e., it is unclear that how and where the energy is transferred. We assume that the energy transfer takes place in the different regions of the common envelope to investigate the effects of the region of energy transfer on the structure and evolution of contact binaries. We find that the region of energy transfer has significant influence on the structure and evolution of contact binaries, and conclude that the energy transfer may occur in the outermost layers of the common convective envelope for W-type systems, and this transfer takes place in the deeper layers of the common envelope for A-type systems. Meanwhile, if we assume that the energy transfer takes place in the outermost layers for our model with low total mass, and find that our model steadily evolves towards a system with a smaller mass ratio and a deeper envelope, suggesting that some A-type W UMa systems with low total mass could be considered as the later evolutionary stages of W-subtype systems, and that the surface temperature of the secondary excesses that of the primary during the time when the primary expands rapidly, or the secondary contracts rapidly, suggesting that W-subtype systems may be caused by expansion of the primary, or by the contraction of the secondary. ",
        "introduction": "It is probable that more than 50 per cent of all stars are in binary or multiple systems. An unknown, but possibly large, percentage of these systems are sufficiently close that sometime during their lifetime interacting as a result of Roche lobe overflow (RLOF). The W UMa-type contact binaries are the most common ones, comprising some 95 per cent of eclipsing variables in the solar neighborhood \\citep{shap48} or one stars in every 1000--2000 in the same spectral range \\citep{egg67}. Allowing for selection effects, W UMa stars may even contribute 1 per cent of all F and G dwarfs \\citep{vant75a}. A more recent discussion of contact binaries in the solar neighborhood is carried out by \\citet{ruc02}. He considers the complete sample of 32 EW, EB, and ellipsoidal (ELL) variables with $V<7.5$, and gives the frequency as 1 per 500 stars with $-0.5<V<5.5$, including a serious estimate of the uncertainty.  W UMa stars are found both in young and in old galactic clusters \\citep{vant75b}. Among these clusters are: 1. NGC 2602, NGC 6383, NGC 7235 ($\\la 10^{7}$ yr); 2. Pleiades, Coma, Praesepe ($\\la 5\\times 10^{8}$yr); 3. NGC188, M67 ($\\ga 5\\times 10^{9}$ yr). For this reason it seems we should be able to construct zero-age as well as evolved contact models. Further, there would be a natural explanation for the existence of W UMa stars in young as well as in old cluster if zero-age contact models could be shown to evolve on a nuclear timescale. Structure and evolution of contact binaries are complex and by no means well understood. Although theoretical investigation on early-type systems is almost absent, our discussion will be restricted to late-type systems. The observational properties \\citep{mon81,ruc93} that provide the best clues to the nature of late-type W UMa systems are that the systems are of fairly low total mass, and are in shallow contact, that no equal mass systems exist, that the mass-luminosity relation is unusual, and that most of the W-type systems are to be un unevolved or slightly evolved only, and some of them are known to be unevolved.  \\citet{luc68} had the key idea of a common convective envelope in which the entropy is constant and energy is transferred from the primary to the secondary in the common convective envelope. Thermal equilibrium turned out to be usually impossible. Lucy thermal equilibrium model (violating the period-colour relation) only in a limited mass region and for extreme Population I composition ($Z\\ga 0.04$). \\citet{mos70} and \\citet{whe72} encountered similar difficulties. Indeed, thermal equilibrium was found to be in conflict with the equal entropy condition in zero age contact binaries \\citep{kah95}.  Renouncing the restriction to thermal equilibrium, \\citet{luc76,fla76}, and \\citet{rob77} obtained contact binary solution evolving in thermal cycles about a state of marginal contact. The solutions are in agreement with the period-colour relation, but they have bad light curves because of a large temperature difference between the two components for a considerable part of the time in a cycle. Some authors encountered the same difficulty, the so-called light curve paradox, also for systems evolving without loss of contact. \\citet{haz01} pointed out that the present theory of cyclic contact binaries (thermal relaxation oscillation, TRO) is not in a position to resolve the light curve paradox. The difficulty of explaining the EW-type light curves (i.e. light curve paradox) reflects a basic conflict between the treatment of the internal structure and the observed properties of the outermost layers.  \\citet{kah02a,kah02b} has discussed the structure equations of contact binaries, and assumed that the energy sources/sinks caused by the interaction of the components occurs only in the secondary's/primary's outer layers to obtain the models which are applied to the typical late-type systems. First he imposed the restriction that the fractional extent in mass of the sources/sinks in the layers above the critical surface is the same in the both components, and found solutions evolving in thermal cycles. And if the energy transfer is assumed to be sufficiently effective, loss of contact is avoided. Otherwise, the cycles consists of a long quiet phase in good contact and a short violent phase with rapid changes between contact, semi-detached, and detached configurations. And in both cases a large potential difference between the surfaces of both components occurs during a part of the contact phase, that is to say, the system deviates hydrostatic equilibrium during this time. Meanwhile, the maximum contact degree is extremely small ($\\sim 2\\%$) in both cases, suggesting that late-type contact binaries must be shallow contact. However, the contact degree of most of real W UMa systems which were observed to be in good thermal contact is about $10\\%$, only a little of real late-type contact systems (ER Cep, AO Cam, Maceroni \\& van't Veer, 1996) is observed to be in good thermal contact with a smaller contact degree ($\\sim$ 2 per cent).  In a series of papers, Shu and his colleagues have extended the unequal entropy model by putting an equal entropy common envelope on top of interiors with the entropy difference, $\\Delta S = 0$, \\citep{shu76, shu79, shu80,lub77,lub79}. These models  satisfy the light curve constraint but requires a temperature discontinuity region between the common envelope and the interior of the secondary component. This so-called DSC model has been strongly attacked by other investigators \\citep{haz78,pap79,smi80} since it was thought to be violating the second law of thermodynamics.  The structure and evolution of W UMa stars still comprises many unsolved questions although many progresses have been achieved. The most difficult problem concerns the energy transfer between the components, i.e., it is not clear where and how the energy is transferred. The mechanism causing energy transfer is still a largely unsolved problem. In addition, we do not confirm that the transfer occurs in the base, or the outermost layers, or the whole of common envelope although it seems probable that the transfer occurs in the common envelope, above the inner Roche critical surface, where the stars are in good contact.  In present paper, we have calculated a relative-radius grid of contact binaries with different mass ratios and different contact degrees according to the Roche potential to ensure the surfaces of two components lying on the same equipotential by interpolation using this radius grid. In addition, we assume that the energy transfer occurs in the different regions of common envelope to investigate the effect of region of energy transfer on the structure and evolution of contact binaries. As result of our investigation, we find that the region of energy transfer has a significant influence on the structure and evolution of the contact binaries, and that the energy transfer may take place in the outermost layers of the common envelope for W-type systems, and this probably takes place in the deeper layers of the common envelope for A-type systems. ",
        "conclusions": "The most uncertainties in the theoretical investigation of contact binaries concerns mainly the energy transfer, i.e., it is not clear that where and how this energy is transferred from the primary (more massive star) to the secondary. Although it seems probable that the transfer occurs in the common envelope, above the Roche lobe, where the two stars are in good contact. However, we can not confirm that the energy transfer occurs in the base or the outermost region or the whole of the common envelope. Therefore, most of results of the previous investigators are based on the assumption that the energy transfer between the components occurs in the adiabatic part of the common convective envelope (adiabatic transfer). The modification due to superadiabatic transfer for W-type systems is small \\citep{whe72,bie73}. Extreme superadiabatic transfer \\citep{mos73} gives more freedom, but \\citet{haz74} showed that this case must be excluded because the heat capacity of the secondary's subphotospheric layers is too small. In order to investigate the effect of the region of energy transfer on the structure and evolution of contact binaries, we assume the energy transfer takes place in the different regions of the common envelope, and find that the region of energy transfer has significant influence on the evolution of the contact binaries. There is no possibility of the temperature of the secondary exceeding that of the primary at any time of a cycle if the energy transfer occurs in the base or the whole of the common envelope, and the temperature of the secondary can excess that of the primary if the energy transfer takes place in the outermost layers of the common envelope, suggesting that during contact phase of the cycle the models are reasonably good agreement with the observed properties of A-type W UMa systems when the energy transfer is assumed to take place in the base or the whole of common envelope, and the model is very consistent with the observed properties of W-type systems when the energy transfer is assumed to take place in the outermost layers of the common envelope. Therefore, we conclude that the energy transfer may take place in the base of the common envelope for A-type systems, and the transfer takes place in the outermost layers for W-type systems. Since the surfaces of both components are radiative rather than convective, it follows that convection is by no means essential to heat transport in contact envelopes, although it may well have an important influence.  The W UMa systems have been divided by \\citet{bin70} into A-type and W-type systems according as the primary minimum in the light curve is a transit, in which the smaller star partially eclipses the larger, or an occultation, when the larger star is in front. The observations of W UMa stars have been reviewed by \\citet{mon81}. The main differences between these subclasses are as the followings: the A-type systems have longer periods, are hotter, smaller mass ratio, and in better contact (i.e. the higher contact depth). Meanwhile, \\citet{wil78} has found that the A-type systems are likely evolved stars. For a low total mass contact binary, we assumed that the energy transfer takes place in the outermost layers of the common envelope. Our model indicates that the mass ratio of the binary becomes smaller and smaller, and the contact depth becomes higher and higher as the evolution of the system, that is to say, the system steadily evolves towards a contact binary with a smaller mass ratio and a deeper common envelope. It suggests that some A-type systems with low total mass could be considered as later evolutionary stages of W-subtypes, although W-type systems show greater activity in the form of period and light curve changes, which were originally thought to occur on a thermal timescale \\citep{ruc74} but now their periods appear to consists of abrupt changes typically in a timescale of a few months; And A-type systems possess more stable light curves and less rapidly varying periods. However, in this scenario the system evolution was towards smaller mass ratio, longer periods and deeper contact, and the better physical contact will lead to the gradual disappearance of the W-type peculiarities (hotter secondary, light curve perturbations, and frequent period changes). Therefore, some of low total mass systems can be regarded as the later evolutionary stages of W-type systems.  In order to obtain a model of W-subtype W UMa system in which the temperature of the less massive star is higher than that of the massive one, the investigators assumed that the energy transfer takes place in the superadiabatic part or even in extreme superadiabatic region of the envelope of the secondary. However \\citet{haz74} suggested that the extreme superadiabatic transfer must be excluded because the heat capacity of the secondary's subphotospheric layers is too small. In our model, we do not assume that the energy occurs in the superadiabatic or extreme superadiabatic part of the envelope of the secondary, but the surface temperature of the secondary can excess that of the primary during the time when the radius of the primary increases rapidly, or the radius of the secondary (less massive component) decreases rapidly (see figure 9a,b). The temperature of the secondary exceeding that of the primary in our model can be attributed to exhaustion of a part of the nuclear luminosity of the primary due to the expansion of the primary, or the release of the gravitation energy of the secondary because of its contraction. Meanwhile, it suggests that the two subtype W UMa systems are probably caused by the energy exchange between the gravitational (potential) energy and thermal one. \\citet{wan94} has showed that the two types of contact binaries are in two different TRO states: the less massive component of W-types are shrinking whereas the less massive star of A-types are swelling. The contraction of the secondary in W-type systems releases some of its gravitational energy, therefore, it makes the effective temperature of the secondary higher than the primary. However, In our model, the contraction of the secondary is one of the mechanisms which can make the temperature of the secondary higher than the primary, but it only lasts a very short period, so the surface temperature of the secondary higher than the primary is mainly caused by the depletion of a part of the luminosity of the primary due to its rapid expansion.  Our models for contact binary systems exhibit cyclic behavior about a state of marginal contact, with a period of $10^7$ yr. In cyclic evolution, the two components of the system are unlikely in thermal equilibrium because of the difference in their thermal timescales, but they attempt to reach the thermal equilibrium. The attempt of the system to reach an nonexistent thermal equilibrium, coupled with Roche geometry, is driver for cycling behavior \\citep{rah81}. A larger temperature difference ($\\Delta T_{\\rm eff}> 300 {\\rm K}$) between the two components occurs in a part of the time of a cycle (lasting about 30$\\sim$35 percent time of a cycle). Almost all of the previous investigators thought that this requires there to be as many short-period binary with EB light curves as with EW light curves, and that the models for contact binaries encounter a difficulty, so called light curve paradox. \\citet{ruc02} gives 13 EWs and 5 EBs (and 14 ELLs, which have too small an amplitude to be classified as EWs or EBs). It is reasonable to identify the EWs as contact binaries and the EBs as semi-detached. The ratio of 5/13 is not much out of line with TRO theory, so we are not sure there is in fact any light curve paradox. However, the W UMa systems indeed undergo angular momentum loss without doubt, and the most likely angular momentum loss mechanism is magnetic braking \\citep{hua66,mes68}. The {\\it Einstein} X-ray observations and {\\it IUE} ultraviolet observations \\citep{eat83} showed that W UMa systems are strong sources, suggesting surface activity of the kind we observe on the Sun, and so the presence of the magnetic fields. The stellar wind would cause magnetic braking, and we will discuss the evolution of W UMa systems included the angular momentum loss in our future work. Meanwhile, we do not consider the energy source at secondary's atmosphere provided by the accreting matter from the primary at our present work. This energy source can hasten the expansion of the secondary, and shorten the time spent in the semi-detached evolution. We refer to a forthcoming work included this extra energy source for the secondary."
    },
    "0403/astro-ph0403517_arXiv.txt": {
        "abstract": "{ We observed the pre--stellar core L1521F in dust emission at 1.2mm and in two transitions each of \\NTHP, \\NTDP, \\CEIO \\ and \\CSEO \\ in order to increase the sample of well studied centrally concentrated and chemically evolved starless cores, likely on the verge of star formation, and to determine the initial conditions for low--mass star formation in the Taurus Molecular Cloud.  The dust observation allows us to infer the density structure of the core and together with measurements of CO isotopomers gives us the CO depletion. \\NTHP \\ and \\NTDP \\ lines  are good tracers of the dust continuum and thus they give kinematic information on the core nucleus. We derived in this object a molecular hydrogen number density  n(\\HH)$\\sim  10^6$\\percc \\ and a CO depletion factor, integrated along the line of sight, $f_{\\rm D}$ $\\equiv$ 9.5$\\times$10$^{-5}$/$x_{\\rm obs}$(CO) $\\sim 15$ in the central 20\\arcsec , similar to the pre--stellar core L1544.  However, the $N(\\NTDP )/N(\\NTHP )$ column density ratio is $\\sim$ 0.1, a factor of about 2 lower than that found in L1544.  The observed relation between the deuterium fractionation and the integrated CO depletion factor across the core can be reproduced by chemical models if \\NTHP \\ is slightly (factor of $\\sim$2 in fractional abundance) depleted in the central 3000 AU. The \\NTHP \\ and \\NTDP \\ linewidths in the core center are $\\sim$ 0.3 \\kms , significantly larger than in other more quiescent Taurus starless cores but similar to those observed in the center of L1544. The kinematical behaviour of L1521F is more complex than seen in L1544, and a model of contraction due to ambipolar diffusion is only marginally consistent with the present data. Other velocity fields,  perhaps produced by accretion of the surrounding material onto the core and/or unresolved substructure, are present. Both chemical and kinematical analyses suggest that L1521F is less evolved than L1544, but, in analogy with L1544, it is approaching the ``critical'' state. ",
        "introduction": "The pre--stellar core L1544 has recently been the subject of much study because while apparently in hydrostatic equilibrium, there are indications that it is close to the ``critical'' state at which it will become gravitationally unstable and from which it will dynamically collapse (see discussions by Tafalla et al.~\\cite{taf98}; Ciolek \\& Basu~\\cite{cb2000}; Caselli et al.~\\cite{cas02a}; Caselli et al.~\\cite{cas02b}). If correct, this is of fundamental importance because it defines the initial conditions for the formation of a protostar, which affects many theoretical studies of low mass star formation.\\\\ Clearly  trying to extrapolate general trends from a single object is difficult and a larger number of L1544-like cores (preferably with the same external environment) should be studied. Unfortunately there are rather few other objects with similar properties due to the short timescale of this phase. According to the Ciolek \\& Basu (\\cite{cb2000}) model, for example, (contraction of a disk driven by ambipolar diffusion) L1544--like  properties fit the model structure of a core at times~3-$8\\times 10^4$ years prior to the collapse, after an evolution of $2.6\\times10^6$ years. Thus in that particular model, L1544 finds itself in the last few percent of its evolution prior to becoming a protostar. While this may be a somewhat too literal interpretation of the model results, it shows that ``L1544--type cores'' should be relatively rare.  Further progress requires the definition of what is a L1544-like core. One answer is to use current estimates of dust emission and absorption selecting cores of dust extinction upwards of 50 mag. Another approach is to say that cores  which show signs of infalling gas (as does L1544, see Williams et al.~\\cite{Wil99}; Tafalla et al.\\cite{taf98}) are ``L1544-twins''.  This latter indicator is complicated by the fact that at the high densities found in the nuclei of cores similar to L1544, many molecular species and in particular CO and CS freeze--out onto dust grain surfaces (see Kramer et al.~\\cite{kra99}; Caselli et al.~\\cite{cwt99}; Bacmann et al.~\\cite{bac02}; Bergin et al.~\\cite{berg02}; J\\o rgensen et al.~\\cite{jsv02}; Tafalla et al.~\\cite{taf02}); observing such tracers implies observing the low density surrounding envelope.  However, recent studies indicate that species whose abundance is linked to that of molecular nitrogen such as N$_{2}$H$^{+}$ and NH$_{3}$ (as well as their deuterated counterparts) do not condense out in the same  fashion and hence can be used as tracers of the dense gas (Bergin \\& Langer~\\cite{bl97}). The extent to which this is true is debatable but it is a useful hypothesis and substantiated by the general similarity of the spatial distributions seen for example in dust emission and in maps of \\NTHP (Tafalla et al.~\\cite{taf02}, ~\\cite{tmc03}).\\\\  Caselli et al. (\\cite{cas02a},\\cite{cas02b}) have used \\NTHP \\ and \\NTDP \\ to derive the physical, chemical and kinematical properties of L1544. They found that L1544 has a central  \\NTHP ~column density of $1.5 \\times 10^{13}$ cm$^{-2}$ and a column density ratio $N(\\NTDP )/N(\\NTHP )$ of  0.24. The \\NTHP \\ linewidths towards the nucleus (the dust emission peak) are roughly 0.3 km s$^{-1}$ and decrease as one goes to positions away from the center. The  line of sight velocity measured in \\NTHP (1-0) and \\NTDP (2-1) shows a gradient along the minor axis of the elliptical structure seen in 1.3mm dust emission but no clear gradient along the major axis.  In this paper, we will study another core in the Taurus complex, L1521F (at an assumed distance of 140 parsec) using the same approach as in our study of L1544. Repeating the L1544 study carried out by Caselli et al. (\\cite{cas02a}, \\cite{cas02b}) is important because it allows us to check  to what extent L1544 is an exceptional case. In order to do this we need another source which has the same general characteristics as L1544. The source selection was made using some preliminary results we obtained in a survey carried out at the IRAM-30m telescope. L1521F stood out as being the only core in Taurus, besides L1544, with strong \\NTDP (2-1) emission compared to  \\NTHP (1-0). This suggests enhanced deuterium fractionation implying an advanced evolutionary state (Caselli et al.~\\cite{cas02b}).  Previous observations of this object have been carried out by Mizuno et al. (\\cite{moh94}), Onishi et al. (\\cite{omk96}), Codella et al. (\\cite{cod97}), and Lee et al. (\\cite{lom99}). Onishi et al. (\\cite{omf99}) also studied L1521F (which they call MC 27), and found a high central density, suggesting that this is the most evolved starless condensation in Taurus. L1521F was also noted by Lee et al.~(\\cite{lmt99}) as a strong infall candidate, in their survey of CS and \\NTHP \\ lines in starless cores, although later mapping of the two tracers  has shown extended ``red'' asymmetry in the CS(2--1) profiles (Lee et al.~\\cite{lmt01}).  In section 2 of this paper, we describe our observational procedure. In section 3 we present the observational results deriving the physical characteristics of the source and analysing its chemical and kinematical properties. In section 4 we discuss the observational results and the summary can be found in Section 5. ",
        "conclusions": "We have analysed physical and chemical properties of L1521F, a starless core in the Taurus Molecular Cloud, with characteristics similar to the pre--stellar core L1544.  The main similarities and differences between the two cores are listed below:  1.  the dust emission distributions are similar, implying a fairly closely matched density structure, with central densities of $n_0$ $\\sim$ 10$^6$ \\percc , the radius of the ``flat'' region $r_0$ = 20\\arcsec , and similar asymptotic power index $\\alpha$ (see Sect.~\\ref{sdensity} and Tafalla et al.~\\cite{taf02}). In particular, the aspect ratio is quite similar: 1.6 and 1.8 in L1521F and L1544, respectively.  2. The line width decreases with distance from the cloud center ($\\sim$ 0.3 \\kms ) to 80\\arcsec \\ away ($\\sim$ 0.25 \\kms ; see Fig.~\\ref{dv_b}), in analogy with L1544, and consistent with the predictions of ambipolar diffusion models, although any gravity--driven contraction in 2-3D is expected to get localized line broadening. The particular model which best fits the data will be investigated in the future.  3. The amount of CO freeze--out (integrated CO depletion factor $f_{\\rm D}$ = 15) is also comparable to L1544, as is the column density of \\NTHP \\ toward the dust peak ($\\simeq$ 1.5$\\times$10$^{12}$ cm$^{-2}$).  4. The deuterium fractionation toward the L1521F dust peak ($R_{\\rm deut}$ $\\sim$ 0.1) is however a factor $\\sim$ 2 smaller than in L1544, due to the (factor of 2) smaller column density of \\NTDP \\ toward L1521F.  This can be understood if L1521F is less chemically evolved than L1544, with a smaller ($r < 2000$ AU) central molecular hole.  5. Unlike in L1544, the velocity field in L1521F maintains a complex structure even at the small scales traced by \\NTDP \\ and \\NTHP (3--2) (see Figs.~\\ref{fhgrad},\\ref{fdgrad}).  This may be due to the presence of clumps in the central few thousand AU (as deduced by the interferometric observations of Shinnaga et al.~\\cite{sol03}), but could also be caused by normal mode pulsations, as in the case of B68 studied by Lada et al. (\\cite{lba03}). The ambipolar diffusion model with infall of Ciolek \\& Basu (\\cite{cb2000}) has problems in reproducing the whole velocity field observed across the core. This may be due to the fact that part of the observed bulk motions result from residual core contraction, as suggested by Tafalla et al.~(\\cite{tmc03}) in their study of L1517B and L1498, thus preventing a clear determination of the velocity field within the core nucleus.  6. The line profiles in L1521F show asymmetric structure, although the two peaks are not well separated as in L1544. This is consistent with smaller (factor of ~1.5) infall velocities in the central few thousand AU of L1521F.  In any case, the large central density ($\\sim$ 10$^6$ \\percc ) and the evidence of central infall (broader line widths toward the center and central infall asymmetry in \\NTHP (1--0)) indicate that L1521F is another starless core on the verge of star formation, or a pre--stellar core. Although a study of a more complete sample is needed, assuming that L1544 and L1521F are the only two cores in Taurus close to dynamical collapse, and given that the total number of starless cores in Taurus is 44 (Onishi et al. \\cite{omk02}), we can argue that this process of contraction towards the ``critical'' stage, or the ``L1544--phase'', may last about five percent of the core lifetime.  A more detailed analysis of \\NTHP \\ line profiles will be presented in a future paper, where the observed chemical abundances will be introduced in a non spherically symmetric Monte Carlo radiative transfer code. This is needed to understand the nature of the \\NTHP (1--0) line asymmetry, which may be caused by self--absorption plus infall, or to infall plus a molecular hole, or to the relative motion of different clumps, or to a mixture of the above possibilities."
    },
    "0403/astro-ph0403451_arXiv.txt": {
        "abstract": "{The X-Ray Flash (XRF), 031203 with a host galaxy at $z=0.1055$, is, apart from GRB\\,980425, the closest $\\gamma$-Ray Burst (GRB) or XRF known to date.  We have monitored its host galaxy from 1--100\\,days after the burst. In spite of the high extinction to the source and the bright host, a significant increase and subsequent decrease has been detected in the apparent brightness of the host, peaking between 10 and 33\\,days after the GRB.  The only convincing explanation is a supernova (SN) associated with the XRF, SN2003lw.  This is the earliest time at which a SN signal is clearly discernible in a GRB/XRF (apart from SN1998bw). SN2003lw is extremely luminous with a broad peak and can be approximately represented by the lightcurve of SN1998bw brightened by $\\sim0.55$\\,mag, implying a hypernova, as observed in most GRB-SNe.  The XRF--SN association firmly links XRFs with the deaths of massive stars and further strengthens their connection with GRBs. The fact that SNe are also associated with XRFs implies that \\emph{Swift} may detect a significant population of intermediate redshift SNe very soon after the SN explosions, a sample ideally suited for detailed studies of early SN physics. ",
        "introduction": "} It is now firmly established that at least some long-duration $\\gamma$-ray bursts (GRBs) are accompanied by the contemporaneous explosion of a supernova \\citep[SN, e.g.][]{2003Natur.423..847H,2003ApJ...591L..17S,2003A&A...406L..33D,1998Natur.395..670G}, consistent with expectations for some models of GRBs involving the collapse of massive stars \\citep{2001ApJ...550..410M,1999ApJ...524..262M}.  The lack of large numbers of SN--GRB associations may be explained, at least in part, by the difficulty of obtaining the optical spectra of SNe with redshifts usually greater than unity, against the combined backgrounds of the fading afterglow and the host galaxy.  X-Ray Flashes (XRFs), a class of very soft bursts, was discovered with \\emph{BeppoSAX} \\citep{2001grba.conf...16H}. They are intense, short-lived flashes of soft X-rays of extragalactic origin \\citep{2003ApJ...599..957B,2003astro.ph.11050S,2004astro.ph..2085P}, and may be defined by a larger X-ray than $\\gamma$-ray fluence in the burst \\citep[$S_{\\rm X}/S_\\gamma>1$,][]{2003astro.ph.12634L}.  The similarity in the durations of XRFs and GRBs \\citep{2001grba.conf...16H,2003A&A...400.1021B}, the continuum of spectral properties observed between the two classes \\citep{2003astro.ph.12634L}, their cosmological origins in each case \\citep{2003ApJ...599..957B,2003astro.ph.11050S}, and the similarity of their optical and X-ray afterglows \\citep{2004astro.ph..2240F,2004astro.ph..1225W}, makes it seem probable that XRFs and GRBs have a similar origin. While GRBs and XRFs are located at cosmological distances, few have been located at redshifts $<0.3$.  They are: GRB\\,030329 at $z=0.1685$ \\citep[associated with SN2003dh,][]{2003Natur.423..847H,2003ApJ...591L..17S}, XRF\\,020903 with its probable host galaxy at $z=0.251$ \\citep{2003astro.ph.11050S}, GRB\\,980425 probably associated with SN1998bw at $z=0.0085$ \\citep{1998Natur.395..670G} and recently, the XRF referred to as GRB\\,031203 at $z=0.1055$ \\citep{2004astro.ph..1225W,2004astro.ph..2085P}.  Of these, XRF\\,020903 had a very low peak spectral energy and a low luminosity \\citep{2004ApJ...602..875S,2003astro.ph.11050S} and GRB\\,980425 had an extraordinarily low luminosity \\citep{1998Natur.395..663K}.  That GRB\\,031203 was in fact an XRF was discovered because of the detection of a transient, outwardly moving ring of X-ray emission surrounding the afterglow \\citep{2004ApJ...603L...5V}.  This was interpreted as reflection of the original burst event and early afterglow off dust sheets in the Galaxy, from which strong lower limits on the prompt soft X-ray fluence were obtained \\citep{2004astro.ph..1225W}.  The high extinction toward GRB\\,031203 \\citep[$(\\,\\rm E(B-V)\\approx1$,][]{1998ApJ...500..525S,2004astro.ph..2085P}, though instrumental in allowing the detection of the dust reflection halo, also hampered attempts to follow the afterglow at optical wavelengths and no optical or infrared afterglow was detected. The location of the burst is therefore determined from X-ray and radio detections.  Both of these locate GRB\\,031203 unambiguously on a sub-luminous, blue, and strongly star-forming galaxy at fairly low redshift \\citep[$z=0.1055$,][]{2004astro.ph..2085P}, where the probability of a chance association with such a galaxy is not very significant \\citep{2004astro.ph..1225W,2004astro.ph..2085P}.  Strong evidence for the association of XRFs with the deaths of massive stars was present in the lightcurve of XRF~030723 \\citep{2004astro.ph..2240F}, but the analysis of that burst was complicated by the lack of a redshift. Because of the low redshift there was considerable interest in attempting to discover a SN associated with GRB\\,031203, in particular since it has been found to be an XRF and the host galaxy was therefore monitored independently by a number of groups in order to find the photometric variability that would indicate a SN \\citep[e.g. this paper;][]{2003GCN..2544....1B,2003GCN..2545....1T,2004astro.ph..3510C,2004astro.ph..3608G}. This would be the first SN associated with an XRF with a spectroscopic redshift, firmly establishing the association of XRFs with the deaths of massive stars and confirming the suspected link between XRFs and GRBs.  In Sect.~\\ref{observations} we describe I-band imaging observations of the host galaxy of GRB\\,031203 (HG\\,031203) over the first 100\\,days since the burst, and in Sects.~\\ref{results} and \\ref{discussion} the discovery of a SN \\citep[named SN2003lw, ][]{2004IAUC.8308....1T} associated with the XRF and the implications of this discovery. This paper supersedes an earlier preliminary report on some of these observations \\citep{2003GCN..2493....1H}.  A cosmology where $H_0=75$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_\\Lambda = 0.7$ and $\\Omega_{\\rm m}=0.3$ is assumed throughout. ",
        "conclusions": "} We have monitored the host galaxy of the XRF, GRB\\,031203 in the near-infrared, from 1--100\\,days following the burst.  In spite of the bright host galaxy and high extinction, we have discovered positive evidence of a SN, peaking $\\sim20$\\,days after the XRF and can clearly trace the early SN rise.  At $z=0.1055$, this is the closest GRB/XRF-associated SN discovered so far, after SN1998bw, and the first SN associated with an XRF with known redshift. This confirms the strong case for an association between XRFs and SNe found in XRF\\,030723 \\citep{2004astro.ph..2240F}. The SN appears to have a somewhat higher peak luminosity than observed in SN1998bw, but the lightcurve is otherwise fairly similar, implying that the SN accompanying GRB\\,031203 was a hypernova.  It is likely that \\emph{Swift} will detect a significant population of faint bursts like GRB\\,031203 and hence allow the study of (Type Ib/c) core-collapse SN at much earlier times than what has been possible so far; this may have a substantial impact on SN research."
    },
    "0403/astro-ph0403498_arXiv.txt": {
        "abstract": "It is well known that the energy input from massive stars dominates the thermal and mechanical heating of typical regions in the interstellar medium of galaxies. These effects are amplified tremendously in the immediate environment of young massive star clusters, which may contain thousands of O and B stars. We present models of star cluster formation that attempt to account for the interplay between feedback and self-gravity in forming clusters, with application to systems ranging from the Orion Nebula Cluster, the Galactic Center clusters like the Arches, and super star clusters. ",
        "introduction": "Star formation in disk galaxies appears to be dominated by a clustered mode (Lada \\& Lada 2003), occurring in regions of the disk that are gravitationally unstable (Martin \\& Kennicutt 2001). The most unstable mass scales are about $1.3\\times 10^7 (\\sigma_{\\rm gas}/6\\;{\\rm km\\:s^{-1}})^4 (\\Sigma_{\\rm gas}/10\\;M_\\odot {\\rm pc}^{-2})^{-1} M_\\odot$, roughly consistent with the observed cutoff of the Galactic Giant Molecular Cloud (GMC) mass function (Williams \\& McKee 1997). GMCs are gravitationally bound (Solomon et al. 1987), and there is a comparable amount of atomic gas associated with these structures (Blitz 1990). Altogether, a significant fraction ($\\sim 1/2$) of the total gas mass inside the solar circle is organized into bound structures.  The lifetimes of GMCs are not well known. Blitz \\& Shu (1980) used observational and theoretical arguments to infer a GMC lifetime of a few times $10^7$~yr; their argument based on photoevaporation of the clouds was confirmed in a more extensive calculation by Williams \\& McKee (1997). The rocket effect associated with the photoevaporation displaces the molecular gas from regions of active star formation at the same time the gas is being transformed from molecular form to atomic and ionized form. Leisawitz, Bash \\& Thaddeus (1989) found that open star clusters older than about $\\sim 10$~Myr were not associated with molecular clouds, which is consistent either with post-star-formation cloud lifetimes shorter than this age or with relative velocities of star clusters and their parent clouds of about $10\\:{\\rm km\\:s^{-1}}$. Elmegreen (2000) presented a number of arguments for star formation to occur in about $1-2$ dynamical timescales, particularly from the small age spreads in clusters such as the Orion Nebula Cluster and from the statistics of the presence of young stars in regions of dense gas. Hartmann (2003) has argued for such rapid star formation in Taurus. On the other hand, the observation that the angular momentum vectors of GMCs in M33 are small and are both pro- and retro-grade (Rosolowsky et al. 2003) may indicate a relatively long lifetime so that angular momentum can be shed by magnetic braking (e.g., Mestel 1985) and/or cloud-cloud collisions in a shearing disk (Tan 2000).  Most star formation in GMCs is concentrated in {\\it clumps} that occupy a relatively small fraction of the volume and have a small fraction of the mass. Within clumps are over-dense regions that we refer to as {\\it cores} that form individual stars or binaries (Williams, Blitz, \\& McKee 2000). Star-forming clumps are identified by $\\rm H_2O$ masers, outflows, and far infrared and radio continuum emission (Plume et al.  1997; Hunter et al. 2000; Sridharan et al. 2002; Beuther et al. 2002; Zhang et al. 2002; Mueller et al. 2002; Shirley et al. 2003). Most of the 350~$\\rm \\mu m$ emission maps of Shirley et al. (2003) have morphologies consistent with quasi-spherical, virialized distributions.  Typical properties of clumps are masses $M\\sim 100-10^4\\: M_\\odot$, diameters $\\sim 1$~pc and surface densities $\\Sigma\\sim 1\\:{\\rm g\\:cm^{-2}}$.  Clump properties are very similar to those of the smaller infrared dark clouds (IRDCs) (Carey et al. 1998), which are dense, cold regions embedded in GMCs. Thus, understanding the origin of IRDCs is crucial for progress in the fields of both star cluster formation and galaxy formation and evolution. Carey et al. noted nongaussian emission line profiles of $\\rm H_2CO$ from 9 out of 10 IRDCs. This may indicate short formation timescales via a triggering mechanism. IRDC morphologies are also more filamentary than the sub-mm emission maps of star-forming clumps, which suggest that the latter represent a more evolved stage.  Possible triggers for IRDC formation include compression of parts of GMCs (which are likely to contain clumpy substructure) by shock waves driven by cloud-cloud collisions (Scoville, Sanders \\& Clemens 1986; Tan 2000), spiral density waves (Sleath \\& Alexander 1996), ionization fronts (Elmegreen \\& Lada 1977; Thompson et al. 2004), supernovae (e.g. Palous, Tenorio-Tagle \\& Franco 1994) or stellar winds (e.g. Whitworth \\& Francis 2002). Clumps that were previously pressure-confined and gravitationally stable may suddenly become unstable. The contraction of the cloud can be halted by the onset of star formation (McKee 1989). We expect that the protocluster will come into an approximate equilibrium in which pressure support balances self-gravity. Some pieces of evidence in support of this view are the approximately spherical morphologies of star-forming clumps (Shirley et al. 2003); the timescales of star cluster formation estimated from outflow momentum generation rates (Tan \\& McKee 2002); and the empirical age spreads of stars in young clusters, such as the Orion Nebula Cluster (Palla \\& Stahler 1999), which are quite long compared to the dynamical or free-fall timescales of clumps, $\\bar{t}_{\\rm ff}=(3\\pi/32G\\bar{\\rho})^{1/2} =1.0\\times10^5 (M/4000M_\\odot)^{1/4}\\Sigma^{-3/4}\\:{\\rm yr}$. This last point requires comment, since Elmegreen (2000) uses the Orion Nebula Cluster as an example of star formation on a dynamical time. He estimated the density in the cluster prior to star formation as $n_{\\rm H}=1.2\\times 10^5$ cm$^{-3}$, corresponding to a free-fall time of $1.25\\times 10^5$ yr. If the star formation occurred over a time $t_{\\rm sf}=10^6$ yr, then $\\eta\\equiv t_{\\rm sf}/\\bar t_{\\rm ff} \\simeq 8\\gg 1$. He used the dynamical time $t_{\\rm dyn}\\equiv R/\\sigma$ to compare with $t_{\\rm sf}$.  If the virial parameter $\\alpha_{\\rm vir}\\equiv 5\\sigma^2 R/GM\\sim 1$, as observed in star-forming regions (McKee \\& Tan 2003), then $R/\\sigma =2.0\\bar t_{\\rm ff}/\\alpha_{\\rm vir}^{1/2} \\simeq 2.5\\times 10^5$ yr. For $t_{\\rm sf}\\simeq 10^6$~yr, star formation in the Orion Nebula Cluster was rapid, but not so rapid that a quasi-equilibrium treatment is invalid. ",
        "conclusions": ""
    },
    "0403/astro-ph0403384_arXiv.txt": {
        "abstract": "We present a direct detection of the growth of large-scale structure, using weak gravitational lensing and photometric redshift data from the COMBO-17 survey. We use deep $R$-band imaging of two $0.5\\times0.5$ square degree fields, affording shear estimates for over 52000 galaxies; we combine these with photometric redshift estimates from our 17 band survey, in order to obtain a 3-D shear field. We find theoretical models for evolving matter power spectra and correlation functions, and fit the corresponding shear correlation functions to the data as a function of redshift. We detect the evolution of the power at the $7.7\\sigma$ level given minimal priors, and measure the rate of evolution for $0<z<1$.  We also fit correlation functions to our 3-D data as a function of cosmological parameters $\\sigma_8$ and $\\Omega_\\Lambda$. We find joint constraints on $\\Omega_\\Lambda$ and $\\sigma_8$, demonstrating an improvement in accuracy by a factor of 2 over that available from 2D weak lensing for the same area. ",
        "introduction": "The evolution of the distribution of matter is a topic of central concern to cosmology. We seek to understand the history of matter environments at all epochs, from the near uniform distribution of matter at $z\\simeq 1000$, to the highly skewed distribution at present. Detailed knowledge of this evolution will lead to precise measurements of cosmological parameters such as the dark energy density $\\Omega_v$, and its equation of state $w$ (e.g. Refregier et al 2003, Benabed \\& van Waerbeke 2003). Furthermore, a comparison of the evolution of dark matter and baryonic matter will afford a deep understanding of the formation of baryonic structures.  There are several approaches available for studying matter evolution. For example, one can compare the matter power spectrum inferred from the Cosmic Microwave Background with that inferred from low-redshift galaxy surveys (e.g. Percival et al 2002, Spergel et al 2003); or if one is principally interested in baryonic structure evolution, one can survey galaxy counts out to higher redshifts and directly observe evolution in the count power spectra (eg Phleps \\& Meisenheimer 2003). Here, we will study another approach: direct measurements of matter power spectrum evolution via weak gravitational lensing combined with photometric redshifts.  Weak gravitational lensing has proved to be a highly useful method of obtaining cosmological information. The weak lensing phenomenon is observed as the slight alignment of neighbouring galaxy images, due to light rays being deflected by gravitational potential fluctuations along their light paths. As the deflection is due to the gravitational environment through which the light passes, the phenomenon is sensitive to all matter, both dark and baryonic. Lensing therefore allows measurement of the distribution of the entire matter content of the universe.  The use of weak lensing (often without redshift estimates for individual galaxies) to constrain the matter power spectrum is now well developed. The shear estimates for large numbers of galaxies have been used to measure the statistics of an effective 2-dimensional shear field; theoretical calculations of these statistics can be fit to the data in order to estimate cosmological parameters (see e.g. van Waerbeke et al 2001, Hoekstra et al 2002, Bacon et al 2003, Brown et al 2003, Jarvis et al 2003) or the statistics can be used to measure the dark matter power spectrum (e.g. Pen et al 2003, who measure the 3-D power spectrum at a fixed redshift).  Now, with the increase in availability of reliable photometric redshifts, there is a great deal of interest in combining redshifts with weak lensing. For example, Wittman et al (2001, 2003) have used shear and redshifts to measure the mass and redshift of two clusters directly from the 3-D shear field. In addition, several recent theoretical studies examine how to directly reconstruct the full 3-D gravitational potential from lensing together with photometric redshifts (Taylor 2001; Hu \\& Keeton 2003; Bacon \\& Taylor 2003); this will allow examination of peak statistics to obtain cosmological parameters, and studies of galaxy formation as a function of environment.  A more directly statistical approach to the shear estimators is to use redshifts to permit shear power spectrum tomography (e.g. Seljak 1998, Hu 1999, 2002, Huterer 2002, King \\& Schneider 2002b, Heavens 2003). In this methodology, the galaxies are divided into several redshift bins, in each of which shear correlation functions or power spectra are measured; by comparing these to theoretical models, cosmological parameters can be estimated. This paper will utilise an extension of this approach. This methodology also forms the starting point for removal of systematic pollutants of the cosmological shear signal such as intrinsic galaxy alignments (Heymans et al 2004, Heymans \\& Heavens 2003, King \\& Schneider 2002a,b). An alternative 3-D statistical method (particularly for measuring $w$ and $w'$) is provided by Jain \\& Taylor (2003), which is further investigated by Bernstein \\& Jain (2004).  In this paper, we seek to measure directly the evolution of the matter power spectrum. We will achieve this using a maximum likelihood fit between data and theoretical power spectra evolving with redshift. We will also fit evolving power spectra expected for various values of cosmological parameters, in order to demonstrate the power of this method for precise determinations of dark energy parameters with future large lensing surveys.  We will use the COMBO-17 survey (Wolf et al 2001) to measure the dark matter evolution. This survey has already been extensively studied from a weak shear perspective; in particular, it has yielded precise measurements of mass-to-light correlations for a supercluster (Gray et al 2002), measurements of the shear power spectrum arising from large-scale structure (Brown et al 2003), and 3-D gravitational potential maps (Taylor et al 2004). One of the great assets of the COMBO-17 survey is the existence of accurate photometric redshifts ($\\Delta z \\simeq 0.05$) for galaxies with $z\\la 1$. Our previous statistical lensing study with COMBO-17 only used these redshifts to obtain median redshifts for lens and source planes; in this study, we will use the redshifts for each individual galaxy to compare each galaxy pair's distortion with what would be expected theoretically for galaxies at these 3-D positions.  The paper is organised as follows. In section 2, we will describe the construction of the theoretical power spectra which we require. In particular, we discuss the modeling of the underlying evolving matter power spectrum; we go on to discuss the shear power spectra expected given this matter power spectrum. This includes not only power spectra for shear in a redshift shell, but also for cross power spectra for shears at two different redshifts. We also find inversion equations for moving between shear and matter power spectra, and construct phenomenological models to measure evolution of the power spectra.  We proceed in section 3 to describe the COMBO-17 survey, giving details regarding the observations, photometric redshift estimation and weak shear analysis.  Section 4 describes the analysis of the data. We measure the rate of evolution of the matter power spectrum using a simple model. We also find constraints upon the cosmological constant and amplitude of the power spectrum by fitting models with varying values of these parameters to the 3-D shear field data.  Finally, in section 5 we discuss the implications of our results, and discuss the future development of the method on large shear datasets. ",
        "conclusions": "In this paper, we have explored the evolution of large-scale structure for redshifts $z<1$. This has been achieved using weak gravitational lensing together with photometric redshifts for a set of galaxies, both measured upon two fields from the COMBO-17 survey (Wolf et al 2001).  We have described the construction of theoretical models of evolving matter power spectra, both in terms of phenomenology of the growth of structure, and in terms of characterisation of this growth with cosmological parameters. We have then shown how to calculate the shear power spectrum from this matter power spectrum, and vice versa. In particular, we have discussed the importance of the cross power spectrum for shears between redshift shells, and have related phenomenological models for evolution between the matter and shear cross power spectra.  In order to practically implement this theoretical framework, we have described the usefulness of the COMBO-17 survey for 3-D shear analysis, due to its accurate photometric redshifts for galaxies in the survey (with $\\Delta z\\simeq 0.05$ for $0<z<1.0$) and its well-studied shear catalogue.  We proceeded to apply a least squares estimator analysis to the data, finding that best-fit matter power spectra included significant evolution. In particular, we found that the growth rate expected for $\\Lambda$CDM was approximately 50 times more likely than a zero evolution model, in the case where we ignore galaxies with unknown redshifts. However, we found that the best-fit rate of growth understandably depended sensitively upon the present-day amplitude of the power spectrum. We calculated constraints upon $\\Delta^2(z)$, finding a reasonable match with that predicted by $\\Lambda$CDM for $0<z<1$, but allowing larger present-day amplitude of the power spectrum with more dramatic growth rates.  As an alternative route to extracting information from the 3-D shear field, we obtained least squares estimates of the cosmological constant and the normalisation of the power spectrum. We find low values of $\\sigma_8$ in this small area of sky, with best-fit constraints $\\Omega_\\Lambda=5.6-7.7(\\sigma_8\\pm0.07)$ and $\\Omega_\\Lambda=1-0.15(\\sigma_8\\pm0.04)^{-1.5}, \\sigma_8<0.72$ if we use WMAP priors on $\\Omega_m$ or $\\Omega_k$ respectively. We find that the latter constraint is a factor of 2 improvement upon the constraint obtained by 2D lensing for the same area of sky.  These results demonstrate that 3-D statistical lensing is a useful means of examining dark matter evolution. If one is interested in measuring only cosmological parameters, this method is a complementary alternative to measuring 2-D higher-order statistics of the shear field; the question of the comparative (and combined) merits of the two approaches is the subject of ongoing research (c.f. Benabed \\& van Waerbeke 2003, Refregier et al 2003). However, if one seeks a direct measurement of the dark matter evolution itself, 3-D methods such as the one explored here are essential.  With large multi-colour imaging surveys in the next few years such as that of SNAP (Rhodes et al 2004), the potential of this approach will be fully realised, affording measurements of the cosmological constant and dark energy equation of state parameters to unprecedented accuracy, together with precise reconstructions of the growth of the power spectrum as a function of redshift."
    },
    "0403/astro-ph0403667_arXiv.txt": {
        "abstract": "We present a study of McNeil's Nebula, a newly appeared reflection nebula in the L1630 cloud, together with photometry and spectroscopy of its source. New IR photometry compared to earlier 2MASS data shows that the star has brightened by about 3 magnitudes in the near-infrared, changing its location in a J-H/H-K$'$ diagram precisely along a reddening vector. A Gemini NIRI K-band spectrum shows strong CO-bandhead emission and Br$\\gamma$ is in emission, indicative of strong accretion. A Gemini GMOS optical spectrum shows only a red, heavily veiled continuum, with H$\\alpha$ strongly in emission and displaying a pronounced P~Cygni profile, with an absorption trough reaching velocities up to 600 km~s$^{-1}$. This implies significant mass loss in a powerful wind. However, no evidence is found for any shocks, as commonly seen in collimated outflows from young stars. Apparently the eruption has dispersed a layer of extinction and this, together with the intrinsic brightening of the IRAS source, has allowed an earlier outflow cavity to be flooded with light, thus creating McNeil's Nebula. ",
        "introduction": "Wide-field photographic and CCD images have documented the presence of numerous small compact reflection nebulae in star-forming dark clouds. Such nebulae may possibly represent a transitional stage in which a young star goes from being an embedded infrared source to a visible T Tauri star or Herbig Ae/Be star. This process is widely believed to involve powerful bipolar outflow activity, as reflected in the morphology of the nebulae, which are often cometary, with the partially embedded star illuminating an outflow cavity. When observed over longer time-scales, such small reflection nebulae are often found to vary in brightness and occasionally in illumination pattern. Perhaps the earliest known example is NGC 2261, which reflects light from the young star R~Mon and was found to vary by Hubble (1916). More recently, Cohen et al. (1981) demonstrated the variability of the fan-shaped nebula emanating from PV~Cephei. At some stage, the outflow activity of a star for the first time punches a hole in its surroundings, suddenly allowing a beam of light to sweep across the surroundings. Subsequent eruptions in the source will again flood the outflow channel with light. The emergence of the nebula Re~50 in the Lynds 1641 cloud near the embedded source IRAS 05380--0728 was documented by Reipurth \\& Bally (1986). Another case may have been the nebula IC~430 in L1641, which at present is quite small, but at the end of the 19th century appears to have been a very large and much brighter object (Pickering 1890; Strom \\& Strom 1993). ",
        "conclusions": "The two major classes of eruptive variables among pre-main sequence stars are the FUors and the EXors (Herbig 1966, 1977, 1989). In the following we discuss whether LMZ~12 can be linked to one of these two categories.  FUors are characterized by large-amplitude ($\\Delta$V$\\sim$5-6 mag.) brightenings lasting several or many decades, and in the optical they show F-G type spectra without emission lines. In the K-band region, FUors display deep CO-bandhead absorption (e.g., Reipurth \\& Aspin 1997).  Although LMZ~12 has about the right amplitude, it remains to be seen how long it stays bright.  Spectroscopically, however, LMZ~12 looks very different from mature FUors both in the optical and in the infrared. The prominent H$\\alpha$ emission and the strong CO emission that are presently seen do not suggest a classification as a FUor as we currently understand the phenomenon. On the other hand, Brice\\~no et al. (2004) emphasize that the earliest optical spectrum after outburst of the FUor V1057 Cyg had some similarities to that of LMZ~12.  Further photometric and spectroscopic monitoring of the star is required to settle this.  EXor eruptions, which may occur repeatedly in a given star, have amplitudes that can be comparable to those of FUors, but they have durations which are much shorter, from a few months to a few years. Spectroscopic studies of EXors are limited, because of the rarity and shorter durations of these eruptions. The few such studies that have been made (e.g., Lehmann, Reipurth, \\& Brandner 1995; Herbig et al. 2001; Parsamian et al. 2002) all show that EXors in eruption have at least the lower Balmer lines in emission, together with emission lines of He~I and other lines characteristic of the most active T~Tauri stars (Herbig 1962). A high-resolution spectrum of the H$\\alpha$ line of SVS~13 in outburst (Eisl\\\"offel et al. 1991) shows a P~Cygni profile rather similar to the one seen in Fig.~5.  However, we do not see any other emission lines than H$\\alpha$ in the spectral range observed (4800-6800~\\AA) in LMZ~12.  Infrared spectra of EXors are even more limited, one case is the eruption of SVS~13, for which a K-band spectrum shows the CO bandheads strongly in emission as well as Br$\\gamma$ in emission (Carr \\& Tokunaga 1992). However, a similar spectrum of EX~Lup shows Br$\\gamma$ in emission, but the CO bandheads and Na~I in absorption (Herbig et al. 2001). It thus does not seem that EXors have a unique infrared spectral signature. Based on a comparison with these somewhat limited observations of EXors, we conclude that LMZ~12 does have a certain resemblance to EXors.  The spectroscopic data give us some insight into the physical processes behind the observed eruption.  The detection of Br$\\gamma$ in emission testifies to a region of hot gas. Br$\\gamma$ emission has been shown to correlate tightly with accretion luminosity, so its presence indicates that the eruption is linked to an episode of accretion (Najita et al. 1996; Muzerolle et al. 1998).  The pronounced P~Cygni profile seen at H$\\alpha$ in the optical spectrum is likely to be formed in a strong wind that has sufficient optical depth to produce the deep blueshifted absorption trough (e.g., Muzerolle et al. 2001). Only few T~Tauri stars show such well-developed P~Cygni profiles at H$\\alpha$. The well-defined blue edge of the absorption trough indicates wind velocities of up to 600 km~s$^{-1}$, even more than the extreme wind that emanates from FU~Orionis (e.g., Herbig et al. 2003).  It is noteworthy that we see no spectral features indicative of shocks: the infrared spectrum shows no evidence for H$_2$ emission, and the optical spectrum does not display the [SII] 6717/6731 lines characteristic of Herbig-Haro jets. Our H$_2$ 2.122 $\\mu$m image of LMZ~12 shows no presence of any extended jet flow. It thus appears that LMZ~12, in common with most FUors and EXors, has powerful mass loss, but apparently not in the well-collimated fashion that enables shocked HH jets.   One final insight into the LMZ~12 EXor event comes from the changes of the infrared colors from before to after the eruption. The pre- and post-outburst colors show that LMZ~12 has moved precisely along a reddening vector, indicating that the star has brightened partly because its visual extinction diminished by about 4.5 magnitudes. Using the reddening curve of Rieke \\& Lebofsky (1985), we find that this corresponds to A$_J$=1.26, A$_H$=0.81, and A$_K$=0.50. Since we know that the infrared colors changed by $\\Delta$J=3.64, $\\Delta$H=3.16, and $\\Delta$K=2.87, we see that the intrinsic brightening of the star corrected for the change in extinction is 2.4 magnitudes in each of the J, H, and K filters. Through a combination of this intrinsic brightening and the clearing away of an obscuring layer, the star has been able to illuminate its outflow channel, thus creating McNeil's Nebula."
    },
    "0403/hep-ph0403101_arXiv.txt": {
        "abstract": "We study, analytically and with lattice simulations, the decay of coherent field oscillations and the subsequent thermalization of the resulting stochastic classical wave-field. The problem of reheating of the Universe after inflation constitutes our prime motivation and application of the results.  We identify three different stages of these processes. During the initial stage of ``parametric resonance'', only a small fraction of the initial inflaton energy is transferred to fluctuations in the physically relevant case of sufficiently large couplings.  A major fraction is transfered in the prompt regime of driven turbulence.  The subsequent long stage of thermalization classifies as free turbulence.  During the turbulent stages, the evolution of particle distribution functions is self-similar. We show that wave kinetic theory successfully describes the late stages of our lattice calculation.  Our analytical results are general and give estimates of reheating time and temperature in terms of coupling constants and initial inflaton amplitude. ",
        "introduction": "Field theoretical systems which are a long way from thermal equilibrium have been studied intensely in recent years. The particular problem of how and when such systems approach equilibrium stretches beyond obvious fundamental interest and finds many practical applications. In high-energy physics understanding of these processes is crucial for applications to heavy ion collisions and to cosmology of the early universe. The first topic gains further importance in light of the current and future experimental search for a quark-gluon-plasma at RHIC and at the forthcoming LHC. The second application, our main interest in this paper, is related to the problem of reheating of the universe after cosmological inflation.  Inflation provides a solution to the flatness and the horizon problems of standard cosmology \\cite{lindebook,kolbbook,lythbook} and explains the generation of initial density perturbations - the seeds of galaxies and large-scale structure in our universe. During inflation the universe is in a vacuum-like state. At the end of inflation all energy density is stored in a Bose condensate, the coherently oscillating \"inflaton\" field. This state is highly unstable: parametric, tachyonic or strong non-adiabatic particle creation triggers a fast and explosive decay of the inflaton. This process, dubbed preheating \\cite{Kofman:1994rk,Shtanov:1995ce}, is currently well understood \\cite{Khlebnikov:1996mc,Khlebnikov:1997wr,Khlebnikov:1997zt,Prokopec:1997rr, Kofman:1997yn,Garcia-Bellido:1998wm, Felder:2000hj,Giudice:1999fb}. A generic feature is a strong and fast amplification of fluctuation fields at low momenta, which may lead to various interesting physical effects during and after preheating. These include non-thermal phase transitions \\cite{Kofman:1995fi,Tkachev:1996md, Riotto:1996vi,Khlebnikov:1998sz} with possible formation of topological defects \\cite{Tkachev:1998dc,Kasuya:1997ha,Rajantie:2000fd,Ray:2001cd,Copeland:2002ku}, creation of super-heavy particles \\cite{Felder:1998vq,Giudice:1999fb}, generation of high-frequency gravitational waves \\cite{Khlebnikov:1997di}, etc.  The explosive stage of inflaton decay ends when the rate of interactions of created fluctuations among themselves and with the inflaton becomes comparable to the inflaton decay rate \\cite{Khlebnikov:1996mc,Khlebnikov:1997wr,Khlebnikov:1997zt,Prokopec:1997rr}. The understanding of the subsequent stages of relaxation towards equilibrium, of the thermalization processes and the calculation of the final equilibrium temperature is important for various applications as it links the inflationary phase with that of standard cosmology. Among those one can list baryogenesis \\cite{Kolb:1996jt,Kolb:1998he,Giudice:1999fb,Garcia-Bellido:1999sv, Garcia-Bellido:2001cb} and the problem of over-abundant gravitino production in supergravity models \\cite{Ellis:1982yb,Ellis:1985er,Kallosh:1999jj,Giudice:1999yt,Giudice:1999am, Maroto:1999ch}.  It determines the abundances of other relics, like super-heavy dark matter \\cite{Chung:1998zb,Kuzmin:1998uv,Kuzmin:1998kk,Chung:1999ve}, or axino dark matter \\cite{Covi:2004rb}.  Knowledge of the reheating temperature is also important for fixing constraints on the inflationary model from Cosmic Microwave Background (CMBR) anisotropy \\cite{Liddle:1993fq,Dodelson:2003vq,Liddle:2003as,Feng:2003nt}.  In some models cosmologically important curvature perturbations may be even generated during the process of thermalization \\cite{Hamazaki:1996ir,Dvali:2003em,Kofman:2003nx,Enqvist:2003uk, Matarrese:2003tk}. Last but not least: the reheating temperature should be larger than about one GeV to ensure that the standard Big Bang Nucleosynthesis \\cite{kolbbook,Olive:1999ij} is not hampered.  There have been many efforts and successes in the understanding of the non-equilibrium dynamics and relaxation of field theories, see e.g. Refs. \\cite{Son:1996uv,Semikoz:1996ty,Aarts:2000mg,Davidson:2000er, Salle:2000hd,Calzetta:2000kn,Borsanyi:2000ua,Felder:2000hr,Borsanyi:2000pm, Bodeker:2000pa,Aarts:2002dj,Berges:2002cz,Borsanyi:2003ib,Boyanovsky:2003tc, Baacke:2003bt,Ikeda:2004in}.  However, the leading asymptotic dynamics towards equilibrium remained rather less understood and developed.  The main problem for the theoretical understanding of reheating is that initially the occupation numbers are very large, of order of the inverse coupling constant. In addition, in many inflationary models the zero mode does not decay completely during preheating. Therefore, a simple perturbative approach is not justified.  On the other hand, in this regime, a description in terms of classical field theory is valid \\cite{Khlebnikov:1996mc}, and the whole process (including preheating), can be studied by classical lattice simulations.  Recently, we employed this method to show \\cite{Micha:2002ey,Micha:2003ws} that the classical reheating of a massless $\\Phi^4$-theory in 3+1 dimensions is characterized by a \\textit{ turbulent} and \\textit{ self-similar} evolution of distribution functions towards equilibrium. The shape of the spectra, as well as the self-similar dynamics, could be understood within the framework of wave kinetic theory.  This made it possible to estimate reheating time and temperature, which turned out to coincide parametrically with the results of the simple perturbative approach.  Turbulence appears in a large variety of non-equilibrium-phenomena in nature (see Refs.~\\cite{ZakBook,lvovbook,FrischBook} for a general introduction). It was first discussed for fluids, in the regime of large Reynolds numbers (velocities), where viscosity is subdominant.  Kolmogorov identified turbulence in this regime \\cite{K41:1,K41:2} as a statistically scale invariant flow of spectral energy mediated by vortex interactions.  The same dynamical structure may appear in systems of coupled waves, e.g. on fluid surfaces or for coupled fields in a plasma \\cite{Z67,ZakBook,lvovbook}. In this case the cascade is mediated by wave interactions and the phenomenon has been called \\textit{wave turbulence}.  If there exists an active (stationary) source of energy in momentum space, the turbulence is called \\textit{driven} (stationary).  When the source is switched off after the stage of activity, the freely propagating energy cascade is often referred to as \\textit{free} turbulence. If the kinetic description is applicable, the energy cascade is called \\textit{weak turbulence}. Otherwise one is facing a \\textit{strong turbulence}.  One may expect that the concept of turbulence should be relevant for the problem of reheating \\cite{Khlebnikov:1996mc,Son:1996uv} already on general grounds. Indeed, the source of energy, localized in the \"infra-red\" is present initially.  It is represented by the inflaton field in the problem at hands. To complete the argument, we note that as the final outcome of the evolution one should expect cascading of energy towards a significantly separated region of ``ultra-violet'', high momentum modes.  The goal of the current paper is twofold. First, we want to apply the wave kinetic theory of turbulence to the problem of Universe reheating after inflation. We derive general formulas for the spectra of turbulent distributions and for the self-similar evolution towards equilibrium. This enables us to give asymptotic estimates of reheating time and temperature in Minkowski space as well as in Friedmann universe.  Second, we want to test and confront these ideas to numerical lattice calculations. For our numerical integrations we have chosen the simplest ``chaotic'' inflationary model \\cite{Linde:1983gd}. While the initial ``preheating'' stage in other inflationary models, e.g. in hybrid inflation \\cite{Linde:1994cn} may exhibit important differences \\cite{Garcia-Bellido:1998wm,Felder:2000hj,Micha:1999wv} with this model, we expect the subsequent turbulent stages to be more universal.  We start lattice integration from ``vacuum'' initial conditions for fluctuations in a background of classical oscillating inflaton field.  We observe the initial parametric resonance stage when the energy in fluctuations is growing exponentially with time. This stage terminates when re-scattering of waves out of the resonance band becomes important \\cite{Khlebnikov:1996mc,Khlebnikov:1997zt}.  In the physically relevant case of sufficiently large couplings this happens rather early, when only a small fraction of initial inflaton energy is transferred to fluctuations \\cite{Khlebnikov:1997zt,Prokopec:1997rr,Kofman:1997yn}. At this point a state of stationary turbulence should be established that is driven by the zero-mode. On general grounds, it can be deduced that during this stage the energy in fluctuations should grow linearly with time. This behavior is confirmed by the results of our numerical simulations.  The stage of stationary turbulence should terminate when the energy left out in the zero-mode becomes smaller than the energy stored in created \"particles\".  From this moment of time, the transport of energy from the source is negligible and we observe free turbulence with self-similar evolution of particle distributions towards thermal equilibrium.  The first stage of driven turbulence is prompt and gives the main mechanism by which energy is drawn out of the zero-mode, e.g. out of the inflaton field. The identification of this constitutes one of the new results of the present paper, as opposed to the common opinion that the main mechanism is a \"parametric resonance''. The second stage of free turbulence is very long and can be analytically described as self-similar evolution. This is another new result and diffuses some existing claims and hopes that \"parametric resonance\" may bring a system to thermal equilibrium on a very short time scale.  Overall, the kinetic description and the results of lattice simulations are in rather good agreement with each other.  This indicates that the regime of weak wave turbulence may be already achieved on the lattice.  The paper is organized as follows. In Section \\ref{NumSim:One-Field} we review the results of our numerical simulation of reheating in the simplest $\\lambda \\Phi^4$ model to get familiar with concepts, problems and the typical dynamical behavior of the systems of interest. In Sections \\ref{Turbulentreheating}, \\ref{sec4} we apply the theory of wave turbulence to the problem of reheating in general.  In Section \\ref{TwoFields} we present our numerical simulations.  In Section \\ref{CompLatt2Kin} we compare lattice results with the kinetic approach and discuss the applicability of the latter. In Section \\ref{PhysApplications} we discuss some physical applications of our results, in particular the thermalization in the self-similar regime.  In Appendix \\ref{NumericalProcedure} we give the details of our numerical procedure. In Appendix \\ref{WavesKinetics} we review the derivation of the kinetic equation for a system of weakly interacting classical waves. ",
        "conclusions": "\\label{conclusions}  We have studied the process of thermalization of classical systems, which at some point in their evolution are in a highly non-equilibrium state with energy being concentrated in a deep ``infra-red'' region of momenta. Such states naturally appear e.g.  during reheating of the Universe after cosmological inflation.  We have shown that the process of relaxation in such systems can be divided, in the general case, into three distinct stages.  In the models of the type we have considered in this paper, the initial stage of preheating \\cite{Kofman:1994rk} is powered by  \\textit{parametric resonance}.  During this initial linear stage the rate of energy transfer is the fastest. The energy in particles grows exponentially. However, in the physical situation of reheating after inflation, the coupling constants have to be sufficiently large to insure an acceptably short time-scale of the subsequent thermalization, while with large couplings, only a negligible fraction of the initial inflaton energy is transfered into fluctuations during the parametric resonance stage \\cite{Khlebnikov:1997zt,Prokopec:1997rr}.  We have shown that in such situations the linear stage is followed by the regime of a \\textit{driven stationary turbulence}. During this stage, the energy in particles grows linearly in time. The regime of stationary turbulence stops as soon as the energy in particles starts to dominate the overall energy balance. Therefore, this regime is a major mechanism of energy transfer from the oscillating inflaton zero-mode into other species in realistic models of the type we have considered here.  This period of evolution is also prompt. It should be noted that the source which drives the turbulence is powerful because coherence effects are still strong in the relevant region of momenta.  The subsequent long stage of \\textit{thermalization} classifies as \\textit{free turbulence}.  This stage should be generic. The energy in particles is conserved during this epoch, while the shape of the particle distribution function changes in a self-similar way with the front of the distribution propagating into the ultra-violet. This stage continues until the quantum regime is reached and particles can relax to Bose-Einstein distributions. Applying conventional kinetic theory we have calculated analytically the time needed to equilibrate a system and the resulting temperature in terms of coupling constants and initial inflaton amplitude. The result coincides parametrically with the ``naive'' perturbative estimates \\cite{lindebook} .  We made a comparison of kinetic theory with the numerical integration of scalar field models on the lattice. We show that, at late times, the kinetic approach is applicable, resulting in a weak wave turbulence regime \\cite{ZakBook}.  In the models considered numerically, the evolution is driven by three-particle scattering in the background of  zero-mode oscillations. The characteristic exponents calculated within the framework of wave kinetic theory are consistent with the results of our lattice simulations."
    },
    "0403/astro-ph0403058_arXiv.txt": {
        "abstract": "{ XMM-Newton observations of 19 galaxy clusters are used to measure the elemental abundances and their spatial distributions in the intracluster medium. The sample mainly consists of X-ray bright and relaxed clusters with a cD galaxy. Along with detailed Si, S and Fe radial abundance distributions within 300-700~kpc in radius, the O abundances are accurately derived in the central region of the clusters. The Fe abundance maxima towards the cluster center, possibly due to the metals from the cD galaxy, are spatially resolved. The Si and S abundances also exhibit central increases in general, resulting in uniform Fe-Si-S ratios within the cluster. In contrast, the O abundances are in general uniform over the cluster. The mean O to Fe ratio within the cluster core is sub-solar, while that of the cluster scale is larger than the solar ratio. These measurements indicate that most of the Fe-Si-S and O in the intracluster medium have different origins, presumably in supernovae Ia and II, respectively. The obtained Fe and O mass are also used to discuss the past star formation history in clusters. ",
        "introduction": "Metals in the universe are being produced as a product of stellar evolution and ejected into intergalactic space via stellar winds and supernovae. The location of these metals can be measured by X-ray observations of the intra-cluster medium (ICM).  These observations indicate that the ICM contains an amount of metals comparable to the total amount of metals found in galaxies (e.g. Tsuru \\cite{tsuru91}; Arnaud et al. \\cite{arnaud92}). Therefore, X-ray measurement of the ICM metallicity distribution is a unique way to study the chemical history of galaxies.  Most importantly, the total amount of metals and the abundance ratios among elements (e.g.  O/Fe ratio) constrain the total number of stars that have formed as well as their origin (i.e.  the relative ratios of supernova types; Renzini et al. \\cite{renzini93}).  Measurements from {\\it ASCA} and {\\it BeppoSAX} observations revealed important properties of the ICM metallicity. These include an excess of Fe around cD galaxies (e.g. Fukazawa et al. \\cite{fukazawa00}; De Grandi and Molendi \\cite{grandi01}) and variations of the Si/Fe ratio within a cluster (e.g. Finoguenov et al. \\cite{finoguenov00}) and among clusters (e.g. Fukazawa et al. \\cite{fukazawa98}). The variations in Si/Fe suggest that the metals in the ICM have at least two different origins, presumably SN Ia and SN II. The spatial resolution of these previous instruments was limited, causing the derived total amount of metals to depend on their assumed spatial distribution. The low spatial resolution combined with the limited spectral resolution also introduced a systematic uncertainty in the temperature structure. This is particularly important in the central cool regions. This ambiguity in turn caused severe errors in the abundances of Fe and other elements. Furthermore, in most cases these measurements are limited to the Fe, Si and S abundances.  Recent observations with {\\it Chandra} and {\\it XMM-Newton} have allowed a detailed study of the ICM metals, using an improved spatial resolution and sensitivity not only for the Si, S, and Fe lines but also for the O emission.  For example, Tamura et al.  (\\cite{tamura01b}) measured the spatial distribution of the O, Si, S, and Fe in Abell~496 and found a difference between O and Si-S-Fe. David et al. (\\cite{david01}) used the Chandra data of the Hydra-A cluster to measure the Si and Fe abundances. In addition, the abundance distribution of several elements in the central region (70~kpc in radius) of the Virgo cluster was obtained with high accuracy (Finoguenov et al.  \\cite{finoguenov02}; Gasteldello \\& Molendi \\cite{gas02}; Matsushita et al.  \\cite{matsushita03}).  In this paper, we use the {\\it XMM-Newton} EPIC and RGS instruments and measure the spatial distribution of heavy elements in 19 clusters. This systematic analysis enables us to compare the metal distribution in the sample. In addition, we may reduce the statistical errors on any given measurement when the variation in the sample is negligible. We selected the sample based on their X-ray flux and spatial size. Therefore the sample is limited to nearby X-ray bright clusters with a mean redshift of 0.06.  We obtained the radial distribution of Fe on a sub-arcmin scale for each cluster and confirmed the central increase of the Fe abundance within a typical radius of 100~kpc in most cases.  Furthermore, radial profiles of O, Si, and S relative to Fe have been measured in a fair number of clusters.  We found that the O and Si-S-Fe abundances have different spatial distributions when we combine the results from several clusters.  Throughout this paper, we assume the Hubble constant to be $H_0 = 100h$ km\\,s$^{-1}$\\,Mpc$^{-1}$ and use 90\\% confidence levels unless stated otherwise. Solar abundances are taken from Anders \\& Grevesse (\\cite{anders}) with Fe/H=$4.68\\times 10^{-5}$. ",
        "conclusions": "\\subsection{Present results}  Using spatially-resolved EPIC spectra, we have resolved the temperature and abundance distribution on a sub-arcmin scale in a sample of 19 clusters. Furthermore, using high spectral resolution RGS spectra of 11 cluster cores, we have resolved the soft X-ray spectrum dominated by the Fe-L and the \\ion{O}{viii}~Ly$\\alpha$ lines.  This provides not only tight constrains on the thermal structure but also a robust measurement of the O abundance. Below we summarize our main results.  \\begin{enumerate}  \\item We found no significant variation in the Fe abundance and the Si, S, and O ratios to the Fe among the systems (see Fig.~\\ref{fig:tem-metal}).  \\item The O/H and O/Fe ratio in the cluster cores, (20--60)~$h^{-1}$~kpc in radius, are 0.34$\\pm 0.03$ and 0.63$\\pm 0.05$, respectively.  The r.m.s. deviation for individual clusters as compared to these average values is comparable to the measurement errors on these ratios.  \\item In all clusters with a temperature less than 6~keV, except for the faint cluster MKW~9, we detected a central increase in the Fe abundance; the abundance starts to increase at (50--100)~$h^{-1}$~kpc in radius from 0.2--0.4~solar value to 0.6--0.8~solar value towards the center.  \\item The Si/Fe and S/Fe ratios in cool and medium temperature clusters are radially uniform within (200--500)~$h^{-1}$~kpc with mean values of 1.4$\\pm$0.2 and 1.1$\\pm$0.3, respectively, relative to the solar ratio.  \\item Contrary to the Fe, Si, and S abundances, the O abundance shows no spatial variations. The errors on the O abundance for each cluster are large. However, when we combine the results from several clusters and assume no correlation of the systematic errors of the measurements between clusters, we detected a significant radial variation in the O/Fe ratio; the mean O/Fe ratios are $0.5^{+0.3}_{-0.1}$, $1.7^{+0.9}_{-0.4}$, and $4.1^{+3.5}_{-1.9}$ for (0--50)~$h^{-1}$~kpc,(50--200)~$h^{-1}$~kpc, and (200--500)~$h^{-1}$~kpc, respectively.  \\end{enumerate}  \\subsection{Similarity of the abundance among clusters} {\\it ASCA} and {\\it BeppoSAX} observations have established the Fe increase towards the cluster center in relatively cool clusters with cD galaxies (e.g. Fukazawa et al. \\cite{fukazawa00}; De Grandi and Molendi \\cite{grandi01}). Our measurement not only confirms these results in the cool and medium temperature clusters but also suggests a similar enhancement in hot clusters without a prominent cD galaxy (e.g. Coma), probably because of the improved spatial resolution.  Beside the Fe abundance, we found no variation in the Si and other element ratios to Fe among the clusters. This implies a homogeneous production of the metals in our sample. However, the sample is biased towards medium-temperature X-ray luminous clusters. Therefore we make no conclusion about the universal distribution of the relative abundances. In fact Fukazawa et al. (\\cite{fukazawa98}) claimed an increase in Si/Fe from cool to hot systems using 40 clusters.  \\subsection{The O abundance at the cluster core}  The theoretical models for supernova (SN) types Ia and II predict the O/Fe ratio to be $<0.05$ and $1.5-4$ times solar values, respectively.  On the other hand, the observed O/Fe ratio at the cluster cores are between these two predicted values.  This suggests that the ICM metals at the cluster cores could be produced by a mix of both types of SNe. A value of the O/Fe abundance ratio in the ICM close to solar values supports the idea that the metals in the ICM and the Galaxy were produced in a similar way. Assuming the O and Fe yields of SN Ia and II from Tsujimoto et al.  (\\cite{tsujimoto95}), the observed O/Fe ratio ($\\sim 0.65$) can be reproduced by the relative frequencies of the two SNe, $N_{\\rm Ia}/N_{\\rm II}$, of $\\sim 0.6$. In this model, the mass fractions of the O and Fe originated from SN Ia are 0.05 and $\\sim 0.8$, respectively.  The observed O abundance provides an accurate way to determine the total mass of O in the ICM, $M_{\\rm O}^{\\rm ICM}$.  We estimate $M_{\\rm O}^{\\rm ICM}$ within 50~$h^{-1}$ kpc to range from $10^8$ (NGC~533) to $2\\times 10^9$ (A~1795) in units of $h^{-2.5}$~M$_{\\odot}$ for our RGS sample.  We assume that all oxygen was produced in SNe II.  Then, combining $M_{\\rm O}^{\\rm ICM}$ with a SN II model O yield ($\\sim 2$~M$_{\\odot}$, Tsujimoto et al.  \\cite{tsujimoto95}), gives a total number of SNe II of $10^8-2.5\\times10^9$ for $h=0.7$.  Further assuming that the duration of SNe II activity is $10^7$~yr ( a typical life time of a 20~M$_{\\odot}$ star), we can estimate the SN II rate to be 10--200~SN IIe~yr$^{-1}$.  This rate is at least one or two orders of magnitude larger than that of a typical starburst galaxy.  In other words, the observed O mass requires a large amount of massive star formation, producing the corresponding number of SNe II.  The observed O/H ratio in cluster cores (20--60~kpc in radius) is about half that observed in the central regions (5--15~kpc) of giant elliptical galaxies, like in M87 (Sakelliou et al.  \\cite{sakelliou02}; Gastaldello and Molendi \\cite{gas02}; Finoguenov et al.  \\cite{finoguenov02}) and NGC~4636 (Xu et al. \\cite{xu02}).  The overabundance of oxygen in galaxy centers suggests that in those regions oxygen has been produced via massive star formation and that metals have not been completely mixed with the ambient gas (Gastaldello and Molendi \\cite{gas02}).  \\subsection{The total Fe mass in clusters}  \\begin{figure}[!ht] \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{femass1.qdp.ps}} \\caption{The Fe number density in the ICM. Same notation as Fig.~\\ref{fig:epic-radial-ab}.} \\label{fig:femass1} \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{femass2.qdp.ps}} \\caption{Same as the previous figure, but for the integrated Fe mass in the ICM. } \\label{fig:femass2} \\end{center} \\end{figure}  \\begin{table}[!ht] \\caption{The iron mass to light ratio within $250~h^{-1}$~kpc.} \\label{tbl:imlr} \\begin{center} \\begin{tabular}{lccc} \\hline cluster & $M_{\\rm Fe} $ & $N_{\\rm 0.5}^{\\mathrm{a}}$ & $M_{\\rm Fe}/L_{\\rm B}$ \\\\ & ($10^9$ M$_{\\odot}$) & & ($10^{-2}$M$_{\\odot}$/L$_{\\odot}$)\\\\ \\hline MKW~9\t& 0.5  \t& 13\t& 0.3 $\\pm 0.15$\\\\ MKW~3s\t& 1.5\t& 13\t& 0.9 $\\pm 0.2$\\\\ A~496\t& 2.8\t& 14\t& 1.5 $\\pm 0.5$\\\\ A~1795\t& 6.0\t& 27\t& 1.0 $\\pm 0.2$\\\\ A~754\t& 3.8\t& 29\t& 0.7 $\\pm 0.3$\\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] Number of galaxies within a radius of 250~$h^{-1}$~kpc. \\end{list} \\end{center} \\end{table}  We also calculated the total Fe density and mass in the ICM as a function of radius, based on our measurements of the emission measure and the Fe abundances (Figs.~\\ref{fig:femass1} and \\ref{fig:femass2}).  To compare this Fe mass and the predicted Fe production rate in galaxies, we estimate the iron mass to light ratio (IMLR).  We use the galaxy number within a radius of 250~$h^{-1}$~kpc, $N_{0.5}$ (Bahcall \\cite{bahcall81} and reference therein), and estimate the total optical luminosity of member galaxies based on the relation $L_B = 4.5\\times 10^9 N_{0.5}^{1.42} h^{-2}$~L$_{\\odot}$ (Edge and Stewart \\cite{edge91}).  Table~\\ref{tbl:imlr} shows the results for representative clusters, where $N_{0.5}$ is available.  The IMLR of MKW~9 is relatively smaller than for other clusters. The potential of MKW 9 might not be as deep as in other systems and a large part of the metal-enriched gas has been driven out already from the potential well.  An IMLR of order $10^{-2}$ in other systems is consistent with previous measurements (e.g.  Tsuru \\cite{tsuru91}) and at least a few times larger than the predicted value based upon the current rate of SN Ia in elliptical galaxies, as argued previously (e.g., Tsuru \\cite{tsuru91}; Renzini et al. \\cite{renzini93}).  We presume that a large fraction of Fe has been produced also in the early phase of the galaxy formation with a high production rate. Relatively low metal abundances observed in elliptical galaxies (sub-solar at most; e.g.  Matsushita et al.  \\cite{matsushita00}) also support this idea.  \\subsection{Spatial distribution of the O, Si, S, and Fe abundances} The most simple model to explain the results (3) and (4) is that Si, S, and Fe have a common origin for a large part, while O has a different origin. This contrast between Si-S-Fe and O has reported separately in A~496 (Tamura et al.  \\cite{tamura01b}). Qualitatively this trend is consistent with the current idea that the contribution from SN Ia relative to that from SN II, $N_{\\rm Ia}/N_{\\rm II}$, increases towards the cluster center and the O/Fe ratio is more sensitive to $N_{\\rm Ia}/N_{\\rm II}$ than the Si/Fe ratio.  We do not attempt a further detailed comparison between the observation and models, because of large uncertainties in both the data (e.g.  the O/Fe ratio in the outer regions and for hot clusters) and the theory (e.g.  Gibson et al.  \\cite{gibson97})."
    },
    "0403/astro-ph0403091_arXiv.txt": {
        "abstract": "RX J1211+2242 is an optically faint ($B \\simeq 19.2 \\, \\rm mag$) but X-ray bright ($f_{2 - 10 \\keV} = 5 \\times 10^{-12} \\ecs$) AGN, which has been shown to be a BL Lac object at redshift $z = 0.455$. The {\\it ROSAT} X-ray, Calar Alto optical, and NVSS radio data suggest that the peak of the synchrotron emission of this object is at energies as high as several keV. {\\it BeppoSAX} observations have been carried out simultaneously with optical observations in order to extend the coverage to higher energies. The new data indeed indicate a turn-over in the $2 - 10 \\keV$ energy region. We propose that RX J1211+2242 is the counterpart of the unidentified EGRET source 3EG J1212+2304, making it a gamma-ray emitter with properties similar to, for example, Markarian 501 in its bright state, though being at a much larger distance. % ",
        "introduction": "The nature of astronomical sources in the gamma-ray domain above 20 MeV is still not well understood. While we have identified most of the sources seen at energies below 1 MeV, the identification of sources at higher energies suffers from large uncertainties in their positions and in many cases strong variability. The Energetic Gamma-Ray Experiment Telescope (EGRET, Kanbach et al. 1988) on board the {\\it Compton Gamma Ray Observatory} ({\\it CGRO}) covered the energy range between 20 MeV to over 30 GeV. But, even ten years after the mission, most of the 271 sources listed in the 3rd EGRET catalog are still unidentified \\cite{hartman}. While many sources in the Galactic plane are associated with populations of massive stars \\citep{kaaret,yadi}, the largest fraction of high Galactic latitude sources turned out to be AGNs. Among this latter group, blazars are dominant.  The common model % of the blazar phenomenon is that we are looking down % a highly relativistic jet when observing these sources. This leads to the characteristic properties of blazars, such as high polarization, high variability, and, for some objects, featureless optical spectra. The overall spectral energy distribution (SED) can be described by two components \\cite{pado}. The first component is thought to be due to synchrotron radiation from relativistic electrons (in some models also positrons) interacting with the magnetic field of the jet.  The peak emission of this component occurs in the infrared region for Flat Spectrum Radio Quasars, in the UV range for low frequency cut-off BL Lacs (LBL), and mostly in the X-ray domain for high frequency cut-off BL Lacs (HBL). The second component can be well modeled by inverse Compton scattering of photons by relativistic electrons. Here the seed photons can be either the synchrotron photons themselves (SSC; e.g. Mastichiadis \\& Kirk 1997) or external photons produced by the accretion disk \\cite{EC} and/or the broad line region \\cite{sikora}. Due to relativistic beaming towards the observer, the emission of blazars is amplified, which makes them the brightest extragalactic objects, especially around the observed peak frequencies. This also explains their dominance in the EGRET energy band, where the peak of the inverse Compton emission is expected.  Here we present the blazar RX J1211+2242 and show that it is likely to be associated with the unidentified EGRET source 3EG J1212+2304. ",
        "conclusions": "The RX J1211+2242 data presented here, together with previous observations\\\\ \\cite{hrxbllac} indicate that this source is a BL Lac object with a high synchrotron peak frequency in the $2 - 10 \\keV$ energy range. The object exhibits strong variations in the X-ray domain on long term time scales (years), though no significant variability was detected during the individual \\sax observations.  RX J1211+2242 is located inside the 95\\% confidence radius of 3EG J1212+2304, which has been shown to be a highly variable source (see Fig.~\\ref{fig:EGRETlc}). Galactic sources, such as pulsars and supernova remnants show little variability in the EGRET observations, while AGN exhibit strong flux variations \\cite{variabilityEGRET}. This is a first indication that 3EG J1212+2304 might be an AGN. At a Galactic latitude of $b = 80.02^\\circ$, this source is far out of the Galactic plane, making an extragalactic origin more likely, and thus supporting the blazar identification. RX J1211+2242 is the only plausible counterpart to this source inside the error circle. Sowards-Emmerd et al. (2003) claim that 3EG J1212+2304 is not a blazar but base this on the evidence that there is no radio flux enhancement towards the center of the (large) EGRET error circle and that there was no known blazar inside the error circle, which we have just shown is not the case. %  Fitting the spectral energy distribution to a Synchrotron Self Compton model shows that physically meaningful parameters can explain the overall data, including the EGRET measurements. Although BL Lacs are known to exhibit a strong variability in the inverse Compton branch, it is unlikely that the {\\it BeppoSAX} observations detect the source in an outburst phase due to an expected duty cycle $\\ll 0.5$. On the other hand, most of the measurements of 3EG J1212+2304 are upper limits and the modeled IC branch falls well within the region of possible flux variations. Furthermore, no other X-ray source inside the error circle is likely to be the counterpart, and no strong radio source has been found there either. Since all the identified EGRET blazars have also been detected in the RASS, a connection between a strong {\\it ROSAT} source and an EGRET detection is even more likely.  Assuming that the BL Lac and the gamma-ray emitter are in fact the same source, this yields a SED similar to that of Markarian 501 (z=0.034) in its bright state. An emitting region of $R = 5 \\times 10^{15} \\rm \\, cm$, a magnetic field of $B = 0.1 \\rm \\, G$, Doppler factor of $\\delta = 20$ and a Lorentz factor for electrons at the break $\\gamma_b = 3 \\times 10^4$ reproduce the SED of RX J1211+2242 under the assumption that the flux during the {\\it BeppoSAX} and optical observation was below the EGRET upper-limit of 3EG J1212+2304. The SED would reach the EGRET detections by using $\\delta = 10$ and $\\gamma_b = 6 \\times 10^4$. In order to match the different detections and non-detections, the IC branch flux needs to vary by a factor of $\\sim 50$ at the peak of the emission. Again, this behavior is similar to that observed in Markarian 501 that showed variability in the hard X-rays (10-100 keV) by a factor of $\\sim 100$ in different {\\it BeppoSAX} observations coupled with a simultaneous shift of the synchrotron peak from the soft X-rays up to the soft gamma-rays. It also showed a strong variability, by a factor of $\\sim 100$, in the EGRET band \\cite{mrk501var}.  The luminosity of the inverse Compton branch seems to be much larger than the synchrotron luminosity, at least during those periods in which the source was detected by EGRET. If the association of the EGRET source with the BL Lac is correct, this would make RX J1211+2242 a peculiar HBL, which usually show $L_c/L_s \\sim 1$ (e.g. Fossati et al. 1998). A value of $L_c/L_s \\gg 1$ has already been seen in some HBLs (e.g. bright state of Markarian 501 and Markarian 421), and 1426+428 has shown an even larger dominance of the IC branch ($f_{IC} > 10^{-10} \\ecs$; Aharonian et al. 2003) over the synchrotron component ($f_s \\simeq 2 \\times 10^{-11}  \\ecs$; Costamante et al. 2001). It also must be taken into account that the bright gamma-ray state of this object is likely to be accompanied by increased synchrotron emission, which would decrease the $L_c/L_s$ ratio to values previously observed in HBL. With respect to the blazar sequence \\cite{fossati}, RX J1211+2242 is an HBL with remarkably strong gamma-ray emission, making it an interesting object for further studies at energies above $\\sim 1 \\rm \\, MeV$.  The interesting region where the inverse Compton branch becomes dominant over the synchrotron branch lies in the energy range covered by the {\\it INTEGRAL} instruments ($20 - 8000 \\keV$). Given the X-ray spectrum shown by RX J1211+2242 in the {\\it BeppoSAX} data, it will not be easily detectable by this mission, even when assuming that the power law with $\\alpha_X \\simeq 1.0$ extends up to several hundred keV. We have performed simulations that show that a 500 ksec observation would give a $< 5 \\sigma$ detection by the SPI spectrometer. However, the upcoming gamma-ray missions {\\it GLAST} \\cite{GLAST} and {\\it AGILE} \\cite{AGILE}, with their improved sensitivity, should be able to detect RX J1211+2242 in the EGRET band.  Even though this source is likely to be a strong TeV emitter, it will not be seen by most Cherenkov telescopes as the high energy emission of an object at redshift $z = 0.455$ will be suppressed by interaction of the TeV photons with the extragalactic infrared background \\cite{kneiske}. Nevertheless the {\\it MAGIC} telescope \\cite{MAGIC} with its low energy threshold of $\\sim 30 \\, \\rm GeV$ is suitable to observe the peak of the IC component making RX J1211+2242 one of the most promising targets in the gamma-ray sky."
    },
    "0403/astro-ph0403572_arXiv.txt": {
        "abstract": "We derive absolute dimensions of the early B-type detached eclipsing binary V453\\,Cygni (B0.4\\,IV + B0.7\\,IV, $P=3.89$\\,d), a member of the open cluster NGC\\,6871. From the analysis of new, high-resolution, spectroscopy and the $UBV$ light curves of Cohen (1974) we find the masses to be $14.36 \\pm 0.20$\\Msun\\ and $11.11 \\pm 0.13$\\Msun, the radii to be $8.55 \\pm 0.06$\\Rsun\\ and $5.49 \\pm 0.06$\\Rsun, and the effective temperatures to be $26\\,600 \\pm 500$\\,K and $25\\,500 \\pm 800$\\,K for the primary and secondary stars, respectively. The surface gravity values of $\\logg = 3.731 \\pm 0.012$ and $4.005 \\pm 0.015$ indicate that V453\\,Cyg is reaching the end of its main sequence lifetime. We have determined the apsidal motion period of the system to be $66.4 \\pm 1.8$\\,yr using the technique of Lacy (1992) extended to include spectroscopic data as well as times of minimum light, giving a density concentration coefficient of $\\log k_2 = -2.226 \\pm 0.024$. Contaminating (third) light has been detected for the first time in the light curve of V453\\,Cyg; previous analyses without this effect systematically underestimate the ratio of the radii of the two stars. The absolute dimensions of the system have been compared to the stellar evolution models of the Granada, Geneva, Padova and Cambridge groups. All model sets fit the data on V453\\,Cyg for solar helium and metal abundances and an age of $10.0 \\pm 0.2$\\,Myr. The Granada models also agree fully with the observed $\\log k_2$ once general relativistic effects have been accounted for. The Cambridge models with convective core overshooting fit V453\\,Cyg better than those without. Given this success of the theoretical predictions, we briefly discuss which eclipsing binaries should be studied in order to further challenge the models. ",
        "introduction": "\\label{introduction}  Theoretical models of high-mass stars are very difficult to construct, due to a number of poorly understood astrophysical phenomena which become important at higher stellar masses. Convective core overshooting is an essential ingredient of stellar models and can have a large effect on the lifetimes and luminosities of high-mass stars. The degree of overshooting may depend on metallicity (Cordier \\etal\\ 2002, Palmieri \\etal\\ 2002) and stellar mass (Young \\etal\\ 2001). Stellar rotation is known to cause enhanced mixing in the interiors of stars which can mimic a small amount of convective overshooting and strongly affect the evolution of stars (Maeder \\& Meynet 2000). Mass loss is also an important effect (Maeder 1997) but is usually represented in theoretical models by a parameterisation such as that of Reimers (1975) (Woo \\etal\\ 2003).  Until recently there was a discrepancy between masses derived from spectroscopic observations and inferred from evolutionary models (Herrero \\etal\\ 1992). Recent advances in theoretical modelling, and the increased sophistication of observational techniques, appears to have removed this mass discrepancy (Hilditch, 2004), but confirmation of this requires detailed comparisons between models and as many physical properties of individual stars as possible.   Detached eclipsing binaries (dEBs) are a vital source of observational data on the radii and masses of high-mass stars (Andersen 1991) but studies of such stars are hindered by the rarity of useful systems, the difficulty of measuring accurate velocities from spectra of rapidly-rotating multiple stars (Sana, Rauw \\& Gosset 2001), and the large amount of telescope time needed to obtain light curves which are complete over the whole orbital period. This has caused a shortage of data on eclipsing binaries with masses above 10\\Msun\\ and sufficiently large orbital periods to be detached and therefore to have evolved as single stars.  Whilst dEBs can provide accurate observed stellar masses, radii, effective temperatures and equatorial rotational velocities, the ages, chemical compositions and distances of such systems are not, in general, directly observable. This means that when fitting observed quantities to theoretical predictions, there are several free parameters which can be adjusted to find the best fit to the observational data. This makes it very difficult to assess the success of important but more subtle physical effects such as rotational mixing, mass loss, convective overshooting and even the type of turbulent convective theory.  Eclipsing binaries in open clusters and associations can be analysed to derive accurate masses, radii, temperatures and rotational velocities. Their membership of a star cluster can also provide knowledge of the age, chemical composition, and distance of the dEB, allowing a more complete physical description of a single stellar system. This reduces the number of free parameters adjustable when fitting theoretical models to observations, which allows a discriminate test of the representation of different physical effects and indeed the success of different sets of stellar evolutionary models. For example, from the analysis of two main sequence dEBs which are members of the open cluster h\\,Persei (NGC\\,869), Southworth, Maxted \\& Smalley (2004, hereafter Paper\\,I) were able to provide a first estimate of the bulk metallicity of one of the most important and well-studied open clusters in the Northern Hemisphere.   \\subsection{V453\\,Cygni in NGC\\,6871}     \\label{v453cyg}  \\begin{table} \\begin{center} \\caption{\\label{tablephotdata} Identifications, location, and combined photometric parameters for the V453\\,Cygni eclipsing system. \\newline {\\bf References:} (1) Cannon \\& Pickering (1923); (2) Argelander (1903); (3) H{\\o}g \\etal\\ (1998); (4) Hoag \\etal (1961); (5) Popper (1980); (6) Zakirov (1992); (7) Cohen (1969); (8) Reimann (1989).} \\begin{tabular}{lr@{}lr} \\hline \\hline &     & V453 Cygni            & Reference \\\\ \\hline Henry Draper number     &     & HD 227696             & 1   \\\\ Bonner Durchmusterung   &     & BD\\,+35\\degr 3964     & 2   \\\\ Hoag number             &     & NGC 6871 13           & 3   \\\\ \\hline $\\alpha_{2000}$         &     & 20 06 34.967          & 4   \\\\ $\\delta_{2000}$         &   + & 35 44 26.28           & 4   \\\\ Spectral type           &   & B\\,0.4\\,IV + B\\,0.7\\,IV & 5   \\\\ \\hline $V$                     &     & 8.285                 & 6   \\\\ $B-V$                   &   + & 0.179                 & 6   \\\\ $U-B$                   & $-$ & 0.61                  & 6   \\\\ $V-R$                   &   + & 0.254                 & 6   \\\\ $\\beta$                 &     & 2.590                 & 7,8 \\\\ \\hline \\hline \\end{tabular} \\end{center} \\end{table}  \\begin{table*} \\begin{center} \\caption{\\label{pubspecorbits} Published spectroscopic orbits of V453\\,Cygni. BMM97 originally fitted their data with a circular orbit. We have refitted their radial velocities with an eccentric orbit to increase the accuracy of our determination of the apsidal motion (see section~\\ref{perioddet} for details). A colon after a number indicates that it is uncertain. Quantities without quoted errors or a colon were not determined by that investigation. When quantities are given separately for each star we have quoted a weighted mean of the two values. Symbols have their usual meanings. Times are written as (HJD $-$ 2\\,400\\,000). \\newline $^*$\\,The reference time, $T_0$, refers to a time of periastron passage, not a time of minimum light.} \\begin{tabular}{l r@{\\,$\\pm$\\,}l r@{\\,$\\pm$\\,}l r@{\\,$\\pm$\\,}l r@{\\,$\\pm$\\,}l r@{\\,$\\pm$\\,}l r@{\\,$\\pm$\\,}l } \\hline \\hline & \\mc{Pearce}  & \\mc{Abt, Levy and} & \\mc{Popper and} &\\mc{Simon and}   & \\mc{BMM97}      & \\mc{BMM97}          \\\\ & \\mc{(1941)}  & \\mc{Gandet (1972)} & \\mc{Hill (1991)}&\\mc{Sturm (1994)}& \\mc{}           & \\mc{(our solution)} \\\\ \\hline $P$ (days) & \\mc{3.87972} & \\mc{3.8890}        & \\mc{3.8898128}  & \\mc{3.88982309} &\\mc{3.8898128}   &\\mc{3.889825}        \\\\ $T_0$ (HJD) &30231.0843&0.0543$^*$ & \\mc{40495.027$^*$} & \\mc{39340.099} & \\mc{36811.7296} &48141.82&0.01$^*$& 48500.64 & 0.66$^*$ \\\\ $K_{\\rm A}$ (\\kms) & 181.8 & 1.13 & \\mc{152:}          & 171 & 1.5       & 171.7 & 2.9     & 173.2 & 1.3     & 173.7 & 1.4         \\\\ $K_{\\rm B}$ (\\kms) & 237.4 & 2.78 & \\mc{}              & 222 & 2.5       & 223.1 & 2.9     & 213.6 & 3.0     & 212.4 & 3.4         \\\\ $e$ & 0.07 & 0.007 & \\mc{0.05:}         &\\mc{0.0}         &\\mc{0.0}         & \\mc{0.0}        & 0.011 & 0.015       \\\\ $\\omega$ (degrees) & 175.2 & 5.06 & \\mc{99:}           & \\mc{}           & \\mc{}           & \\mc{}           & 88.6 & 6.0          \\\\ \\Vsys\\ (\\kms) & $-$15.0&0.94 & \\mc{$-$22.7:}      &\\mc{$-$14}       & \\mc{$-$7:}      & $-17.6$ & 1.0   & $-$18.0 & 1.6       \\\\ \\hline \\hline \\end{tabular} \\end{center} \\end{table*}  V453\\,Cygni is a high-mass dEB with an orbital period of 3.89\\,days. Its membership of the young open cluster NGC\\,6871 means that its age and distance can be found independently. The primary component of V453\\,Cyg is approaching the terminal age main sequence (TAMS) and its large radius causes the eclipses to be total, allowing a very accurate determination of the radii of both stars. Table~\\ref{tablephotdata} contains identifications and some photometric properties of the system.  The eclipsing nature of V453\\,Cyg was discovered by Wachmann (1939) and an early spectroscopic orbit was calculated by Pearce (1941). A period study by Cohen (1971) provided a determination of the orbital longitude of periastron, $\\omega$, inconsistent with that derived by Pearce. In a period study by Wachmann (1973) this was correctly interpreted as apsidal motion. Wachmann derived an apsidal period of $U = 72$ years using measurements of the time differences between several groups of adjacent primary and secondary eclipses. A more recent period study, including parabolic and periodic terms, was undertaken by Rafert (1982).  Excellent photoelectric $UBV$ light curves were observed by Wachmann (1974) and analysed using the Russell-Merrill method (Russell \\& Merrill 1952) which involves the procedure of rectification. His work contains a plot of the light curves adjusted for the effects of orbital eccentricity and apsidal motion using parameters updated from that of his previous work, but the data themselves have so far been unobtainable. It is possible that they are in an unlabelled file in the IAU Variable Star Archives (Breger 1988), but no record of them exists at Hamburg Observatory, where the light curves were observed (A.\\ Reiners, private communication).  Cohen (1974) published complete photoelectric $UBV$ light curves which contain fewer datapoints and more observational scatter than those of Wachmann (1974). He analysed these using the Russell-Merrill method but stated that his observations were not definitive. They have since been analysed by Cester \\etal\\ (1978) using the light curve analysis code {\\sc wink} (Wood 1971). This is the only previous photometric study to use modern techniques.  A recent investigation using photoelectric $UBVRI$ light curves has been published by Zakirov (1992). He analysed his light curves using the ``direct machine method of Lavrov (1993)'', which is based on rectification. The results of the four photometric analyses of V453\\,Cyg are substantially in agreement about the basic photometric parameters of the system.  Recent spectroscopic orbits have been published by Popper \\& Hill (1991), Simon \\& Sturm (1994) and Burkholder, Massey \\& Morrell (1997, hereafter BMM97). These results are collected in Table~\\ref{pubspecorbits}. Simon \\& Sturm used seven spectra to demonstrate their spectral disentangling algorithm, which decomposes observed composite spectra into the separate spectra of two stars. The total secondary eclipse of V453\\,Cyg allowed them to directly compare their disentangled primary spectrum with a spectrum observed during secondary eclipse. Disentangling can be used to determine accurate orbital semiamplitudes (Hynes \\& Maxted 1998, Harries \\etal\\ 2003) but it is not clear if there is a robust method by which to estimate uncertainties in the derived quantities. BMM97 derived a good spectroscopic orbit from 25 high signal-to-noise spectra and compared the dEB to models to investigate the discrepancy at higher masses between models and observations. The rotational velocities of the components of V453\\,Cyg were determined by Olson (1984) to be $107 \\pm 9$\\kms\\ and $97 \\pm 20$\\kms\\ for the primary and secondary stars respectively. A preliminary single-lined spectroscopic orbit was also given by Abt, Levy \\& Gandet (1972).  An abundance analysis of V453\\,Cyg was undertaken by Daflon \\etal\\ (2001) using both LTE and non-LTE calculations. The results suggest that V453\\,Cyg has a slightly sub-solar metallicity. These authors derived an effective temperature of 29\\,100\\,K using the $Q$ parameter based on $UBV$ magnitudes (Johnson 1957), and a surface gravity of $\\logg = 4.45$ (\\cms) from profile fitting of the H$\\gamma$ 4340\\,\\AA\\ spectral line. Both values are larger than expected and inconsistent with previous analyses. Their surface gravity value is in fact somewhat larger than appropriate for the ZAMS, and is inconsistent with significant main sequence evolution.    \\subsection{NGC\\,6871}              \\label{ngc6871}  The open cluster NGC\\,6871 is a concentration of bright OB stars which forms the nucleus of the Cyg\\,OB3 association (Garmany \\& Stencel 1992). This makes it an important object for the study of the evolution of high-mass stars. The cluster itself has been studied photometrically several times but its sparse nature means determination of its physical parameters is difficult.  $UBV$ photometry of the 30 brightest stars was published by Hoag \\etal\\ (1961). Crawford, Barnes \\& Warren (1974) observed 11 stars in the Str\\\"omgren $uvby$ system and 24 stars in the Crawford $\\beta$ system, finding significantly variable reddening and a distance modulus of 11.5\\,mag. This $uvby\\beta$ photometry was extended to 40 stars by Reimann (1989), who found reddening $\\Eby$ with a mean value of 0.348\\,mag and an intracluster variation of about 0.1\\,mag. His derived distance modulus of $11.94 \\pm 0.08$\\,mag and age of 12\\,Myr are both greater than previous literature values.  Massey, Johnson \\& DeGioia-Eastwood (1995) conducted extensive $UBV$ CCD photometry of 1955 stars in the area of Cyg\\,OB3. Their values of distance modulus, $11.65 \\pm 0.07$, and of reddening, $\\EBV = 0.46 \\pm 0.03$\\,mag with individual values between 0.04 and 1.11 mag, agree well with previous determinations. They find an age of 2 to 5 Myr for stars with spectral types earlier than B0 but give evidence for a significant spread of stellar ages in the cluster. Whilst the highest-mass unevolved cluster members have main sequence lifetimes of 4 to 5 Myr, NGC\\,6871 contains evolved 15\\Msun\\ stars despite their main sequence lifetimes being of the order of 11\\,Myr. ",
        "conclusions": "\\label{discussion}  We have derived the absolute dimensions of the components of the high-mass detached eclipsing binary V453\\,Cygni, a member of the open cluster NGC\\,6871.  Effective temperatures were found using the helium ionisation balance derived from high-resolution spectra, which also suggest an enhanced photospheric helium abundance relative to solar. Radial velocities were derived from the spectra using only the weak spectral lines and the {\\sc todcor} two-dimensional cross-correlation algorithm.  The apsidal motion rate of the system has been determined using an extended version of the photometric method of Lacy (1992), which includes times of minimum light and spectroscopic determinations of eccentricity and $\\omega$. The apsidal period is well constrained, and allow the derivation of eccentricity and $\\omega$ to a greater accuracy than possible with the light curves and radial velocity curves.  We have reanalysed the $UBV$ light curves of Cohen (1974) in order to determine the radii of the components of the dEB. The best-fitting parameters include a small amount of third light, which was previously undetected. Robust parameter uncertainties were derived using bootstrapping, allowing us to quantify and illustrate the effect of correlations between different photometric parameters. The ratio of the radii and the amount of third light are strongly correlated, due to their dependence on the depths of the eclipses; previous photometric studies which did not include third light are systematically biased towards values of the stellar radii which are 1\\% higher and 5\\% lower for primary and secondary respectively.  The accurate absolute dimensions presented here allow V453\\,Cyg to be added to the list of dEBs with the best-determined values of mass, radius and effective temperature (Andersen 1991). However, our analysis would clearly be much improved with better observational data. The inclusion of only a few new times of minima would greatly increase the accuracy of the results of the apsidal motion analysis, and more accurate rotational velocities would allow a more accurate derivation of the internal structure constant $\\log k_2$. A definitive spectroscopic orbit will require observations with a higher signal-to-noise than those presented here, and should give masses determined to accuracies of better than 1\\%. Definitive light curves of the system would allow determination of the limb darkening coefficients for both stars, providing an important test of model atmosphere codes.  The absolute masses, radii and effective temperatures of the components of V453\\,Cyg have been compared to several stellar models in the mass--radius and $\\log\\Teff$--\\logg\\ planes, assuming the same age for both stars. Not only is there impressive agreement betwen different theoretical models, all model sets are able to fit the observational data for a solar helium and metal abundance. Moreover, the Granada models (Claret 1995) provide a perfect match to the observed apsidal motion rate once the relativistic contribution has been subtracted from the overall effect. Stellar models have for a long time appeared to predict that the central condensations of stars are lower than that found using observations of apsidal motion (e.g., Clausen, Gim\\'enez \\& Scarfe 1986, Barembaum \\& Etzel 1995). This apparent discrepancy has been reduced by the discovery that the internal structure constants change significantly through a star's evolution. The current generation of theoretical models, incorporating OPAL opacity data (Rogers \\& Iglesias 1992), are in good agreement with observations. It is noticeable that some observers have not removed the general relativistic effect from their observed $\\log k_2$ values before comparison with theory; in many cases this will have a negligible effect but for the stars of V453\\,Cyg it causes about 6\\% of the observed apsidal motion, changing $\\log k_2$ by an amount similar to its uncertainty.  The normal helium abundance implied by stellar model fits also conflicts with the slight overabundance noted in our spectral synthesis analysis. We note that the photospheric helium abundance is not directly comparable to the initial internal helium abundance used in model calculations.  Fits to the Cambridge stellar models support the inclusion of a moderate amount of overshooting in most stellar evolutionary models. Whilst models without overshooting were able to fit the masses and radii of the stars, the predicted effective temperatures are slightly lower than that determined from the helium ionisation balance.  The stellar models were extremely successful in fitting the absolute dimensions and effective temperatures of a high-mass slightly-evolved dEB, with component masses and radii differing by ten and twenty-five times their combined uncertainties, respectively. For observational stellar astrophysicists, this fact implies that we must either observe systems so thoroughly that their masses and radii are known to accuracies of 0.5\\% and the effective temperatures to 2\\%, or target particular types of stars to critique the success of one set of stellar models compared to another. Such targets include low-mass, high-mass, pulsating, and Population II stars, as well as eclipsing systems found in Local Group galaxies. Eclipsing binaries in open clusters can satisfy this requirement if the cluster they belong to is well-studied or otherwise interesting (Paper\\,I), and further observations are being undertaken towards this goal."
    },
    "0403/astro-ph0403434_arXiv.txt": {
        "abstract": "{In the proceedings of this, and of several recent close binary conferences, there have been several contributions describing smoothed particle hydrodynamics simulations of accretion disks. It is apposite therefore to review the numerical scheme itself with emphasis on its advantages for disk modelling, and the methods used for modelling viscous processes. } \\addkeyword{accretion, accretion disks} \\addkeyword{binaries:close} \\addkeyword{line:profiles} \\addkeyword{x-rays:binaries} \\addkeyword{hydrodynamics} \\addkeyword{methods:numerical}  \\begin{document} ",
        "introduction": " ",
        "conclusions": "One of the great advantages of the smoothed particle hydrodynamics technique is its simplicity. An SPH code is rarely more than a thousand lines long. This, in combination with established methods for incorporating new physics into the algorithm, makes it an ideal tool for studying systems where the important physics is uncertain. In the case of accretion disk research, we can not yet describe with any confidence the form of the anomalously large dissipation. We can however make use of SPH simulations to investigate various postulated forms for disk viscosity (e.g. the viscoelastic description developed by Ogilvie, 2003). The encouraging message is thus that if there is viscosity in an SPH simulation of a disk, it is there specifically to model the unknown process driving accretion in close binaries. It is without doubt possible to have shocks and a prescribed viscous dissipation term in the same simulation. We conclude by referring the reader to presentations of SPH disk simulations at this conference by Hayasaki, Kunze, Manson, Truss and Okazaki."
    },
    "0403/astro-ph0403328_arXiv.txt": {
        "abstract": "A point mass at the center of an ellipsoidal homogeneous fluid is used as a simple model to study the effect of rotation on the shape and external gravitational field of planets and stars.  Maclaurin's analytical result for a homogenous body is generalized to this model. The absence of a third order term in the Taylor expansion of the Maclaurin function leads to further simple but very accurate analytical results connecting the three observables: oblateness ($\\epsilon$), gravitational quadrupole ($J_{2}$), and angular velocity parameter ($q$).  These are compared to observational data for the planets.  The moments of inertia of the planets are calculated and compared to published values.  The oblateness of the Sun is estimated. Oscillations near equilibrium are studied within the model.  \\\\ \\\\ ",
        "introduction": "The rotation induced oblateness of astronomical bodies is a classical problem in Newtonian and celestial mechanics (for the early history, see Todhunter \\cite{todhunter}).  It has twice played an important role in the history of science.  In the early eighteenth century measurements indicated a prolate shape of the Earth, in strong conflict with the Newtonian prediction.  This was later shown to be wrong by more careful measurements by Maupertuis, Clairaut, and Celsius in northern Sweden in 1736.  Then, in 1967, measurements of the solar oblateness were published and according to these it was much larger than the Sun's surface angular velocity would explain.  The confirmation of general relativity by Mercury's perihelion precession would then be lost.  Also this problem is now gone and the modern consensus is that the solar oblateness is too small to affect this classic test of general relativity \\cite{will,godier,laskar}.   The subject of the flattening of rotating astronomical bodies is thus quite mature.  The classical theory is due mainly to Clairaut, Laplace, and Lyapunov. Also Radau, Darwin, de Sitter and many others have made important contributions. More recent accounts of the theory can be found in, for example, Jeffreys \\cite{jeffreys}, Zharkov et al.\\ \\cite{zharkov}, Cook \\cite{cook}, Moritz \\cite{moritz} and, partly, in Chandrasekhar \\cite{chandrasekhar}. Some pedagogical efforts can be found in Murray and Dermott \\cite{murray}, or in Kaula \\cite{kaula}.  As is plain from these references the theory is quite involved.  Only the unrealistic assumption that the body is homogeneous gives compact analytical results.  Otherwise a specified radial density distribution is needed and one resorts to cumbersome series (multipole) expansions, or purely numerical methods, for quantitative results.  Here we will present analytical results based on the assumption that the body consists of a central point mass surrounded by a homogeneous fluid, the so called point core model.  By varying the relative mass of the fluid and the central point particle one can interpolate between the extreme limits of Newton's homogeneous body and the Roche model \\cite{kippenhahn} with a dominating small heavy center.  The point core model goes back to the work of G.~H.\\ Darwin.  More recently it has been used to study the shape of outer planet moons, see Hubbard and Anderson \\cite{hubbard}, Dermott and Thomas \\cite{dermott}.   Apart from the basic point core approximation (i) several further approximations are assumed here.  These are: (ii) that the shape is determined by hydrostatic equilibrium and (iii) that the shape is ellipsoidal.  These are not consistent. According to Hamy's theorem, see Moritz \\cite{moritz}, the exact shape is not ellipsoidal, so we regard an ellipsoidal shape as a constraint and find the equilibrium shape, among these, by minimizing the energy.  A further approximation (iv) is the neglect of differential rotation.  On the other hand fixed volume (``bulk incompressibility'') need not be assumed; the equilibrium volume problem separates from the shape problem.  In spite of these approximations, which are standard in the literature, the mathematics can be quite involved.  In this article I hope to clarify and simplify it as much as possible.    Mathematically our model then becomes a three degree of freedom mechanical system for which we can calculate the kinetic and potential (gravitational plus centrifugal) energies exactly.  Multipole expansions in terms of spherical harmonics are not needed.  The three degrees of freedom correspond to the three semi-axes of the ellipsoid ($a, b, c$) but these are transformed to three generalized coordinates that describe size (or volume), $R$, spheroidal, $\\xi$, and triaxial, $\\tau$, shape changes, see Eq.\\ (\\ref{eq.def.xi.eta}).  The statics problem of equilibrium shape is solved by minimizing the potential energy, $U(R,\\xi,\\tau)$ for a fixed $R$, Eq.\\ (\\ref{eq.energy.xi.tau}).  For slow rotation the shape will not be triaxial so $\\tau=0$.  Finding the shape, or flattening, is then only a matter of minimizing a dimensionless potential energy, \\begin{equation} \\label{eq.dim.less.pot.en.xi} u(\\xi)=\\psi(\\xi)-\\frac{1}{2}k \\xi^2 , \\end{equation} given by the last two terms of Eq.\\ (\\ref{eq.energy.xi}).  Here $\\psi(\\xi)$ is defined by Eq.\\ (\\ref{eq.basic.phi.xi}) and plotted in Fig.\\ \\ref{FIG1} while $k$ is a constant that depends on rotational parameter $q$ and dimensionless moment of inertia $\\kappa^2$, Eq.\\ (\\ref{eq.kappa2.nu}).  The root $\\xi(k)$ of the equation $\\psi'(\\xi) -k\\xi =0$ thus gives the flattening.  This root is to very high accuracy given by \\begin{equation} \\label{eq.xi.root.of.k} \\xi(k)=\\left[1-(7/2)\\left(\\sqrt{1+15 k/7} - 1\\right) \\right]^{-1/6} , \\end{equation} see Eqs.\\ (\\ref{eq.maclaurin.approx.inv}) and (\\ref{eq.eccentricity.relation.xi}), for the rotational parameters that can be found in the solar system.  From this result, simple formulas, Eqs.\\ (\\ref{eq.maclaurin.approx.relation.q}) -- (\\ref{eq.maclaurin.approx.relation.epsilon}), relating the observables, rotation parameter, $q$, gravitational quad\\-ru\\-pole, $J_{2}$, and excentricity squared, $e^2$, are obtained.  These appear to be, partly, new, and their usefulness is demonstrated by comparing with empirical data for the Sun and the rotating planets of the solar system.  Variational methods have been used before to study similar problems, see for example Abad et al.\\ \\cite{abad}. Denis et al.\\ \\cite{denis1} have pointed out that variational methods generally fail to provide estimates of their accuracy. Therefore the agreement of our formulas with empirical data, as demonstrated in Table /cite(table), is important and demonstrates that our model catches the essential physics of rotational flattening.  Finally small amplitude oscillations near the equilibrium are investigated, starting from the Lagrangian \\begin{equation} \\label{eq.lagrangeian} L=T(R,\\xi,\\tau,\\dot R, \\dot\\xi,\\dot\\tau)-U(R,\\xi,\\tau), \\end{equation} for our three degree of freedom model system.  This gives useful insight into the physics of free stellar or planetary oscillations and their coupling to rotation.  Our approximations are, however, too severe for these results to be of quantitative interest. ",
        "conclusions": "Some apparently new results relating to the classic theory of the figure of rotating bodies have been presented.  The basic model, a point mass at the center of a homogeneous fluid, is characterized by their mass ratio, and interpolates between the limits of an ellipsoidal homogeneous fluid and a body dominated by a small central mass concentration.  It allows simple analytic treatment but is still flexible enough to correctly describe the essential hydrostatics of real rotating planets as well as stars.  Such models are always useful, especially when one wishes to analyze, compare, and understand large numbers of observational data.  In spite of the simplicity there is no perturbation order to which the results are valid; the nonlinearity of the basic equations can be retained.  To be more precise Eq.\\ (\\ref{eq.maclaurin.approx}) shows that the formulas are valid to seventh order in the eccentricity (within the basic model with its simplified mass distribution).  As demonstrated by the numerical experiments on Jupiter and Saturn data, this is essential. In fact Table~\\ref{table} indicates that the geometric oblateness of the surface equipotential surface of Mars and Uranus determined from observed $q$ and $J_{2}$ values using (\\ref{eq.maclaurin.approx.relation.epsilon}) probably are more reliable than current observational data.  Finally the dynamics of the model reveals the essentials of the coupling and rotational splitting of the most basic free oscillations of a planet without recourse to expansion in spherical harmonics.\\\\ \\\\ \\noindent {\\bf Acknowledgement} Constructive comments from Dr.\\ John D.\\ Anderson on a previous version of this manuscript are gratefully acknowledged."
    },
    "0403/astro-ph0403487_arXiv.txt": {
        "abstract": "The recent detection of variable infrared emission from Sagittarius A*, combined with its previously observed flare activity in X-rays, provides compelling evidence that at least a portion of this object's emission is produced by nonthermal electrons. We show here that acceleration of electrons by plasma wave turbulence in hot gases near the black hole's event horizon can account both for Sagittarius A*'s mm and shorter wavelengths emission in the quiescent state, and for the infrared and X-ray flares, induced either via an enhancement of the mass accretion rate onto the black hole or by a reorganization of the magnetic field coupled to the accretion gas. The acceleration model proposed here produces distinct flare spectra that may be compared with future coordinated multi-wavelength observations. We further suggest that the diffusion of high energy electrons away from the acceleration site toward larger radii might be able to account for the observed characteristics of Sagittarius A*'s emission at cm and longer wavelengths. ",
        "introduction": "The observation of stellar motions within light-days of Sagittarius A*, a compact radio source at the Galactic Center \\citep{Balick74}, has provided compelling evidence that this source is the radiative manifestation of a $\\sim$ four million solar mass black hole \\citep{Schodel02, Ghez031}. The recently detected infrared emission and flare activity have provided an additional evidence that this source is powered by a hot gas accreting onto the black hole \\citep{Baganoff01, Goldwurm03, Baganoff032, Porquet03, Zhao04, Ghez03}. The quasi-periodic near-infrared variability may be an indication that the gas flared up before spiraling into the black hole \\citep{Genzel03}. It is now generally agreed that the radio and infrared emission, and the flares, are likely produced by nonthermal high energy electrons (Liu \\& Melia 2001; Genzel et al. 2003;  see also Mahadevan 1998).  However, the exact nature of the mechanism responsible for the acceleration of the electrons has not been addressed. This has given rise to diverse interpretations of the observations with assumed spectra of the accelerated electrons \\citep{Markoff01, Liu02a, Nayakshin03, Yuan03}.  In this letter, we show that the mechanism producing high-energy particles in solar flares works equally well in hot plasmas near the black hole. The solar flare model is based on a second order Fermi acceleration process or a stochastic acceleration (SA) of particles by interacting resonantly with plasma waves or turbulence (PWT) generated via an MHD dissipation process (see e.g. Miller \\& Ramaty 1987; Hamilton \\& Petrosian 1992; Petrosian \\& Liu 2004). In Sagittarius A*, nonthermal particles can be produced by the turbulence expected to be induced by the magneto-rotational instability in the accretion torus (Balbus \\& Hawley 1991; Melia, Liu \\& Coker 2001). The radiation emitted at the acceleration site can explain the quiescent state mm and shorter wavelength observations. Solar flares are energized by the process of magnetic reconnection during the dynamical evolution of the coronal magnetic field. Similar processes in Sagittarius A* can release energy in a small region and produce what we call a local event. A global flare can be induced by an MHD fluctuation in the accretion torus or an enhancement of the accretion rate onto the black hole. The emission spectra from these two energization mechanisms are quite different, and can explain the distinct flare behaviors.  In \\S\\ \\ref{SA}, we outline the theory of SA. Its application to Sagittarius A* is presented in \\S\\ \\ref{sgra}. The main results of this letter are summarized in \\S\\ \\ref{dis}, where we also discuss consequences of energetic electrons escaping the acceleration site. If such electrons can diffuse toward larger radii, they may account for Sagittarius A*'s emission at cm and longer wavelengths and has the potential to explain the observed linear and circular polarization characteristics, and variabilities of this source. ",
        "conclusions": "\\label{dis}  Based on the theory of stochastic particle acceleration, we have built a model to account for Sagittarius A*'s emission at mm and shorter wavelengths. The quiescent state emission is attributed to electrons accelerated by turbulence in a magnetized accretion torus and flares can be produced via two distinct mechanisms. The IR and X-ray flares with harder spectrum are likely induced by a local MHD process, such as magnetic reconnection. Global fluctuations generally produce flares with IR and X-ray spectral indexes close to their corresponding value in the quiescent state. The model not only accounts for the varied spectra of the flares, but also explains their relative occurrence rate observed at IR and X-ray.  Radio emission at longer wavelengths cannot be produced within such a small emission region \\citep{Liu01a}. This is not surprising, given that the radio emission is variable on a longer time scale of tens of hours to one week, suggesting that this radiation is produced at larger radii than what we have been considering here \\citep{Zhao04, Zhao03, Zhao93, Bower02}. Very interestingly, our model suggests a significant outflow of high-energy electrons \\citep{Liu02b}, which may very well be the particles that eventually produce Sagittarius A*'s cm spectrum via synchrotron emission on a spatially larger scale. In the quiescent state (Model A), the power carried away by electrons with $\\gamma>100$ (electrons with lower energy may be trapped by the gravitational potential of the black hole and will not diffuse to larger radii) is about $2\\times10^{37}$ ergs s$^{-1}$, which is more than enough to power the observed radio emission, whose luminosity is about $10^{34-35}$ ergs s$^{-1}$. (Note that in this model, the mass accretion rate is $\\sim 10^{18-19}$g s$^{-1}$, for which a few percent of the dissipated gravitational energy is carried away by the outwardly diffusing electrons; see the caption of Table \\ref{tab:mods}.) For the {\\it XMM-Newton} flare on October 3 2002 (Model C), the total energy carried away by the escaping high-energy electrons is about $7\\times 10^{40}$ ergs which for a few days can sustain the radio flare observed 13 hours after the X-ray event \\citep{Zhao04}. This picture also explains the weak correlation between the {\\it Chandra} X-ray flares and the radio emission from Sagittarius A* \\citep{Baganoff032}, because the flux of high energy electrons produced during these flares is much smaller than that for the {\\it XMM-Newton} flare;  no strong enhancement of radio emission is expected following a {\\it Chandra}-type of X-ray flare.  One of the intriguing properties of Sagittarius A*'s radio emission is that it is circularly polarized below $100$ GHz, even though no linear polarization has been observed there \\citep{Bower02}. However, in the mm and sub-mm range only strong linear polarization is observed \\citep{Aitken00, Bower03}. The mm/sub-mm polarization is likely associated with the structure of the magnetic field at small radii \\citep{Agol00, Melia00, Bromley01}, which is consistent with the spectral formation at these wavelengths discussed in this paper. The observed circular polarization may be due to anisotropy of the escaping electrons as they are transported along magnetic field lines toward larger radii under the influence of synchrotron losses \\citep{McTiernan90}. Relativistic electrons beamed along magnetic field lines produces synchrotron emission with significant circular polarization while the degree of linear polarization can be suppressed by irregularities in the source magnetic field \\citep{Epstein73a, Epstein73b}. Work to demonstrate this feature self-consistently is in progress, and the results will be reported elsewhere. The model we have presented here shows promise in being able to account not only for the spectral characteristics of Sagittarius A*, but also for its polarization characteristics and time variation properties.  Finally, we emphasize that although we suggested that $\\eta$, $k_{\\rm max}/\\Omega_{\\rm e}$, $\\alpha$, and $a$ may not change significantly over time, and have fixed them in all the models discussed above, given the dramatic changes of the physical conditions during a flare state, these parameters may also vary slightly, giving rise to different emission spectra than those presented here. Indeed, if the infrared upper limits reported by Hornstein et al. (2002) are confirmed, one then needs to decrease the high frequency cutoff of the synchrotron emission by adjusting these parameters. This will be investigated in a more comprehensive study of particle acceleration in Sagittarius A*."
    },
    "0403/astro-ph0403164_arXiv.txt": {
        "abstract": "We re-formulate cosmological perturbations in the decaying cold dark matter model, and calculate cosmological microwave background (CMB) anisotropies. By comparing our theoretical predictions with recent observational data from the Wilkinson Microwave Anisotropy Probe (WMAP), we derive a new bound on the abundance and lifetime of decaying dark matter particles. We show that the data of WMAP alone do not prefer the decaying cold dark matter model:  the lifetime is constrained to $\\Gamma^{-1} \\ge 123$ Gyr at $68\\%$ C.L. ($52$ Gyr at 95.4$\\%$ C.L.) when cold dark matter consists only of such decaying particles. We also consider a more general case in which cold dark matter consists of both stable and decaying particles, and show that the constraint generalizes down to $\\Omega_{\\rm DDM}h^2 \\la -0.5(\\Gamma^{-1}/1{\\rm Gyr})^{-1}+0.12$ for $\\Gamma^{-1} \\ge 5\\mbox{Gyr}$ at $95.4\\%$ C.L.. These bounds are robust and widely applicable, because they are derived only from gravitational effects and therefore do not rely on the details of decay channels or decay products. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403352_arXiv.txt": {
        "abstract": "We compute the properties of dark matter halos with mass $10^{6.5}-10^9\\msun$ at redshift $z=6-11$ in the standard cold dark matter cosmological model, utilizing a very high resolution N-body simulation. We find that dark matter halos in these mass and redshift ranges are significantly biased over matter with a bias factor in the range $2-6$. The dark matter halo mass function displays a slope of $2.05\\pm 0.15$ at the small mass end. We do not find a universal dark matter density profile. Instead, we find a significant dependence of the central density profile of dark matter halos on halo mass and epoch with $\\alpha_0=0.4-1.0$; the high-mass ($M\\ge 10^8\\msun$) low-redshift ($z\\sim 6$) halos occupy the high end of the range and low-mass ($M\\sim 10^{7}\\msun$) high-redshift ($z\\sim 11$) halos occupy the low end. Additionally, for fixed mass and epoch there is a significant dispersion in $\\alpha_0$ due to the stochastic assembly of halos. Our results fit a relationship of the form $\\alpha_0=0.75((1+z)/7.0)^{-1.25}(M/10^7\\msun)^{0.11(1+z)/7.0}$ with a dispersion about this fit of $\\pm 0.5$ and no systematic dependence of variance correlated with environment. The median spin parameter of dark matter halos is $0.03-0.04$ but with a large lognormal dispersion of $\\sim 0.4$. Various quantities are tabulated or fitted with empirical formulae. ",
        "introduction": "The reionization epoch is now within the direct observational reach thanks to rapid recent observational advances in two fronts --- optical quasar absorption from Sloan Digital Sky Survey (SDSS) (Fan \\etal 2001; Becker \\etal 2001) and the Wilkinson Microwave Anisotropy Probe (WMAP) experiment (Kogut \\etal 2003). The picture painted by the combined observations, perhaps not too surprisingly, strongly suggests a complex cosmological reionization process, consistent with the double reionization scenario (Cen 2003). It may be that this is the beginning of a paradigm shift in our focus on the high redshift universe: the star formation history of the early universe can now be observationally constrained.  It thus becomes urgent to theoretically explore galaxy and star formation process at high redshift in the dark age ($z\\ge 6$). In the context of the standard cold dark matter model it is expected that stars within halos of mass $10^7-10^9\\msun$ at high redshift play an important, if not dominant, role in determining how and when the universe was reionized. Furthermore, these fossil halos may be seen in the local universe as satellites of giant galaxies. This linkage may potentially provide a great leverage to nail down the properties of the high redshift galaxies.  In this paper, as a step towards understanding galaxy formation at high redshift, we investigate the properties of dark matter halos at $z\\ge 6$, using very high resolution TPM N-body (Bode \\etal 2001; Bode \\& Ostriker 2003) simulations. While there is an extensive literature on properties of halos at low redshift, there is virtually no systematic study of dark halos at $z\\ge 6$. The LCDM simulation has a comoving box size of $4h^{-1}$Mpc with $512^3=10^{8.2}$ particles, a particle mass of $m_p=3.6\\times 10^4 \\ h^{-1}\\msun$, and comoving gravitational softening length of $0.14 \\ h^{-1}$kpc. These resolutions allow us to accurately characterize the properties of halos down to a mass $10^{6.5} \\ h^{-1}\\msun$ (having about $100$ particles within the virial radius). The outline of this paper is as follows. The simulation details are given in \\S 2. In \\S 3 we quantify properties of dark matter halos in the mass range $10^{6.5}-10^9\\msun$, including the mass function, bias and clustering properties, density profile distribution, angular momentum spin parameter distribution, internal angular momentum distribution and peculiar velocity distribution. We conclude in \\S 4. ",
        "conclusions": "Using a high resolution TPM N-body simulation of the standard cold dark matter cosmological model with a particle mass of $m_p=3.57\\times 10^4 \\ h^{-1}M_\\odot$ and a softening length of $\\epsilon=0.14 \\ h^{-1}$kpc in a $4h^{-1}$Mpc box, we compute various properties of dark matter halos with mass $10^{6.5}-10^9\\msun$ at redshift $z=6-11$. We find the following results.  \\noindent (1) Dark matter halos at such small mass at high redshifts are already significantly biased over matter with a bias factor in the range $2-6$.  \\noindent (2) The dark matter halo mass function displays a slope at the small end $2.05\\pm 0.15$.  \\noindent (3) The central density profile of dark matter halos are found to be in the range $(0.4-1.0)$ well fitted by $\\alpha_0=0.75((1+z)/7.0)^{-1.25}(M/10^7\\msun)^{0.11(1+z)/7.0}$ with a dispersion of $\\pm 0.5$, in rough agreement with the theoretical arguments given in Ricotti (2003) and Subramanian \\etal (2000).  \\noindent (4) The median spin parameter of the dark matter halos is $0.03-0.04$ but with a lognormal dispersion of $\\sim 0.4$. The angular momentum profile at the small end is approximately linear with the fraction of mass in a halo having specific angular momentum less than a certain value is roughly $0.5$ times the ratio of that value over the average specific angular momentum of the halo.  \\noindent (5) The dark matter halos move at a typical velocity in excess of $30-40$km/s."
    },
    "0403/astro-ph0403678_arXiv.txt": {
        "abstract": "When the ion mean free path much exceeds the Larmor radius in a plasma, the viscous stress tensor is altered dramatically, and depends only upon quantities measured along the field lines.  This regime corresponds to typical interstellar medium conditions in galaxies and protogalaxies, even if the magnetic field is extremely weak, with a negligible Lorentz force on all scales of interest. In this work, the only role of the magnetic field is to channel angular momentum transport along its lines of force. We show that differential rotation in such a gas is highly unstable, with a maximum growth rate exceeding that of the magnetorotational instability.  The regime of interest has been treated previously by plasma kinetic methods. Where there is overlap, our work appears to be in agreement with the kinetic results.  The nonlinear outcome of this instability is likely to be a turbulent process, significantly augmenting the magnetorotational instability, and important to the initial phases of the amplification of small galactic magnetic fields. ",
        "introduction": "The magnetorotational instability, or MRI, has become central to our understanding of turbulent angular momentum transport in accretion disks (e.g. Balbus 2003).  The instability is important even when (indeed, especially when) the magnetic energy density is small compared with the thermal energy density.  It is this feature, the fact that the stability of the gas is hypersensitive to the presence of subthermal magnetic fields of any geometry, that endows the MRI with its special significance. In fact, the MRI is merely one manifestation of much more general behavior.  By imparting new degrees of freedom to a fluid, magnetic fields allow free energy gradients to become sources of instability, with important consequences for a variety of astrophysical systems: accretion disks become turbulent when the angular velocity (not angular momentum) decreases outward, and stratified dilute gases are destabilized when the temperature (not entropy) decreases upwards \\citep{b01}. %  The thermal destabilization caused by magnetic field is especially noteworthy because it is engendered by what seems at first a purely dissipative process: thermal diffusivity.   It is the extreme anisotropy of the conductivity tensor parallel and perpendicular to the field lines that lies at the heart of the instability.  By channeling heat only along lines of magnetic force, small-amplitude ripples along initially isothermal field lines grow into large fluid displacements parallel to the temperature gradient.  This is not, as is sometimes thought, analogous to classical double-diffusive instabilities, such as ocean layer ``salt-fingering''\\footnote{See Menou, Balbus, \\& Spruit (2004) for a true salt-fingering analogy involving the MRI.}.  Rather, it depends wholly upon the properties of anisotropic conductivity.  In this paper we show that the anisotropy introduced into the viscous stress tensor by a weak magnetic field sharply destabilizes dilute astrophysical disks, even without Lorentz forces appearing in the fluid equations.  By ``dilute,'' we mean the limit in which the ion Larmor radius is small compared to a mean free path, which in turn is small compared with the characteristic macroscopic length scales of the disk. We shall refer to the resulting viscous stress tensor in this limit as Braginskii viscosity \\citep{b65}.   The Braginskii limit is appropriate for interstellar and galactic disks (especially protogalactic disks, cf. Malyshkin \\& Kulsrud 2002). A remarkable property of the instability is that in the absence of Lorentz forces, when the viscous diffusivity much exceeds the resistive diffusivity, rapid growth times are associated with arbitrarily high wave numbers.  Except for isolated field geometries (e.g. precisely azimuthal), there is no formal high wavenumber dissipation of the linear magneto-viscous instability.  By way of contrast, for the ultra-weak magnetic fields considered here, the wavelength of maximum growth in the classical MRI would be strongly damped by an ordinary isotropic viscosity. (Note that in real systems, at sufficiently high wave numbers, Lorentz forces will ultimately stabilize.)  The magnetic stability of a dilute plasma may also be studied using a plasma kinetic approach.  This offers rigor in a problem where it is clearly of some benefit, but at a price of greater mathematical complexity.  \\citet{qdh02} analyzed the MRI in the collisionless regime, and noted that the character of the instability changes when the pressure tensor becomes significantly anisotropic, with growth rates in excess of the classical MRI maximum.  In the current work, we have in essence abstracted the anisotropic components of the pressure, and labeled them as a viscosity tensor.  More recently, Sharma, Hammett \\& Quataert (2003) reanalyzed the kinetic problem using a Krook collision operator, and showed that the collisional and collisionless behavior are on the same branch of the dispersion relation.  In the short mean free path limit, the pressure tensor associated with the equations used by Sharma et al. 2003 reduces to a scalar pressure plus the dominant parallel components of the Braginskii viscosity (Snyder, Hammett, \\& Dorland 1997).  The current work, a purely gasdynamical treatment of the problem, affords relative mathematical simplicity and a physically transparent interpretation.  In either its kinetic or gasdynamical guise, this vigorous instability may be important in the early stages of magnetic field amplification in disk galaxies, when densities are low, temperatures are relatively high, and the field is likely to be very weak.  Quataert et al.\\ (2002) and Sharma et al.\\ (2003) discuss applications to black hole accretion flows of low radiative efficiency.  An outline of the paper is as follows.  Section 2 is a discussion of the regime of the applicability of this work and a presentation of the magnetic viscosity formalism.  Section 3 is the heart of the paper, formulating and solving the problem, and checking the validity of the results.  Section 4 is a discussion of the applicability of our results to galactic magnetism, and section 5 summarizes our findings. ",
        "conclusions": "The instability presented in this paper is generic and powerful.  Its most rapidly growing behavior is exhibited at large vertical wavenumbers, and even in the presence of finite resistivity, it is not easily quenched. Long radial and azimuthal wavelengths are highly unstable, imparting large scale coherence in the disk plane already in the linear stages of instability.  In this section we discuss the possibility that this mechanism could be an important part of the process that gives rise to galactic magnetic fields.   Generally, galactic magnetic amplification processes are divided into two principal categories \\citep{beck96}: (1) differential wind-up of a primordial field (Kulsrud 1986); and (2) dynamo amplification of a similar seed field.  The latter category itself consists of at least two very distinct mechanisms; (2a) a classical $\\alpha\\Omega$ dynamo (Parker 1979), and (2b) small scale turbulent amplification (Schekochihin et al. 2004). The weak field limit of these mechanisms are thought to be kinematic in nature.  It is difficult to see how the process outlined in this paper could be unimportant, at least quantitatively, to any of these processes.   Simple wind-up of a seed field by differential rotation fails to incorporate properly the true dynamics of either the MRI or the magneto-viscous instability described here.  Differential rotation and essentially any field geometry is highly unstable, resulting in large radial motions and ultimately MHD turbulence.  The shearing of radial fields of course readily occurs, but it is not the primary amplification process of a differentially rotating system.  Traditional $\\alpha\\Omega$ galactic dynamo models rely on the presence of favorable properties of interstellar turbulence to generate the required mean helicity.  Curiously, these theories of magnetic field amplification generally do not incorporate weak-field MHD instabilities, despite the latter's obvious potential benefits (e.g. Brandenburg et al. 1995; Hawley, Gammie, \\& Balbus 1996).  Conversely, without something like the MRI or the magneto-viscous instability, it is not so obvious that turbulence will generally amplify the field, at least at low to moderate values of ${\\cal P}$.  An example of dissipative behavior is seen in a simulation of Hawley et al.\\ (1996).  The combination of hydrodynamical shear-layer turbulence plus a magnetic field lead not to dynamo activity, but to field energy loss.  This result emerged despite the fact that calculation was done not in the kinematic limit, but fully in the MHD regime. However, the combination of local Coriolis, tidal, and Lorentz forces dramatically altered the dynamo properties of the ensuing turbulence: rapid and significant field amplification was observed, driven by the MRI turbulence, in a hydrodynamically stable background.  More recently, the notion that non-helical homogeneous small scale turbulence may play a key role in galactic dynamos, particularly in the early stages, has been investigated in a series of numerical simulations (Schekochihin et al. 2004; see also Zeldovich et al. 1984, Kulsrud \\& Anderson 1992), which include a scalar viscosity and study the non-kinematical regime.  The idea is that certain magnetic field configurations (termed ``winning'' by Schekochihin et al.) align themselves smoothly with the stretching direction of the strain tensor of the turbulent flow, but fluctuate along the corresponding null axis, so that the work done on the field by fluid element stretching is not undone by relaxation.  A sort of turbulent ratchet thereby ensues, growing the field with enormous efficiency at small scales.  Numerical simulations of homogeneous white noise forcing conducted by Schekochihin et al. 2004 resulted in exponential field amplification, but, interestingly, only in the regime of large ${\\cal P}$. What the relationship is between these small scale structures and a galactic scale field has yet to be established.  Further discussion of the pros and cons of this mechanism would take us too far afield, but we may note that the generation of coherent magnetic field structure on scales larger than the disk scale height requires time scales at least as long as $1/\\Omega$.  And while it may well be that intrinsic interstellar turbulence plays a role in the amplification of galactic magnetic fields, we emphasize here a phenomenon that is certainly unavoidable: differential rotation.  In fact, there is no reason to restrict oneself to galactic radius scales: if sub-galactic, differentially rotating turbulent vorticies are also present, magneto-viscous modes should be seen on these scales as well. The key point is that a kinematic approach will miss this process, which is intrinsically MHD.  One compelling origin of the seed field for the galactic amplification process is the stars themselves (Biermann 1950, Rees 1993), in which both battery processes and a truly powerful, rapid dynamo are likely to be present.  A 0.1 G azimuthal surface field diluted in a stellar wind to interstellar scales would give rise to a seed field of $\\sim 2\\times 10^{-9}$G. With $T\\simeq 10^4$ K, $n\\simeq 1$ cm${}^{-3}$, one finds $$ \\omega_{ci}\\tau_{ci} \\simeq 10, \\quad {\\cal P} \\sim 10^{11} $$ and the regime of the magneto-viscous instability is valid. The ratio of gas pressure to magnetic pressure is $3.5\\times 10^7$, an extremely weak field by this measure. The linear amplification factor per orbit of the magnetic energy, as noted in \\S 3.2, is huge: $5.2\\times 10^7$. The makings of an MHD dynamo would seem to be present.  If a combination of magneto-viscous and magneto-rotational processes is to be a viable candidate for galactic field production, it needs to be shown that (1) the bulk of the field energy emerges in the largest scales; and (2) the saturated field energy density can grow to levels comparable to the thermal pressure.  Definitive answers will require a numerical treatment, but in the meantime the following points may be considered.  A magnetic energy spectrum dominated by the largest scales appears to be a universal outcome of MRI simulations (e.g. Hawley, Gammie, \\& Balbus 1995), despite the fact that the most unstable modes occur at scales much smaller than the scale height.  In such calculations, the available dynamical range is limited.   To the extent it can be measured, however, the inertial range is Kolmogorov-like in the magnetic energy. In hydrodynamics, Kolmogorov scaling for the kinetic energy power spectrum is universal, if the dissipation rate per unit volume is the sole constant characterizing the cascade, as often seems to be the case. Universal processes may also be at work in MHD turbulence (Kraichnan 1963; Goldreich \\& Sridhar 1997).  An important caveat, however, is that most MHD turbulence studies have been based on wave-wave interactions (which may be of secondary importance in linearly unstable rotating systems), and they have yet to address the role of a magnetized viscosity.  Indeed, a magnetized viscosity may be crucial to understanding why interstellar fields are thermal strength (or even slightly above), whereas all numerical MRI simulations to date have yielded subthermal saturated field strengths.  It has been argued (Balbus \\& Hawley 1998) that the outcome may be well be sensitive to ${\\cal P}$.  The point is that large scale reconnection proceeds relatively easily in simulations when the resistive scale is comparable to or larger than the viscous scale.  In ``ideal MHD,'' both scales are grid based.  But matters are likely to be very different if the viscous scale is, say, five orders of magnitude larger.  In that case, enormous viscous stresses would occur in the course of setting up a reconnection front, and prevent its formation. With reconnection stifled, MHD turbulence could amplify the field to its natural dynamical limit: the thermal energy density.  Beyond this point, buoyant effects would make it difficult for a suprathermal field to remain in the disk and be further amplified (Parker 1979).  Large Prandtl number simulations have in fact recently begun (Schekochihin et al. 2004), and treatments of the anisotropic Braginskii viscosity and conductivity have yet to be attempted.  With the resistivity scale hidden below the viscous dissipation scale, there is no effective small scale sink for the magnetic field, and is therefore not surprising to note that Schekochihin et al.\\ find that the magnetic energy spectra increases with wave number on subviscous scales, before it is ultimately cut-off. By way of sharp contrast, the ideal MRI simulations noted above find a monotonically decreasing energy spectrum for the magnetic field (Hawley, Gammie, \\& Balbus 1995).  The magneto-viscous instability is a powerful and general mechanism to amplify weak magnetic fields in galaxies, or in hot dilute plasmas more generally, and is worthy of more detailed study.  Of particular interest would be a suite of numerical simulations designed to isolate the differences between high Prandtl number turbulent dynamos relying on MHD and differential rotation, those based on random forcing only, and those containing both."
    },
    "0403/astro-ph0403022_arXiv.txt": {
        "abstract": "Spatial distributions of energy deposited by an extensive  air shower in the atmosphere through ionization, as obtained from the CORSIKA simulation program, are used to find the fluorescence light distribution in the optical image of the shower. The shower image derived in this way is somewhat smaller than that obtained from the NKG lateral distribution of particles in the shower. The size of the image shows a small dependence on the primary particle type. ",
        "introduction": "\\label{intro} The fluorescence method of extensive air shower detection is based on recording light emitted by air molecules excited by charged particles of the shower. For very high energies of the primary particle, enough fluorescence light is produced by the large number of secondaries in the cascading process so that the shower can be recorded from a distance of many kilometers by an appropriate optical detector system \\cite{greisen,Bal}. As the amount of fluorescence light is closely correlated to the particle content of a shower, it provides a calorimetric measure of the primary energy.  The particles in an air shower are strongly collimated around the shower axis. Most of them are spread at distances smaller than several tens of meters from the axis, so that when viewed from a large distance, the shower resembles a luminous point on the sky. Therefore, a one-dimensional approximation of the shower as being a point source might be adequate in many cases regarding the shower reconstruction. For more detailed studies, however, the spatial spread of particles in the shower has to be taken into account. This is especially important for nearby showers, where the shower image, i.e.~the angular distribution of light recorded by a fluorescence detector (FD),  may be larger than the  detector resolution.  The image of a shower has been studied in Ref.~\\cite{sommers}, where it was shown that for a disk-like distribution of the light emitted around the shower axis, the shower image has a circular shape, even when viewed perpendicular to the shower axis. Analytical studies including lateral particle distributions parameterized by the Nishimura-Kamata-Greisen (NKG) function or estimates based on average particle distributions taken from CORSIKA~\\cite{heck} were discussed in Ref.~\\cite{dgora} and Ref.~\\cite{giller}, respectively.  In this paper, detailed Monte Carlo simulations of the shower image based on the spatial energy deposit distributions of individual showers are performed. By using the energy deposit of the shower particles as calculated by CORSIKA~\\cite{markus3}, the previous simplified assumption of a constant fluorescence yield per particle is avoided. Assuming a proportionality between the fluorescence yield and ionization density, the light emitted by each segment of the  shower is determined. A concept is developed to treat the shower as a three-dimensional object, additionally taking into account the time information on photons arriving at the FD. In contrast to previous analytical studies, shower fluctuations as predicted by the shower simulation code are preserved and studied. Propagation of the light towards the detector, including light attenuation and scattering in the atmosphere is simulated, so that  the photon flux at the detector is calculated. The resulting  distribution of photons arriving simultaneously at the detector, i.e. the shower image, is compared to results obtained by using the NKG approximation of particle distribution in the shower. The comparison is performed for different shower energies and different primary particles. In particular, it is checked whether the shower width depends on the primary particle type. \\\\ \\\\ The plan of the paper is the following: definition of the shower width and algorithms of fluorescence light production are described in Section 2. In Section 3 the size of shower image in the NKG and CORSIKA approaches is calculated and  its dependence on primary energy, zenith angle and primary particle is discussed. Conclusions are given   in Section 4. ",
        "conclusions": "Shower image simulations more accurate than available until now are presented, which incorporate  a more realistic distribution of fluorescence light emitted by the shower. The image simulations are based on distributions of energy deposited by the shower in  air as derived from CORSIKA. A comparison of the size of the shower image obtained using CORSIKA and that given by the NKG function was made for different energies and primary particles. To a first approximation, the results of these two completely independent methods (analytical versus Monte Carlo) show quite reasonable agreement.  The image spot size derived from CORSIKA is smaller by about 15\\% compared to the NKG approximation. This difference is mainly due to the differences in lateral particle distributions in the NKG and CORSIKA approximation.  The energy deposit distribution from CORSIKA leads to a dependence of the size of shower image on the primary particle, so that studies of the shower image may be helpful  for the  primary particle identification.  {\\it Acknowledgements.} We  would like to thank R. Engel and F. Nerling   for fruitful discussions and careful reading of the manuscript. This work was partially supported by the Polish Committee for Scientific Research under grants  No. PBZ KBN 054/P03/2001  and 2P03B 11024 and by the International Bureau of the BMBF (Germany) under grant No. POL 99/013."
    },
    "0403/astro-ph0403508_arXiv.txt": {
        "abstract": "An elementary kinematic model for emission produced by relativistic spherical colliding shells is studied. The case of a uniform blast-wave shell with jet opening angle $\\theta_j \\gg 1/\\Gamma$ is considered, where $\\Gamma$ is the Lorentz factor of the emitting shell.  The shell, with comoving width $\\Delta r^\\prime$, is assumed to be illuminated for a comoving time $\\Delta t^\\prime$ and to radiate a broken power-law $\\nu L_\\nu$ spectrum peaking at comoving photon energy $\\e_{pk,0}^{\\prime}$.  Synthetic GRB pulses are calculated, and the relation between energy flux and internal comoving energy density is quantified.  Curvature effects dictate that the measured $\\nu F_\\nu$ flux at the measured peak photon energy $\\e_{pk}$ is proportional to $\\e^3_{pk}$ in the declining phase of a GRB pulse. Possible reasons for discrepancy with observations are discussed, including adiabatic and radiative cooling processes that extend the decay timescale, a nonuniform jet, or the formation of pulses by external shock processes. A prediction of a correlation between prompt emission properties and times of the optical afterglow beaming breaks is made for a cooling model, which can be tested with Swift. ",
        "introduction": "In the collapsar scenario for GRBs, pulses in GRB light curves are thought to be produced by collisions between relativistic shells ejected from a central engine (see \\citet{zm04} for a recent review). The interception of a more slowly moving shell by a second shell that is ejected at a later time, but with faster speed and larger Lorentz factor, produces a shock that dissipates internal energy to energize the particles that emit the GRB radiation. This scenario is widely considered to explain pulses in GRB light curves \\citep{kps97,dm98}. Studies of pulses are important to decide if GRB sources require engines that are long-lasting or impulsive \\citep{dm03}, with important implications for the nature of the central engine, which is often argued to be a newly formed black hole powered by the accretion of a massive dense torus.  Here we construct an elementary kinematic model for colliding shells, assumed spherical and uniform within jet opening angle $\\theta_j$. This is the sort of jet that \\citet{fra01} discuss regarding the standard energy reservoir result, where jet opening angles are inferred from the time of achromatic spectral breaks in optical afterglow light curves.  We also perform this study in order to quantify the curvature constraint of a spherically emitting shell traveling with bulk Lorentz factor $\\Gamma$, which implies that the shell radius \\begin{equation} r \\approx 2\\Gamma^2 c t_{var}/(1+z)\\; \\label{r} \\end{equation} in order to produce variability on timescale $t_{var}$ \\citep{rl79,fmn96}.  This study also quantifies the rate at which flux decays at a given energy due to curvature effects, and the range of validity of the approximate relation \\begin{equation} \\Phi_E \\cong c r^2 u_0^\\prime \\Gamma^2/d_L^2\\; \\label{phie} \\end{equation} between internal comoving energy density $u_0^\\prime$ and observed energy flux $\\Phi_E$, where $d_L$ is the luminosity distance (See Appendix A). The accuracy of this relation is important to quantify $\\gamma\\gamma$ opacity constraints \\citep{ls01,der04} applied to GRB pulses as measured with the GRB monitor and Large Area Detector on GLAST\\footnote{glast.gsfc.nasa.gov}, as well as to make estimates of photomeson production in GRB blast waves \\citep{wb97}.  If curvature effects dominate the late-time emission in GRB pulses, then a unique relation is found whereby the value of the $\\nu F_\\nu$ peak flux $f_{\\epk}$(in cgs units of ergs cm$^{-2}$ s$^{-1}$) at peak photon energy $\\e_{pk}$ decays in proportion to $\\propto\\e_{pk}^3$. This relation is generally not observed in long, smooth GRB pulses studied by \\citet{br01}, who find power-law decays $f_{\\epk}\\propto \\epk^\\zeta$, with $0.6 \\lesssim \\zeta \\lesssim 3$. Remarkably, values of $\\zeta$ for different pulses within the same GRB are confined to a rather narrow band of values. The wide range of values of $\\zeta$ are found not only in multi-peaked GRBs, but also in single-peaked GRBs that display smooth fast-rise, slow-decay light curves \\citep{br01,rp02}. The smooth single peak GRBs could arise from curvature effects \\citep{fmn96}, or to external shocks \\citep{dbc99}. For GRB pulses that could be produced by spherically symmetric shell collisions, discrepancy with observations suggest a breakdown of our assumptions.  In the next section, the kinematic model is presented. Calculations based on this model are presented in Section 3. In Section 4, we discuss the possibility that radiative-cooling effects produce the power-law relation, implying a prediction that can be tested with Swift\\footnote{swift.gsfc.nasa.gov}. Alternately, the uniform spherical shell assumption could break down, or the basic model of colliding shells could be in error. The Appendices give derivations of simple, widely-used approximations related to this study, a derivation of the curvature relation $f_{\\epk} \\propto \\e_{pk}^3$, as well as an analytic form for the time-dependent pulse profile, leading to a simple expression for the light curve of a pulse in the curvature limit. A brief summary is given in Section 5. ",
        "conclusions": "The estimate $L \\cong 4\\pi d_L^2 \\Phi_E \\cong 4\\pi r_0^2 c u_0^\\prime \\Gamma^2$, where the received flux is intensified by two powers of $\\Gamma$ for the relativistic time contraction and photon energy enhancement in a blast wave geometry, is generally used to relate bolometric energy flux and internal energy density (Appendix A).  More remarkably, the allowed radius of the radiating spherical shell is $\\approx 2 \\Gamma^2 $ times larger than inferred through causality arguments applied to the measured variability time scale. This effect greatly dilutes the comoving photon density compared with a stationary emitting region, and essentially explains the unusual properties of GRBs. In total, we see that \\begin{equation} \\Phi_{E} \\cong {cr_0^2 u_0^\\prime\\Gamma^2\\over d_L^2}= {4c^3 t_{var}^2\\over (1+z)^2 d_L^2}\\;u_0^\\prime \\Gamma^6\\; \\propto u_0^\\prime \\Gamma^6. \\label{phie1} \\end{equation} For the nominal parameters used in the figures, $\\Phi_E \\cong 4.8\\times 10^{-11} t_{var}^2 ({\\rm s})$ $u_0^\\prime \\Gamma_{300}^6$ ergs cm$^{-2}$ s$^{-1}$.  The most rigorous limits on $\\gamma\\gamma$ attenuation are obtained by determining the {\\it minimum value} of the product $u_0^\\prime \\Delta \\tp$ that can produce a pulse with a measured peak flux $f_{\\e_{pk}}$ and full-width half-maximum duration $t_{1/2}$ for a given value of $\\Gamma$. It is the product $u_0^\\prime \\Delta \\tp$ that enters into the $\\gamma\\gamma$ attenuation (and photomeson) calculations. If the shell is found to be optically thick at some photon energy for a given value of $\\Gamma$ even in this case, then $\\Gamma$ must be larger if photons with the corresponding energies are detected.  Inspection of the various cases shows that the pulse formed in the curvature limit produces the brightest measured flux and shortest duration for a given value of the product $u_0^\\prime \\Delta \\tp$. This is because the measured duration is due entirely to curvature effects, and the radiated energy is compressed into the shortest duration and brightest pulse in this limit. From the results of Appendices C and D, this implies that the expression \\begin{equation} u_0^\\prime \\Delta \\tp \\cong {[2^{1/(3-a)}-1](1+z) d_L^2 f_{\\e_{pk}} \\over 8 c^3 \\Gamma^5 t_{1/2}} \\label{u0cdt} \\end{equation} gives the smallest possible values for $u_0^\\prime c \\Delta \\tp$, and this expression will therefore yield the most reliable minimum Lorentz factors for $\\gamma\\gamma$ attenuation calculations derived from {\\it GLAST} or ground-based air Cherenkov telescope observations.  The corresponding expression for the comoving photon spectral energy density is thus \\begin{equation} u^\\prime_{\\e^\\prime} \\cong {0.26(1+z) d_L^2 f_{\\e_{pk}} \\over 8 c^3 \\Gamma^5 t_{1/2} \\Delta \\tp} [x^a H(1-x)+x^b H(x-1)]\\;, \\label{ueprime} \\end{equation} where $t_{1/2}$ is determined at photon energies near the peak of the $\\nu F_\\nu$ spectrum. Eq.\\ (\\ref{ueprime}) is  a factor of 3 smaller than the expression used for a comoving spherical blob with $\\delta \\rightarrow \\Gamma$ and $t_{var} \\rightarrow t_{1/2}$ (see Eq.\\ [2] in \\citet{der04}).  Three generic types of pulses have been identified for the simple kinematic pulse, namely the curvature case where $r_0 \\gg c\\Gamma\\Delta \\tp$, the causal case where $r_0\\approx\\Gamma\\Delta r^\\prime \\approx c\\Gamma\\Delta \\tp$, and the thin shell case where $\\Delta r^\\prime \\ll c\\Delta \\tp$. In all three types of kinematic pulses, curvature effects dominate the formation of the spectrum at late time $t \\gg (1+z) \\Delta \\tp/2\\Gamma$, so that $f_{\\e_{pk}} \\propto \\epk^3$ if curvature effects dominate pulse formation at late times.  The curvature relationship can be derived from a simple scaling argument by noting that the differential stationary-frame shell volume which contributes to the received flux, given by $dV =2\\pi r_0^2 \\Delta r d\\mu$, remains constant with time. This is because the relation between reception time $t$ and $\\mu$ for a shell that is instantaneously illuminated at comoving time $\\tp_0$ is $t = (1+z)\\Gamma\\tp_0(1-\\beta\\mu)$, so that $d\\mu \\propto dt$. The $\\nu F_\\nu$ flux $f_\\e =\\delta^4 L^\\prime/4\\pi d_L^2 = \\delta^4 V^\\prime \\ep j(\\ep )/4\\pi d_L^2 = \\delta^3 V \\ep j(\\ep )/4\\pi d_L^2$, where $L^\\prime$ is the comoving luminosity of the emitting volume that contributes to the flux at time $t$. For an emission spectrum that is flat, that is, $ \\ep j(\\ep) \\propto \\e^{\\prime 0}$, $f_{\\e_{pk}} \\propto \\epk^3$ because $\\epk \\propto \\delta$ for a uniform shell.  Analysis of BATSE GRB light curves (Borgonovo and Ryde 2001) shows that the peak fluxes of a GRB pulse generally follow a relation whereby \\begin{equation} f_{\\e_{pk}}\\propto \\epk^\\zeta\\;. \\label{fepk} \\end{equation} Values of $\\zeta$ for different GRBs vary over a wide range from $\\approx 0.6$ to $3$, with values of $\\zeta$ roughly constant for pulses within the same GRB or in a GRB consisting of a single smooth pulse.  In most GRBs, therefore, curvature effects do not make a large contribution to the decay phase of a GRB light curve.  An interesting question is the source of the difference of observations from our kinematic model pulses. One possibility is that the jet has angular structure, and varies with directional energy release and baryon-loading on angles $\\theta \\approx $ few$\\times \\Gamma^{-1}$. The angle-dependent speeds in such a system would produce a deformed colliding shocked fluid shell where the spherical symmetry assumption fails, as therefore would the uniform jet model. If this is the case, then GRB prompt emission data can in principle be analyzed to reveal shell structure and to determine whether this behavior is consistent with a universal jet structure (\\citet{zha04}; see \\citet{fra03} for review).  Rather than treat these geometrical effects here, we consider instead whether radiation effects could form a power-law relation between $f_{\\e_{pk}}$ and $ \\epk$. The most naive system considers a fixed volume of shocked fluid within which the peak of the $\\nu F_\\nu$ spectrum is made by a large population of quasi-monoenergetic electrons that radiates most of the power through the synchrotron process in a mean magnetic field of strength $B$. If these electrons mainly have comoving Lorentz factors $\\gamma$, then their luminous power $\\propto B^2\\gamma^2$. Because the peak of the $\\nu F_\\nu$ spectrum is $\\propto B \\gamma^2$, and assuming $B$ is constant, then $ f_{\\e_{pk}}\\propto \\epk$ or $ \\zeta_{syn} = 1$ for this simple synchrotron model with constant magnetic field. This model can therefore only apply in rare cases.  A better treatment must consider the evolution of $\\gamma$ due to synchrotron and adiabatic losses in the expanding shell. The equation of electron energy evolution is given by \\begin{equation} -{d\\gamma\\over d\\tp} = {1\\over V^\\prime_{sh}} {dV^\\prime_{sh}\\over d\\tp}\\; {\\gamma\\over 3} + {\\sigma_{\\rm T}B^2(\\tp )\\over 6\\pi m_ec}\\;\\gamma^2\\;, \\label{dgammadt} \\end{equation} where the comoving shell volume changes with time according to $V^\\prime_{sh} \\propto t^{\\prime 3m}$, with $m = 0$ corresponding to no expansion, and $m = 1$ corresponding to 3-dimensional expansion.  The magnetic field will also change as a result of the expansion of the shell volume. In the flux freezing limit where the magnetic field is randomly oriented, $BR^2 \\propto const$, implying $B \\propto V_{sh}^{\\prime 2/3}\\propto t^{\\prime -2 m}$. The well-ordered magnetic field required to explain the polarization observation of the GRB 021206 observed with RHESSI \\citep{cb03} suggests that there is not an efficient mixing and randomization of the magnetic field directions. For simplicity, we therefore write $B \\propto t^{\\prime -2 v m}$, where $v = 1$ gives the flux-freezing limit.  Eq.\\ (\\ref{dgammadt}) becomes \\begin{equation} -{d\\gamma\\over d\\tau} = m\\; {\\gamma\\over \\tau} + \\nu_0\\tau ^{-4vm}\\gamma^2\\;, \\label{dgammadt1} \\end{equation} where $\\tau \\geq 1$ is a dimensionless time variable, and $\\nu_0$ is a dimensionless synchrotron energy loss rate. Eq.\\ (\\ref{dgammadt1}) is analytic, but it is sufficient to consider two limiting cases of dominant adiabatic losses or dominant synchrotron losses at late times. In the case of dominant adiabatic losses we have (dropping the primes) $\\gamma\\propto t^{ -m}$, $B\\propto t^{ -2 v m}$, and $\\e_{pk} \\propto t^{-2m (1+v)}$, so that $ f_{\\e_{pk}}/ \\epk \\propto B \\propto \\e_{pk}^{v/(v+1)}$. Thus $\\zeta_{adi} = 1 + [v/(v+1)]$. Even for a wide range of values of $v$, $1 \\lesssim \\zeta_{adi} \\lesssim 2$, and $\\zeta_{adi}$ is independent of the geometry factor $m$.  If synchrotron losses dominate the cooling, $-d\\gamma/dt \\propto B^2 \\gamma^2$.  The dependence $B\\propto t^{ -2 v m}$ therefore implies $\\gamma \\propto t^{4vm-1}$, so that $\\e_{pk} \\propto B\\gamma^2 \\propto t^{6vm-2}$. Hence $f_{\\e_{pk}}/\\e_{pk} \\propto B \\propto \\e_{pk}^{vm/(1-3vm)}$, so that $\\zeta_{syn}= 1+ [vm/(1-3vm)]$. Except when $v \\ll 1$, $\\zeta_{syn} \\approx 1$ when $m=0$ and $\\zeta_{syn} \\approx 0.5$ -- 0.67 when $m = 1$. In this simple model, therefore, values of $\\zeta$ between $1/2$ and $2/3$ are only possible for 3-dimensional expansion.  Three-dimensional expansion is more likely to occur for narrowly collimated blast waves than for blast waves with large opening angles, and the narrowly collimated blast waves would have ``beaming breaks\" in the optical afterglow light curves at earlier times. If our conjecture that the $f_{\\e_{pk}}$/$\\e_{pk}$ relationships are due to synchrotron and adiabatic effects in GRB blast waves with different opening angles, then those blast waves with $\\zeta < 1 $ should be correlated with earlier beaming break times. Because this effect is only seen when synchrotron losses dominate the cooling, GRBs with $\\zeta < 1$ should also display cooling spectra with photon indices $\\approx 3/2$ below $\\e_{pk}$.  \\citet{br01} find several GRBs and many pulses in the BATSE sample with statistically significant values of $\\zeta $ less than unity. These GRBs however preceded the afterglow era. For those GRBs which have measured beaming breaks (see Table 1 in \\citet{bfk03}), only GRB 990123 has sufficiently bright BATSE data to provide a data point for such a correlation. GRB 990123 has not yet been analyzed to give $\\zeta$, while analysis of Beppo-SAX data is in progress (F.\\ Ryde, private communication, 2004).  Such a model for the $f_{\\e_{pk}}$/$\\e_{pk}$ relationship would explain why $\\zeta$ is approximately constant for different pulses within a GRB, provided that the opening angle of the GRB jet remains the same throughout the period of activity of the GRB engine.  The adiabatic/synchrotron model would not, however, explain pulses with $2 \\lesssim \\zeta \\lesssim 3$. There are many such pulses in the \\citet{br01} sample, though generally with large error bars. If analysis of Beppo-SAX or Swift data reveal such GRBs, then another explanation is required.  One possibility is that GRB pulses are due to the interactions of a single impulsive blast wave with inhomogeneities in the surrounding medium. This version of the external shock model for the prompt phase can be much more efficient than an internal shell model \\citep{dm99,dm03}, and permits quantitative studies of the statistics of BATSE GRBs (\\citet{bd00}; see \\citet{zm04} for a review of the internal/external controversy).  Predictions for the $f_{\\e_{pk}}/\\e_{pk}$ relationship in an external shock model \\citep{dcb99} can be derived by adapting the equations for blast wave deceleration in a uniform medium with $\\Gamma = \\Gamma_0/[1+(x/x_d)^g]$, where $\\Gamma_0$ is the initial Lorentz factor, $x_d$ is the deceleration distance and $g$ is the radiative index ($g = 3/2$ and 3 for an adiabatic and fully radiative blast wave, respectively). In the deceleration phase, $x \\propto t^{1/(2g+1)}$ and therefore $\\Gamma \\propto t^{-g/(2g+1)}$. In this model, $\\e_{pk} \\propto \\Gamma B \\gamma_{pk}^2$ and $f_{\\e_{pk}} \\propto \\Gamma^2 B^2 \\gamma_{pk}^2$, where $\\gamma_{pk} \\propto \\Gamma^4\\propto t^{-4g/(2g+1)}$ in the slow-cooling regime, and $\\gamma_{pk} \\propto (x\\Gamma)^{-1}\\propto t^{-2/(2g+1)}$ in the fast-cooling regime. Thus $f_{\\e_{pk}}$/$\\e_{pk} \\propto B\\Gamma$.  In the slow-cooling regime, $\\e_{pk} \\propto \\Gamma^4\\propto t^{-4g/(2g+1)}$, and $f_{\\e_{pk}}\\propto\\e_{pk}^{3/2}$. In the fast-cooling regime, $\\e_{pk} \\propto t^{-2/(2g+1)}$ and $f_{\\e_{pk}}\\propto\\e_{pk}^{1+g}$. In the slow-cooling and fast-cooling regimes, therefore, values of $\\zeta_{sc} = 3/2$ and $\\zeta_{fc} = 1+g$, respectively, are predicted. Provided that the surrounding medium is uniform (which can be inferred from afterglow modeling, though at a larger distance scale), the slow-cooling result implies a definite value of $\\zeta_{sc}=3/2$ for fast-rise, smooth decay light curves when spectral analysis demonstrates that the GRB evolves in the slow-cooling regime.  For GRBs in fast-cooling regime, this estimate implies $5/2 < \\zeta_{fc} < 4$, and in these cases cooling spectra should be apparent. Further work will be needed to extend the results to radial density gradients of the circumburst medium, and to verify that these relations hold for deceleration in small density inhomogeneities that form GRB pulses in the external shock model."
    },
    "0403/astro-ph0403214_arXiv.txt": {
        "abstract": "{We present the first {\\it BeppoSAX} observation (0.1 to 220 keV) of the quasar \\object{Mrk~205}. We have searched for the unusual Fe line profile claimed in the {\\it XMM-Newton} spectrum which has been widely discussed in recent literature. We find no evidence for a broad, ionized Fe line component in our data. We detect for the first time a Compton hump in this object. Besides, when this component is included in the fit, the line strength diminishes, in agreement with a recent re-analysis of the {\\it XMM-Newton} data, but with better constraints on the reflection component thanks to the PDS instrument (15-220 keV). We interpret this fact as another indication for illumination of a distant and cold material rather than reprocessing in the highly ionized inner parts of an accretion disk. \\\\ We cannot constrain the presence of a high energy cutoff but we confirm the existence of a variable soft excess (one year timescale). ",
        "introduction": "Discovered by Weedman (\\cite{weedman}), the quasar \\object{Mrk~205} is seen through the southern spiral arm of the nearby spiral, barred galaxy \\object{NGC~4319} at a projected distance of roughly 40\\arcsec\\, from its center. From its radio luminosity at 6 cm, \\object{Mrk~205} is classified as a radio-quiet quasar (Rush et al. \\cite{rush}), but it is still not clear whether it belongs to the radio-quiet quasar class (e.g. Bahcall et al. \\cite{bahcall}) or to the Seyfert~1 class (e.g. Veron-Cetty \\& Veron \\cite{veron}). \\object{Mrk~205} has a redshift $z=0.07085$ (Huchra et al. \\cite{huchra}) while \\object{NGC~4319} has $z=0.00468$ (Bowen et al. \\cite{bowen1}). \\\\ \\\\ In the standard paradigm of (radio-quiet) Active Galactic Nuclei (AGN), X-ray photons result from Compton upscattering of optical-UV photons in a hot thermal corona above the accretion disk surrounding a supermassive black hole. A large fraction of the seed photons are thought to be produced by thermal emission of the accretion disk itself, but \\object{Mrk~205} was found by McDowell et al. (\\cite{mcdowell}) to have a weak blue bump.  One of the possible scenarios which could explain this spectral morphology is that the bump may be highly variable and thus observed in a particularly weak state. In the thermal comptonisation framework, blue bump variability would also induce hard X-ray variations.\\\\ A soft X-ray excess component may be interpreted as the hard tail of the blue bump (Walter \\& Fink \\cite{walterfink}; Brunner et al. \\cite{brunner}). So far, evidence for a soft excess in \\object{Mrk~205} was only reported in two observations; with {\\it EINSTEIN} IPC (Wilkes \\& Elvis \\cite{wilkeselvis}) and {\\it XMM-Newton} (Reeves et al. \\cite{reeves2}).\\\\ {\\it BeppoSAX}, with its unique broad-band capabilities allows to extend the observation of \\object{Mrk~205} for the first time above 20 keV to study the properties of the hard X-ray emission as well as of a potential reflection component. Indeed, some of the X-ray photons can be reprocessed in the surroundings, in particular in the disk itself giving rise to neutral (or low-ionization) Fe~K$\\alpha$ line.  While it is a common feature of Seyfert~1 galaxies (Nandra \\& Pounds \\cite{nandra}), only a few quasars present a clear detection of such a component (Williams et al.  \\cite{williams}; Nandra et al. \\cite{nandra2}). Furthermore, the centroid of the line was often measured at energies close to 6.7 keV suggesting that its origin may be in an ionized layer of the accretion disk (Reeves \\& Turner \\cite{reeves1}), the neutral Fe~K$\\alpha$ being emitted in distant and cold material lying outside the broad line region. Reeves et al. (\\cite{reeves2}), who analysed the {\\it XMM-Newton} data of \\object{Mrk~205}, claimed the existence of a broad ionized component in addition to the neutral one in this object. However, Page et al. (\\cite{page}) reprocessed the data including a reflection hump component in the fit and found little evidence for a broad, ionised component in the data. \\\\ In this paper, we present the analysis of three {\\it BeppoSAX} observations of \\object{Mrk~205} for a total of 200 ks of exposure time. In Section 2, we summarize the knowledge we had on \\object{Mrk~205} prior to our {\\it BeppoSAX} observation. In Section 3, we describe the data and their analysis while in Section 4 we compare our results and discuss them in the framework of the unified model. ",
        "conclusions": "\\noindent {\\bf The nature of \\object{Mrk~205}.} It has been noted that, in contrast to the Seyfert~1 case, the detection of Fe K$\\alpha$ line in quasars was more seldom and that if detected, the line energy was around 6.7 keV instead of 6.4 keV (Reeves \\& Turner \\cite{reeves1}). It was thus suggested that sensitive instruments like the EPIC camera on {\\it XMM-Newton} could observe the Fe K$\\alpha$ emission of quasar in detail. From our X-ray observations of \\object{Mrk~205}, and from the luminosity we obtained from them, we suggest to be careful in classifying this object. Its luminosity is found in the region in which it is difficult to choose for a Seyfert~1 or for a quasar. It seems thus difficult to generalize any findings made on this object to a whole class. \\\\ \\\\\\noindent {\\bf The soft excess variability.} The soft excess seems to be variable in this source although a very strong soft excess was never reported. It seems to be absent/not detected in 1997, strong in 2000, weak in 2001 and absent again in 2002 (thus varied in a 2 weeks timescale). The variability of the soft excess cannot be associated with variability of a cold absorber as we have shown in Fig.~\\ref{Fig1} and Section 3 that the absorbing column density deduced from spectral modeling was roughly constant over all the observations. \\\\ Note that the soft excess component measured here should not be associated with direct thermal emission from the accretion disk because its temperature is much too high. \\\\ \\\\\\noindent {\\bf Thermal comptonisation.} The historic variations of the photon index can be explained in the framework of thermal comptonisation. \\object{Mrk~205} showed large $2-10$ keV flux variations between the 1997 (high flux state), 2000 (low flux state), 2001 (high flux state) and 2002 (low flux state) observations which do not seem to be correlated with variability of the soft excess. \\\\ However, in thermal comptonisation processes, the index of the comptonised component is inversely correlated with the ratio of the photon seed flux to the comptonised power law flux. We can test this hypothesis if we assume that the soft excess strength is proportional to the strength of the soft photon field (Walter \\& Fink \\cite{walterfink}; Walter \\& Courvoisier \\cite{walter1}). The index for 2000 was $1.73\\pm 0.02$ for a ratio of 23\\% while for 2001 the index was $2.00\\pm 0.03$ for a ratio of 3\\%. This is thus in disagreement with our expectations. \\\\ On the other hand, using the {\\sc compha} model to estimate the UV flux, we found that an increase of the photon index goes with a larger UV flux, in agreement with simple thermal comptonization model expectation. Similar conclusions were found by Petrucci et al. (\\cite{pet00}) for the Seyfert~1 galaxy \\object{NGC~5548}. The contradiction between the soft X-ray and UV variations is not very surprising. Indeed the soft excess component is certainly linked to the UV bump through complex radiative transfer (e.g. comptonization in the warm (few keV) upper layer of the accretion disk). Consequently, its expected spectral behavior is not straightforward and may be at odds from the zeroth order expectations. \\\\ \\\\\\noindent {\\bf Constraints on the reflecting material.} The nature of the reprocessing medium can be the accretion disk or some obscuring, distant material. If the Fe K$\\alpha$ line is produced in the inner parts of the accretion disk, its profile should be considerably broader, with very distinct asymmetries due to Doppler and gravitational redshift. This is not possible to test with the {\\it BeppoSAX} data while Page et al. (\\cite{page}) report that they did not find evidence for it in the {\\it XMM-Newton} data either.\\\\ The strength of the Compton reflection $R$ is here systematically higher than what was found in Seyfert galaxies ($R\\simeq 0.8\\pm 0.2$, Nandra \\& Pounds \\cite{nandra}), while on the contrary, the iron line equivalent width EW is smaller than the mean EW value. It is worth noting however that our measures of $R$ are consistent with $R<1$ at 90\\% confidence, see Fig.~\\ref{Fig10}.\\\\ The inconsistency between the best fit values of $R$ and the line EW may then occur by chance and the reflecting material in \\object{Mrk~205} may simply cover a solid angle $<2\\pi$ as seen from the X-ray source. The line energy then also suggests the material to be neutral. A low Fe abundance will make the line weaker and the Compton reflection hump stronger (as observed) by reducing the opacity above the Fe K edge (George \\& Fabian \\cite{geofab}).\\\\ Given these constraints a simple solution could be that the reprocessing is done in a distant material e.g. the dust torus.\\\\ \\\\ However, we cannot rule out the presence of close mildly ionized material. Indeed Nayakshin \\& Kallman (\\cite{naykal}) have shown that in the case of relatively strong illumination by local X-ray flares, the upper layer of the reflecting material may become strongly ionized. In this case the reflection features are mainly produced by the neutral depth layer and the expected line energy is 6.4 keV as we observed. The line EW is however smaller than the one expected in the case of neutral reflector since the line may be partly comptonized in the hot skin. But clearly this skin cannot be highly ionized since in this case it would act like a perfect mirror and the reflection hump would be suppressed contrary to what we observe. Lines in the soft band are also expected in the case of a mildly ionized medium and it may explain the poor fit of the soft excess with a simple black body component. \\noindent \\\\ \\\\ {\\bf The 1 keV line.} The spectrum of January 2002 contains a narrow emission feature centered at an energy of $\\sim$ 1 keV which was also found in at least 3 other AGN. This feature could be explained by a blend of several emission lines, predominantly from ionized species of Ne and possibly from Fe L shell. The line emission could be a result of reflection in photoionized matter. It could also be resulting from recombination in an optically thin thermal plasma."
    },
    "0403/astro-ph0403200_arXiv.txt": {
        "abstract": "We consider the general conditions of quark droplets formation in high density neutron matter. The growth of the quark bubble (assumed to contain a sufficiently large number of particles) can be described by means of a Fokker-Planck equation. The dynamics of the nucleation essentially depends on the physical properties of the medium it takes place. The conditions for quark bubble formation are analyzed within the frameworks of both dissipative and non-dissipative (with zero bulk and shear viscosity coefficients) approaches. The conversion time of the neutron star to a quark star is obtained as a function of the equation of state of the neutron matter and of the microscopic parameters of the quark nuclei. As an application of the obtained formalism we analyze the first order phase transition from neutron matter to quark matter in rapidly rotating neutron stars cores, triggered by the gravitational energy released during the spinning down of the neutron star. The endothermic conversion process, via gravitational energy absorption, could take place, in a very short time interval, of the order of few tens seconds, in a class of dense compact objects, with very high magnetic fields, called magnetars. ",
        "introduction": "Since \\citet{Wi84}, following early proposals by \\citet{It70} and \\citet{Bo71}, suggested that strange quark matter, consisting of $% u$-, $d$- and $s$-quarks is energetically the most favorable state of the matter, the problem of existence of strange quark stars has been intensively investigated in the physical and astrophysical literature. The possibility that some compact objects could be strange stars remains an interesting and intriguing, but still open question. \\citet{Wi84} also proposed two ways of formation of strange matter: the quark-hadron phase transition in the early universe and conversion of neutron stars into strange ones at ultrahigh densities. In the theories of strong interaction quark bag models suppose that breaking of physical vacuum takes place inside hadrons. As a result vacuum energy densities inside and outside a hadron become essentially different, and the vacuum pressure on the bag wall equilibrates the pressure of quarks, thus stabilizing the system. If the hypothesis of the quark matter is true, then some of neutrons stars could actually be strange stars, built entirely of strange matter \\citep{AlFaOl86,HaZdSch86}). However, there are general arguments against the existence of strange stars \\citep{CaFr91}. For an extensive review of strange star properties see \\citet{ChDaLu98}.  Several mechanisms have been proposed for the formation of quark stars. Quark stars are expected to form during the collapse of the core of a massive star after the supernova explosion, as a result of a first or second order phase transition, resulting in deconfined quark matter \\citep{Da}. The proto-neutron star core, or the neutron star core, is a favorable environment for the conversion of ordinary matter to strange quark matter \\citep{ChDa}. Another possibility is that some neutron stars in low-mass X-ray binaries can accrete sufficient mass to undergo a phase transition to become strange stars \\citep{Ch96}. This mechanism has also been proposed as a source of radiation emission for cosmological $\\gamma $-ray bursts \\citep{Ch98a}. Some basic properties of strange stars like mass, radius, cooling, collapse and surface radiation have been also studied \\citep{ChHa,ChHa03}.  The physical mechanisms of the transition from neutron matter to quark matter in an astrophysical background have been studied within several models. The first is due to \\citet{Ol87}, who used a non-relativistic diffusion model. As such, this is a slow combustion model, with the burning front propagating at a speed of approximately $10$ m/$\\sec $. This is determined primarily by the rate at which one of the down quarks inside the neutrons is converted, through a weak decay, to a strange quark: $% d+u\\rightarrow s+u$. The second method of describing the conversion process was first suggested by \\citet{HoBe88}, and analyzed in detail by \\citet{LuBeVu94} and \\citet{LuBe95}, who modelled the conversion as a detonation. In this case the conversion rate is several orders of magnitude faster than that predicted by the slow combustion model. This model is based on the relativistic shock waves and combustion theory. But regardless of the way in which the transformation occurs, an initial seed of quark matter is needed to start the process.  Real neutron stars have two conserved charges- electric and baryonic. Therefore, a neutron star has more than one independent component and in this sense it is a ``complex\" system \\citep{Gl02}. The characteristics of a first-order phase transition, like the deconfinement transition, are very different in the two cases. In the ``complex\" system the conserved charges can be shared by the two phases in equilibrium in different concentrations in each phase. The mixed phase, formed from hadrons and quarks cannot exist in a simple body in the presence of the gravity because the pressure in that phase is a constant \\citep{Gl00}. This causes a discontinuity in the density distribution in the star occurring at the radius where Gibbs criteria are satisfied. The isospin symmetry energy in neutron rich matter will exploit the degree of freedom of readjusting the charges between hadronic and quark phases in equilibrium so as to reduce the symmetry energy to the extent consistent with charge conservation. Regions of hadronic matter will have a net positive charge neutralized by a net negative charge on the quark matter regions \\citep{Gl00}. Coulomb repulsion will prevent the regions of like charge to grow too large, and the surface energy will act in the opposite sense in preventing them to become too small. Consequently, the mixed phase will form a Coulomb lattice so as to minimize the sum of Coulomb and surface interface energy at each proportion of phases. As the quark phase becomes more abundant, the droplets merge to form rods, rods merge to form slabs etc. \\citep{Gl00}. Hence the actual geometric phase of quark nuggets in a neutron star evolves as a function of the pressure and of the surface energy between dense nuclear matter and quark matter.  In order to describe the process of formation and evolution of quark nuclei in the neutron stars one must use nucleation theory. The goal of nucleation theory is to compute the probability that a bubble or droplet of the $A$ phase appears in a system initially in the $B$ phase near the critical temperature \\citep{LaLi80}. Homogeneous nucleation theory applies when the system is pure. Nucleation theory is applicable for first-order phase transitions when the matter is not dramatically supercooled or superheated. If substantial supercooling or superheating is present, or if the phase transition is second order, then the relevant dynamics is spinodal decomposition \\citep{ShMo01}.  The nucleation of strange quark matter inside hot, dense nuclear matter was investigated, with the use of Zel'dovich's kinetic theory of nucleation, by \\citet{HoBeVu92}. By assuming that the newly formed strange quark matter bubble can be described by a simple bag model containing $N_{q}$ ultrarelativistic quarks, and by assuming that the time evolution of the radius of the bubble is described by an equation of the form $dr/dt=(r-R_{c})/\\tau _{w}$, where $R_{c}$ is the critical radius of the bubble and $\\tau _{w}$ the weak interaction time scale, the nucleation rate $\\xi $ is given by \\begin{equation} \\xi =2.2\\times 10^{-2}\\left( T/\\sigma \\right) ^{1/2}N_{qc}^{3/4}\\tau _{w}^{-1}\\exp \\left( -3.1N_{qc}^{1/2}\\sigma /T\\right) , \\end{equation} where $N_{qc}\\leq 300$ is the critical quark number, $\\sigma $ is the surface tension of the bubble and $n_{n}$ and $n_{q}$ are the particle number densities in the neutron and quark phase, respectively \\citep{HoBeVu92}% . In this way a lower bound for the temperature of the nucleation, $T\\geq 2.1MeV$, can be obtained. The effects of the curvature energy term on thermal strange quark matter nucleation have been considered by \\citet{Ho94}, who derived, within the same approach, the following expression for the nucleation rate: \\begin{equation} \\xi =\\frac{r_{c}^{2}n_{q}n_{n}}{4\\pi \\tau _{w}}\\left( T/\\sigma \\right) ^{1/2}\\exp \\left( -4\\pi R_{c}^{2}/3T\\right) . \\end{equation}  All the effects of the curvature enter only through the value of $R_{c}$. As a general conclusion one obtains the result that even the curvature term acts against strange quark matter nucleation, the physical temperature of a just born proto-neutron star is, in any model, more than enough to drive an efficient boiling of the neutron material (the observations of the neutrino flux from SN1987A are consistent with an effective temperature of $T\\sim 4$ MeV). The possibility of stable strange quark matter (both bulk and quasi-bulk) at finite temperature, and some of its properties have been investigated, within the framework of the dynamical density-dependent quark mass model of confinement, by \\citet{Ch93}. The behavior of the surface tension and the stability of quark droplets at $T\\neq 0$ has also been discussed, with reference to strangelet formation in ultra-relativistic heavy-ion collisions.  A different approach to the problem of thermal nucleation was taken by \\citet{OlMa94}. They used the assumption (standard in the theory of bubble nucleation in first-order phase transitions) that bubbles form at a rate given by $R\\approx T^{4}\\exp \\left( -W_{c}/T\\right) $, with $% W_{c}$ the minimum work required to form a bubble with radius $r_{c}$. As a main result it follows that if the bag constant lies in the interval where three-flavor but not two-flavor quark matter is stable at zero pressure and temperature ($145$ MeV$\\leq B^{1/4}\\leq 163$ MeV), then all or parts of a neutron star will be converted into strange matter, during the first second of its existence. For bag constant above the stability interval, a partial transformation is still possible \\citep{OlMa94}.  \\citet{He95} calculated the rate of formation of quark matter droplets in neutron stars from a combination of bubble formation rates in cold degenerate and high temperature matter, taking into account nuclear matter calculations of the viscosity and thermal conductivity. The droplet formation rate is that of \\citet{LaTu73}, given by $% I=\\left( k/2\\pi \\right) \\Omega _{0}\\exp \\left( -W_{c}/T\\right) $, with $% \\Omega _{0}=\\left( 2/3\\sqrt{3}\\right) \\left( \\sigma /T\\right) ^{3/2}\\left( r_{c}/\\xi _{q}\\right) ^{4}$ the statistical prefactor, which measures the phase-space volume of the saddle point around $R_{c}$ that the droplet has to pass, in its way to the lower energy state. $\\xi _{q}$ is the quark correlation length. The dynamical prefactor determines the droplet growth rate and is given by $k=\\left( 2\\sigma /\\Delta w^{2}r_{c}^{3}\\right) \\left( \\lambda T+2\\left( 4\\eta /3+\\zeta \\right) \\right) $, where $\\Delta w$ is the enthalpy difference and $\\lambda ,\\eta $ and $\\zeta $ are the thermal conductivity and the shear and bulk viscosities, respectively. The droplet formation rate can be expressed as \\citep{He95} \\begin{equation}\\label{eqn} I=\\frac{\\sigma _{20}^{7/2}\\mu _{400}^{2}\\eta _{50}}{\\Delta P_{10}\\Delta w_{10}^{2}T_{10}^{3/2}}\\exp \\left( 185-134\\frac{\\sigma _{20}^{3}}{\\Delta P_{10}^{2}T_{10}}\\right) \\textrm{ s}^{-1}\\textrm{km}^{-3}, \\end{equation} with $\\mu _{400}$ is the quark chemical potential in units of $400$ MeV.  The total number $N$ of droplets formed by nucleation is the integrated rate over the volume of the neutron star and the time after the nucleation process starts at the moment $t_0$, $N=\\int_{0}^{R}4\\pi r^{2}dr\\int_{t_{0}}^{\\infty }I\\left[ \\Delta p\\left( r\\right) ,T\\left( t\\right) \\right]dt$, where $R$ is the radius of the neutron star. The pressure and temperature depend sensitively on the equation of state of the nuclear and quark matter. To convert the core of the neutron star into quark matter at least one droplet must be formed, i.e., $N>1$. Then, with the use of Eq. (\\ref{eqn}), the condition for the formation of at least one droplet requires $\\sigma \\leq 24$ MeV fm$^{-2}\\Delta P_{c,10}^{2/3}T_{c,10}^{1/3}$, where $P_{c}$ and $T_{c}$ are the core values of the pressure and temperature, respectively \\citep{He95}. The pressure and enthalpy difference depends strongly on the equation of state of nuclear and quark matter and hence the droplet formation rate cannot be reliably estimated.  The effect of subcritical hadron bubbles on an inhomogeneous first order quark-hadron phase transition was studied in \\citet{ShMoGuGl00}. The transition from nuclear matter (consisting of neutrons, protons and electrons) to matter containing strangeness was considered, within a mean field type model and with the use of Langer nucleation theory, by \\citet{No02}. An estimate of the time the new phase to appear at various densities and times in the cooling history of a protoneutron star was also given.  On the other hand, as a result of an increase of the central density of the star, due for example to accretion or spinning down, a metastable, super-compressed neutron phase can appear, with a star consisting of the new phase surrounded by normal neutron matter \\citep{Gr98}. This situation is similar to a gas undergoing a phase transition to its liquid state if compressed to a volume $V_{c}$ at fixed $T$. If it is slowly compressed, it may stay in the vapor state even for $V<V_{c}$ (similar examples are liquid water freed of dissolved air which may be heated above boiling temperature without the formation of vapor and cooled below freezing temperature without solidification). The transition from the super-compressed neutron phase to the quark phase can take place in a very long time, which in some cases can be longer than the estimated life of the neutron star \\citep{Gr98}. Therefore an increase in the density of the neutron star does not automatically lead to the transition to quark matter inside the star.  The effect of the magnetic field on the quark star structure and on the nucleation process of the quark bubbles has been studied in \\citep{Ch95,ChSa96,Ch96a}. In the presence of strong magnetic fields the equation of state of strange quark matter changes significantly. The strange stars become more compact, the magnetic field reducing the mass and radius of the star. The surface energy and the curvature term of the quark phase both diverge, with the surface tension diverging logarithmically, while the curvature term diverges much faster. Therefore the thermal nucleation of quark bubbles in a compact metastable state of neutron matter is completely forbidden in the presence of a strong magnetic field.  These results for the formation of quark bubbles in neutron matter have been obtained by using the Csernai and Kapusta theory of nucleation \\citep{CsKa92,CsKa92b} and its extension including the thermal conductivity of dense matter \\citep{VeVi94}. In this theory both the hadron and quark materials are considered as substances with zero baryonic number, and are treated by using a relativistic formalism. The principal result in this approach is the suggestion that the prefactor $k$ is proportional to the transport coefficients (viscosity and thermal conductivity) of the neutron matter, and thus when these coefficients vanish a quark bubble necessarily does not form.  However, \\citet{RuFr96} argued that the energy flow does not vanish in absence of any heat conduction or viscous damping. Since the change of energy density $e$ in time is given, in the low velocity limit, by the conservation equation $\\partial e/\\partial t=-\\nabla \\left( w\\vec{v}% \\right) $ \\citep{CsKa92}, where $w$ is the enthalpy and $\\vec{v}$ the velocity, this implies that the energy flow $\\sim w\\vec{v}$ is always present. Therefore, an expression for the prefactor can be derived, which does not vanish in the absence of viscosity. The viscous effects cause only a small perturbation to the prefactor. The differences between the Csernai-Kapusta (CK) \\citep{CsKa92} and Ruggieri-Friedman (RF) \\citep{RuFr96} results are due to the technical differences in the treatment of the pressure gradients.  A generalized approach, following the CK formalism leading to the prefactor in both viscous and non-viscous regimes was developed by \\citet{ShMoGu01}. Unlike in \\citet{CsKa92}, the linearized relativistic hydrodynamics equations have been solved in all regions.  It is the purpose of the present paper to extend the Zel'dovich nucleation theory for the formation of quark droplets in neutron matter, by taking into account both the effects of the energy flow and of thermal conductivity, shear and bulk viscosities. As a result a more general expression for the quark matter droplet rate formation can be obtained, which is also valid in the limiting case of vanishing conductivity and viscosity coefficients. In the nucleation theory of \\citet{CsKa92}, the transition from the neutron to quark state is possible only for a normal matter having viscous properties. For zero bulk and shear viscosities the transition is impossible. This assumption is, however, to restrictive, and the transition can also take place for a perfect (no viscosity or heat conduction) neutron matter fluid. We also derive the expression of the rate formation for this case.  An alternative point of view of the transformation of neutron matter into strange matter can be developed if one assumes that the conversion process is endothermic. In this case the strange quark matter formation inside a neutron star is triggered by an external source of energy (accretion from the companion star or spin-down). Pulsars are born with an enormous angular momentum and rotational energy, which they radiate over a long period of time via electromagnetic radiation and electron-positron pair emission. When rotating rapidly, a pulsar is centrifugally flattened. With decreasing angular velocity the central density of the star is increasing and may attain the critical density necessary for a phase transition. First at the center and then in an expanding region the neutron matter will be converted to a highly compressible quark matter phase \\citep{GlPeWe97}. The conversion of neutron matter to quark matter alters the moment of inertia of the star, and the epoch over which conversion takes place will be signaled in the spin-down characteristics of the pulsar. A measurable quantity, the braking index $\\Omega \\ddot{\\Omega}/\\dot{\\Omega}^{2}$, where $\\Omega $ is the angular velocity of the star, could be an observational indicator of the slow transition from neutron to quark phase \\citep{GlPeWe97}. By using the formalism developed for the nucleation of quark matter we reconsider the idea of phase transition from neutron matter to quark matter in rotating compact stellar objects, by assuming that the nucleation and the growth of the quark droplets and the formation of a quark core in a neutron star is a result of the transfer of the gravitational energy to the quark droplets, due to the spinning down of the rapidly rotating neutron star.  The present paper is organized as follows. The nucleation kinetics of strange matter bubbles is considered, within the framework of the Zel'dovich nucleation theory, in Section II. The expressions for the rate of formation of quark matter are derived, in both the perfect and dissipative neutron fluid case, in Section III. The transition of a neutron star to a quark star, due to increase of the central density as a result of the spinning down is considered in Section IV. In Section V we discuss and conclude our results. ",
        "conclusions": "In the present paper we have considered a kinetic theory of strange matter nucleation in neutron stars. We derived general necessary conditions for neutron to quark matter conversion. There have been several assumptions made in the derivation of the results, mostly related to the form of the growth speed of the bubble. The result is valid only if non-linear effects can be ignored, and the linearized hydrodynamic equations are applicable. Furthermore, we have assumed that the radii of the bubbles are larger than the correlation length and that heating due to dissipation is slow, causing the temperature to vary slowly across the bubble wall. We also considered that the phase transition is strongly first order, releasing considerably heat.  The growth of the critical size quark bubbles nucleated in the first order transition in the neutron star core is governed by the dynamical prefactor $% k $. There are several expressions for the prefactor obtained in the physical literature. In the Csernai-Kapusta approach \\citep{CsKa92}, the prefactor essentially depends on the dissipative properties (shear, bulk viscosity and thermal conductivity) of neutron matter. Therefore, in a non-viscous medium with $\\xi _{n}=\\eta _{n}=\\lambda _{n}=0$, the growth of the quark bubbles cannot take place. The main physical assumption is that the energy flow is provided by the viscous effects only.  On the contrary, one can argue that the energy flow does not vanish in absence of heat conduction or viscous damping \\citep{RuFr96,ShMoGu01}% . In this case, and in the limit of zero viscosity, the prefactor depends only on two scale parameters, the correlation length $\\zeta $ and the critical radius of the quark bubble $R_{c}$. For a viscous medium the prefactor can simply be written as the sum of the viscous and non-viscous term, with the viscous term not affecting the growth process significantly. The previous (viscous) results can be reobtained by using some assumptions for the velocity of sound in the medium around the saddle configuration.  The condition for the formation of at least quark bubble in the neutron matter as a result of the thermal fluctuation is given by Eq.~(\\ref{cond1}). The condition includes both the cases of viscous and ideal neutron matter. It is extremely sensitive with respect to the numerical values of the microscopic parameters characterizing the quark bubble, like the surface tension $\\sigma $, the bag constant $B$ or the curvature coefficient $\\gamma $. In discussing the astrophysical implications of the expression (\\ref {cond1}) we shall consider separately the two limiting cases, corresponding to viscous and non-viscous neutron matter, respectively.  In the case of ideal neutron matter, the condition of the formation of a quark bubble is given by Eq.~(\\ref{cond2}). This equation determines at which temperature the phase transition takes place. For example, in the case of the Walecka mean field equation of state \\citep{Wa75}, with $K=15758$ and $% \\gamma =4.95$ \\citep{LuBeVu94}, the temperature $T_{c}$ necessary for the formation of a stable quark bubble in a second ($\\Delta t=1$ s), in a volume $\\Delta V_{0}$ with radius $\\Delta R_{0}=1$ km ($\\Delta V_{0}=4\\pi /3$ km$% ^{3} $), and at the center of the neutron star, is around $T_{c}=12-13$ MeV, for $\\sigma =(85$ MeV$)^{3}$. Here we have also assumed the standard values for the quark matter chemical potential and for the bag constant, $\\mu _{u}=\\mu _{d}=280$ MeV and $B=60$ MeVfm$^{-3}$ \\citep{ChDaLu98}. Since the temperature of a newborn neutron star is around $10$ MeV \\citep{OlMa94}, this value of the surface tension seems to rule out, for the given equation of state, the possibility of formation of quark nuclei in the neutron matter. However, a small variation in the value of $\\sigma $, $\\sigma =(75$ MeV$% )^{3} $ reduces the temperature for quark nuclei formation to $T_{c}=4-6$ MeV, a value which does not exclude the possibility that quark nuclei formation can be initiated in the early stages of the neutron star evolution. Smaller values of $\\sigma $, of the order of $\\sigma =(50$ MeV$% )^{3}$, can lower even more the phase transition temperature, making it possible even for cold neutron stars. The transition temperature is relatively insensitive to the details of the equation of state of the dense neutron matter.  If the temperature at which the phase transition is initiated is strongly dependent by the microscopic model adopted for describing the neutron matter properties, the total transition time of the neutron star to a quark star, given by Eq.~(\\ref{conv}), is much less dependent on the details of the equation of state or the exact numerical values of the surface tension or curvature of the bubble. Generally, for a neutron star with radius $R=12-15$ km, the conversion time is very long, of the order of $10^{6}$ years for $% \\sigma \\approx (75$ MeV$)^{3}$. For $\\sigma \\approx (50$ MeV$)^{3}$ or even lower, the conversion time could have values of the order of $10^{7}$ years, which are still of the same order of magnitude as the lifetime of a neutron star ($10^{7}$ years). Hence, in the present approach, which is limited to the consideration of linear effects only in the hydrodynamic description of the first order phase transitions for an ideal neutron fluid, the growth of the quark nuclei is very slow. Therefore the results of the present analysis suggest that even quark nuclei could appear inside the neutron star at a very early stage, the transition to a pure quark phase ends only in the final stage of its evolution.  If the neutron matter-quark matter phase transition is mainly driven by viscous processes, like in the Csernai-Kapusta scenario, then the time necessary for a neutron star to convert to a quark star follows from Eqs.~(\\ref{k1}) and (\\ref{tconv}). In order to numerically estimate the conversion time, we need an estimation of viscosity and conduction coefficients of nuclear matter. We shall also consider two limiting cases for the phase transition. Below the critical density $\\rho =\\rho _{c}\\approx 10^{14}$ g/cm$% ^{3}$, cold neutron matter consists of two distinct phases. At low density, neutron star matter consists of nuclei which form a solid lattice for temperatures below the melting temperature and a sea of relativistic electrons. When the density reaches $\\rho \\approx 10^{11}$ g/cm$^{3}$, neutrons begin to ''drip'' from the neutron-rich nuclei and then, in addition to the relativistic electrons there is a sea of non-relativistic neutrons. At the critical density the neutron-rich nuclei dissolve, leaving seas of degenerate neutrons, protons and electrons. The neutrons in cold neutron star matter are quite likely superfluid, and therefore the transport coefficients are dominated by electrons, with the motion of electron determined by electron-electron and electron-proton scattering \\citep{FlIt79}% As a result the viscosity is much lower, but the thermal conductivity is only slightly smaller as compared to the case in which both neutrons and protons form normal fluids \\citep{FlIt79}. Hence, we can neglect the viscosity coefficients with respect to the thermal conductivity and approximate $\\lambda _{n}$ by $\\lambda _{n}\\approx \\lambda _{0}\\rho _{14}/T_{8}$, with $\\lambda _{0}\\approx 10^{23}$ \\citep{FlIt79}. A slightly different expression for the leading term in thermal conductivity has been proposed in \\citet{Da84}, $\\lambda _{n}T\\approx \\lambda _{0}\\left( n/n_{0}\\right) ^{1/4}$, with $\\lambda _{0}=0.15$ fm$^{-3}$ and $n_{0}=0.145$ fm$^{-3}$. Hence, by using this last expression, we obtain for $t_{conv}$ \\begin{equation} t_{conv}\\approx R_{c}\\frac{\\left( \\Delta w\\right) ^{2}}{4\\sigma \\lambda _{0}\\left( n/n_{0}\\right) ^{1/4}}R^{2}.  \\label{therm} \\end{equation}  The transition time is independent of the temperature of the neutron matter. It depends only on the microscopic properties of the quark bubble (surface tension and curvature coefficient), density and the equation of state of the matter inside the star. For $\\mu _{u}=\\mu _{d}=280$ MeV and $B=60$ MeVfm$% ^{-3}$ and for $\\sigma \\approx (75$ MeV$)^{3}$ Eq.~(\\ref{therm}) gives, for a neutron star with radius $R=12$ km and for several equations of state values of the conversion time of the order $t_{conv}\\approx 10^{11}$ years.  In the case of young and hot neutron stars, the main dissipative mechanism which could drive the phase transition is the shear viscosity of the neutron matter. The temperature dependence of the shear viscosity coefficient can be approximated by $\\eta _{n}\\approx \\left( \\eta _{0}/T^{2}\\right) \\left( n/n_{0}\\right) ^{2}$, with $\\eta _{0}=1700$ MeV$^{3}$/fm$^{2}$ \\citep{Da84}. Consequently, by neglecting the heat conduction and the bulk viscosity, the conversion time becomes \\begin{equation} t_{conv}\\approx \\frac{3}{32}R_{c}\\frac{\\left( \\Delta w\\right) ^{2}}{\\sigma \\eta _{n}}R^{2}\\approx \\frac{3}{32}R_{c}\\frac{\\left( \\Delta w\\right) ^{2}}{% \\sigma \\eta _{0}\\left( n/n_{0}\\right) ^{2}}T^{2}R^{2}.  \\label{shear} \\end{equation}  In small temperatures, the shear viscosity becomes very large. Hence, for a shear viscosity driven phase transition the conversion time could be very small.  In Fig.~1 we have represented $t_{conv}$ given by Eq.~(\\ref{shear}) as a function of temperature, for different equations of state of neutron matter. We have considered four equations of state, namely, the equation of state (EOS) of the free neutron gas \\citep{ShTe83}, the Bethe-Johnson EOS \\citep{ShTe83}, the Lattimer-Ravenhall EOS \\citep{LaRa78} and the Walecka EOS \\citep{Wa75}. In all cases we calculated the conversion time for a neutron star with radius $R=10$ km.  \\begin{figure} \\plotone{f1.eps} \\caption{Variation of the conversion time $t_{conv}$ for a neutron star with radius $R=10$ km, for different equations of state of the neutron matter: free neutrons EOS (solid curve), Bethe-Johnson EOS (dotted curve), Lattimer-Ravenhall EOS (short-dashed curve) and Walecka EOS (long-dashed curve). \\label{FIG1}} \\end{figure}  As one can see from the figure, in this case the conversion time is much shorter as compared to the ideal neutron matter case or thermal conduction driven phase transition. $t_{conv}$ is rapidly decreasing with the temperature. For a neutron star with rapidly cooling processes, e.g. direct URCA process, kaon condensation etc., the stellar temperature can decrease to less than $10^{-2}$MeV in a time scale of $10^{-4}$ yr \\citep{Ts98}. According to Fig. 1, the conversion can take place in a time scale less than $10^{-3}$ yr if the spin-down time scale is less than this time scale. This can be case if this neutron star is a magnetar. We should point out that the interior of magnetars are not quantum liquid, i.e. superfluid or superconducting, because the strong magnetic field can destroy the nucleon cooper pairs \\citep{Ch92}.  On the other hand, if the neutrons and protons are superfluid, then the transport properties are determined completely by electron-electron scattering. In fact in this case the viscosity of the nuclear fluid is given by $\\eta ^{-1}=\\left( 15\\pi ^{4}\\alpha ^{2}/2p_{e}^{5}\\right) \\left( 2p_{e}/k_{ft}\\right) \\left[ 5/2+3\\left( m_{e}/p_{e}\\right) ^{2}+\\left( m_{e}/p_{e}\\right) ^{4}\\right] T^{2}$ \\citep{FlIt76}, where $\\alpha $ is the fine structure constant, $k_{ft}=\\left( 4\\alpha p_{e}\\epsilon _{e}/\\pi \\right) ^{1/2}$ is the Thomas-Fermi wave vector, and $m_e$, $p_{e}$ and $\\epsilon _{e}$ are the electron mass, momentum and energy, respectively.  We assume that the electron gas inside the neutron star is extreme relativistic and strongly degenerate. Then it follows that $p_{e}\\approx p_{F}=\\left( 3n_{e}/8\\pi \\right) ^{1/3}$ \\citep{ShTe83}, with $p_F$ the Fermi momentum and $n_{e}$ the electron number density. Since $n_{e}$ can be related to the total particle number $n$ by means of the relation $n_{e}=Y_{e}n$, with $Y_{e}$ the mean number of electrons per baryon, it follows that $p_{e}\\approx \\left( 3Y_{e}n_{0}/8\\pi \\right) ^{1/3}\\left( n/n_{0}\\right) ^{1/3}$. Hence we obtain for the shear viscosity coefficient of the nuclear matter \\begin{equation} \\eta^{-1}\\approx \\frac{15}{2}\\frac{8^{5/3}\\pi^{37/6}\\alpha ^{3/2}}{\\left( 3Y_{e}n_{0}\\right) ^{5/3}\\left( n/n_{0}\\right) ^{5/3}}\\left\\{ \\frac{5}{2}+3% \\left[ \\frac{m_{e}\\left( 8\\pi \\right) ^{1/3}}{\\left( 3Y_{e}n_{0}\\right) ^{1/3}\\left( n/n_{0}\\right) ^{1/3}}\\right] ^{2}+\\left[ \\frac{m_{e}\\left( 8\\pi \\right) ^{1/3}}{\\left( 3Y_{e}n_{0}\\right) ^{1/3}\\left( n/n_{0}\\right) ^{1/3}}\\right] ^{4}\\right\\} T^{2}. \\end{equation}  The variation of the conversion time, as a function of temperature, is represented, for a superfluid neutron and proton core of the star with $Y_e=0.1$, in Fig.~2.  \\begin{figure} \\plotone{f2.eps} \\caption{Variation of the conversion time $t_{conv}$ for a neutron star with radius $R=10$ km to a quark star, in the case of the presence of a superfluid core, for different equations of state of the neutron matter: free neutrons EOS (solid curve), Bethe-Johnson EOS (dotted curve), Lattimer-Ravenhall EOS (short-dashed curve) and Walecka EOS (long-dashed curve). \\label{FIG2}} \\end{figure}  In this case, since the viscosity of the superfluid nuclear matter is much smaller, the conversion time is longer than for a nuclear fluid dominated by the neutron viscosity. However, in this case the phase transition from neutron matter to quark matter can also take place at very low temperatures. Consequently, a very short $t_{conv}$, of the order of hours or a few days, is also allowed in the framework of this model.  It is widely accepted that at asymptotic densities the ground state of QCD with $m_s=0$ is the color-flavor locked (CFL) phase \\citep{Al98,Ra01}. In this phase, color gauge symmetry is completely broken, as are both chiral symmetry and baryon number. The effective coupling is weak and the low-energy properties can be determined by using methods from the theory of superconductivity. For large gaps $(\\Delta \\sim 100MeV)$, the CFL phase is rigorously electrically neutral, despite the unequal quark masses, and even in the presence of electron chemical potential. The transition from nuclear matter to quark matter via bubble nucleation may be greatly simplified if the transition occurs directly to quark matter in the CFL phase \\citep{Ra01}. For a given baryonic chemical potential $\\bar{\\mu}$, electrically neutral nuclear matter and electrically neutral quark matter have different values of the electronic chemical potential $\\mu _{e}$. Since $\\mu _{e}$ must be continuous across any interface, a mixed phase region is formed, within which positively charged nuclear matter and and negatively charged quark matter with the same $\\mu _{e}$ coexist at any given radius. The growth of the quark droplets is at the expense of the nuclear matter. However, if the quark matter is in the CFL phase, an interface between bulk nuclear matter with non-zero electron number $N_{e}\\neq 0$ and CFL quark matter with $N_{e}=0$ may be stable, as long as $\\mu _{e}$ satisfies the condition $\\left| m_{s}^{2}/4\\bar{\\mu}-\\delta \\mu \\right| <\\Delta /\\sqrt{2}$, with $\\delta \\mu $ the variation of the chemical potential \\citep{Ra01}. Therefore, the existence of the CFL phase supports the idea that in neutron stars quark matter and baryonic matter may be continuously connected. Then the growth of the interface between the neutron and CFL quark phase could be the result of exterior energy absorption, via an endothermic mechanism.  In summary, by using the kinetic nucleation theory we have considered the possibility that the neutron matter-quark matter phase transition could be triggered by the gravitational energy released during the spin-down of a pulsar. This process could take place in a very short time interval, of the order of few tens seconds, in a class of dense compact objects, with very high magnetic fields, called magnetars. However, the percentage of neutron star matter that can be converted into quark matter depends on the initial period, the neutron star mass and on the equation of state of the neutron and quark phases. An important and open issue is how could we determine if such a process is taking place or took already place in some pulsars or magnetars, and what are the observational signatures of the endothermic neutron matter-quark matter transition. In particular, it would be interesting to know if such a process could be a possible gamma-ray burst mechanism, as suggested by \\citet{Be03}. All of these issues will be considered in a subsequent paper \\citep{Ta04}."
    },
    "0403/astro-ph0403036_arXiv.txt": {
        "abstract": "A topologically finite universe, smaller than the observable horizon, will have circles-in-the-sky: pairs of circles around which the temperature fluctuations in the cosmic microwave background are correlated. The circles occur along the intersection of copies of the spherical surface of last scattering. For any observer moving with respect to the microwave background, the circles will be deformed into ovals. The ovals will also be displaced relative to the direction they appear in a comoving frame. The displacement is the larger of the two effects. In a Lorenz boosted frame, the angular displacement of a point on the surface of last scattering relative to the comoving frame is proportional to the velocity. For the Earth's motion, the effect is on the order of $0.14^o$ at the very worst. If we live in a small universe and are looking for an identical copy of a spot in the sky, it may be displaced by as much as $0.14^o$ from where we expect. This can affect all pattern based searches for the topology of the universe. In particular, high-resolution searches for circle pairs could be off by this much. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403346_arXiv.txt": {
        "abstract": "Large scale strong magnetic fields in galaxies are generally thought to have been generated by a mean field dynamo. In order to have generated the fields observed, the dynamo would have had to have operated for a sufficiently long period of time. However, magnetic fields of similar intensities and scales to the one in our galaxy, are observed in high redshift galaxies, where a mean field dynamo would not have had time to produce the observed fields. Instead of a mean field dynamo, we study the emergence of strong large scale magnetic fields in the first objects formed in the universe due to the action of a turbulent, helical stochastic dynamo, for redshifts $5 \\leq z \\leq 10$. Ambipolar drift plays an important role in this process due to the low level of ionization of the gas, allowing a large scale stochastic dynamo to operate. We take into account the uncertainties in the physics of high redshift objects by examining a range of values for the parameters that characterize the turbulent plasma. By numerically integrating the nonlinear evolution equations for magnetic field correlations, we show that for reasonable values of the parameters in the time interval considered, fields can grow to high intensities ($\\sim 10^{-6}$ G), with large coherence lengths ($\\sim 2-6$ kpc), essentially independent of the initial values of the magnetic field. ",
        "introduction": "Magnetic fields have been observed in all known structures of our universe, from the Earth to superclusters of galaxies, spanning a wide range of intensities from $\\sim\\mu\\,\\rm{G}$ in galaxies and galaxy clusters to $\\sim 10^{\\,12}\\,\\rm{G}$ in neutron stars. The origin of these fields in large structures, such as galaxies and clusters of galaxies, remains an unsolved problem. Many physical processes have been proposed to explain the origin and evolution of these fields \\citep{prim-b,widrow}. The processes suggested can be divided into two main classes: 1) cosmological mechanisms; and 2) local astrophysical processes. Until now, none of them has provided a satisfactory explanation for the generation of the magnetic fields. \\par A mean field dynamo is commonly invoked to explain the fields observed in our galaxy and in small redshift galaxies (e.g., \\citealp{zeldovich83,moffat78}). In order to have attained the observed intensities, the dynamo would have had to have operated for a time on the order of the age of the universe. However, the presence of equally intense and coherent fields in high redshift galaxies \\citep{carrilli02}, where the mean field dynamo would not have had enough time to amplify the field to the observed values, casts doubt on the mean field dynamo paradigm as the preferred generation mechanism. Fields of similar intensity and coherence to those in the Milky Way, have also been detected in high redshift damped $\\rm{Lyman}-\\alpha $ systems \\citep{oren}. \\par In this paper, we are concerned with the origin of strong, large coherent magnetic fields in high redshift structures. We relate the origin of magnetic fields to the physical conditions in the early universe: 1) a dense low ionized plasma; and 2) appreciable turbulence due to the observed high star formation rate (Lanzetta et al. \\citealp{lanzetta02}). \\par The formation of the first stars and quasars marked the beginning of the transformation of the universe from a smooth initial state to its present lumpy state. In the bottom-up hierarchy of cold dark matter (CDM) cosmologies, the first gaseous clouds collapsed at redshifts $z > 10$ and, subsequently, fragmented into stars due to molecular hydrogen cooling \\citep{loeb01}. These collapsing objects then fragmented into many clumps, which had typical masses of $\\sim 10^{\\,2}-10^{\\,3}M_{\\odot }$. Very massive stars have lifetimes of $\\sim 3\\times 10^{\\,6}\\,\\rm{years}$ and end their lives as supernovae. \\par Recently, Lanzetta et al. \\citep{lanzetta02} showed that the incidence of the highest intensity star formation regions increases monotonically with redshift. Their observations indicate that star formation in the early universe occurred at a much higher rate than was previously believed. Therefore, the rate of occurrence of supernovae would have also been much higher in the past than at the present. Supernovae shocks disturb the plasma in which they are immersed, producing turbulent motions of the gas. If the supernovae rate was much higher in the past than at present, the plasma of the first formed objects must have been much more turbulent than that of presently observed, low redshift, star forming galactic molecular clouds. \\par Turbulence generates stochastic magnetic fields (magnetic noise) at a faster rate than it does mean fields \\citep{kul-and-92}. If the turbulence is strongly non-helical, the fields induced are confined to small scales (Kazantzev \\citealp{kazantzev68}). However, if it is helical, induction of large scale magnetic correlations by the $\\alpha$ effect occurs (Vainshtein \\& Kichatinov \\citealp{vainshtein86}). Astrophysical turbulence is mainly of a helical nature. Hence, we can expect that large scale correlations were induced by the high redshift, turbulent plasma. \\par In this study, we explore the hypothesis that the magnetic fields observed in high redshift galaxies were created by small scale, stochastic, turbulent helical dynamos, rather than by mean field dynamos. We use a simple model of a gas cloud that is assumed to have collapsed at a high redshift $z > 10.$  At $z \\sim 10,$ the cloud would have had a low magnetization level and a high level of turbulence. Thus, it would have been similar to the turbulent, low ionization, star forming molecular clouds observed in our galaxy, albeit with a much smaller initial magnetic field and a much higher turbulence level due to the higher star formation rate in the early universe. \\par It is well known that shock waves produced by supernova explosions accelerate cosmic rays to energies $\\geq\\,\\rm{GeV}$ (e.g., \\citealp{cr-acc}). V\\\"{o}lk et al. \\citep{volk89} showed that in all galaxies, the supernova rate is a direct measure of the cosmic ray intensity. We can, therefore, infer that cosmic rays were already present in considerable intensities in high redshift galaxies. We take into account phenomenologically, the effect on turbulence of cosmic rays, supernova shocks, and powerful stellar winds from massive stars on turbulence by varying the turbulent parameters over a broad range, in order to take into account the uncertainies in our knowledge of high redshift structures. \\par The linear evolution equations for the correlation function of magnetic fields for non-helical turbulence were derived nearly forty years ago by Kazantzev \\citep{kazantzev68}. For helical turbulence, the corresponding equations were obtained twenty years later by Vainshtein \\& Kichatinov \\citep{vainshtein86}. These equations are linear in the magnetic correlations. Recently, Subramanian \\citep{subramanian99} and Brandemburg \\& Subramanian \\citep{bran-sub} derived the non-linear evolution equations for the magnetic correlations by taking into account the back-reaction of the Lorentz force on the plasma charges in the form of ambipolar drift. We solve the nonlinear helical evolution equations numerically for various values of the parameters that characterize the high redshift turbulent plasma. \\par The paper is organized as follows. In section II, we give the  evolution equations for the magnetic correlations. We describe the effects of the main parameters on the integration in section III . Finally, in section IV, we summarize and discuss our results. ",
        "conclusions": "In this work, we studied the problem of the origin of strong coherent large-scale magnetic fields, observed in low-redshift galaxies and previously thought to have been created by the mean field dynamo. Since doubts have been cast on the mean field dynamo as the source of these fields due to their observation in high redshift objects, where the dynamo would not have had suficient time to operate, we investigated here the stochastic helical dynamo as such a source. We showed that these fields  can be generated in the plasmas found in very high redshift objects. \\par The generation of strong large scale coherent fields is a highly non-linear magnetohydrodynamical problem, which depends upon many factors. Here we discussed a possible mechanism for the generation of large scale fields, namely the non-linear evolution and diffusion of magnetic noise. Magnetic noise  becomes coherent on a scale which is larger than that of the turbulence due to the presence of non linear terms in the evolution equations for the magnetic correlations. \\par We found that for realistic turbulent parameters, it is possible to generate the magnetic fields, observed at high redshifts . Between $z\\sim 10$ and $z\\sim 5$, the primordial plasma is strongly turbulent and partially ionized The magnetic field intensities reached are independent of initial correlations. \\par We considered a very simple model for the generation and evolution of the magnetic correlations. The dependence of the resulting magnetic field intensity on the charge composition of the primordial plasma, suggest that the reionization and star formation processes played an important role in determining the features of the magnetic fields detected in high redshift objects. In a forthcoming work we shall address the evolution of the magnetic correlations, considering a time dependent ion density as well as other nonlinear processes, such as the the Hall effect. Other turbulent scenarios, in addition to the homogeneous and isotropic one considered here, will be treated as well."
    },
    "0403/astro-ph0403170_arXiv.txt": {
        "abstract": "{ We study the cooling of isolated neutron stars. The main cooling regulators are introduced: EoS, thermal transport, heat capacity, neutrino and photon emissivity, superfluid nucleon gaps. Neutrino emissivity includes main processes. A strong impact of medium effects  on the cooling rates is demonstrated. With taking into account of medium effects in reaction rates and in nucleon superfluid gaps modern experimental data can be well explained. ",
        "introduction": "\\label{sec:intro} The Einstein Observatory was the first that started the experimental study of surface temperatures of isolated neutron stars (NS). Upper limits for some sources have been found. Then ROSAT offered first detections of surface temperatures. Next $X$-ray data came from Chandra and XMM/Newton. Appropriate references to the modern data can be found in recent works by \\cite{TTTTT02,T04,KYG01,YGKLP03}, devoted to the analysis of the new data. More  upper limits and detections are expected from satellites planned to be sent in the nearest future. In general, the data can be separated in three groups. Some data show very {\\em{``slow cooling''}} of objects, other demonstrate an {\\em{``intermediate cooling''}} and some show very {\\em{``rapid cooling''}}. Now we are at the position to carefully compare the data with existing cooling calculations. ",
        "conclusions": "We have shown that the most up-to-date observed NS cooling data can well be explained. Actually we deal with a many-parametric problem that allows for a variation of many quantities. Besides the data points might be partially shifted from their positions  shown in figures due to a number of uncertainties and assumptions used at their analysis. By above figures we have illustrated different possibilities discriminating more probable explanations from less probable ones. We elaborated the example of the probably most realistic and microscopically supported EoS of the $V18+\\delta v +UIX^*$. Actually we used a simple parameterization of this EoS suggested by \\cite{HJ99} (HHJ). In this EoS the DU process does not show up to $M\\simeq 1.839 M_{\\odot}$. Thus within this model the {\\em ``Standard + DU''} scenario would demonstrate that the majority of experimentally measured cooling points relate to very massive NS, with $M\\geq 1.84 ~M_{\\odot}$. From our point of view such a scenario seems unrealistic.  We exploit in-medium effects in the calculation of the emissivities and the pairing gaps, and, that is less important, in specific heat and heat conductivity. In general, medium effects result in a significant suppression of the superfluid gaps, especially of the $3P_2$ $nn$ gap, and they enhance the cooling  rates of the MMU and the MNB processes through {\\em the pion softening effect with the increase of the density.} Without the latter effect the {\\em ``Standard + PU''} scenario suffers from an internal inconsistency. Pion condensation cannot take place without preliminary softening of the pion mode at lower densities. And v.s., recent argumentation for the pion condensation, cf. \\cite{APR98,SST99}, further motivates the presence of a precursor pion softening, details see in \\cite{MSTV90}. Pion condensation at $n\\gsim 3n_0$ does not contradict the data but the data can be explained also without pion condensation and other so called \"exotics\" (KU, HDU, etc) but with the pion softening. Based on the works \\cite{SVSWW97,BGV01} we have further improved our code. We showed that the results are sensitive to the values and the density dependencies of the nucleon superfluid gaps. {\\em The strong suppression of the $3P_2$ gap} motivated by theoretical evaluations that incorporate medium effects is indeed required for the fit of the data. We show that all three groups of points {\\em ``slow cooling''}, {\\em ``intermediate cooling''} and {\\em ``rapid cooling''} are now well explained on the basis of the {\\em ``Nuclear medium cooling  scenario'' demonstrated here, where an important r\\^ole is played by in-medium effects}, cf. \\cite{V00}.   We may draw the following main conclusions.  i) The normal matter assumption (see Figs. \\ref{fig6} - \\ref{fig10}) seems rather unrealistic, as by itself, as in relation to the cooling data. One could explain the data but at the price of ignoring of medium effects in MMU and MNB. Then the {\\em ``intermediate cooling''} points can only be explained by very low NS masses (as $0.5~M_{\\odot}$) or, together with {\\em ``rapid cooling''} points, by the very high NS masses $M>1.839~M_{\\odot}$, which allow for the DU process. In the latter case the mass window separating the {\\em ``intermediate cooling''} and {\\em ``rapid cooling''} is very narrow. Above price seems us too high and we  drop such a scenario. The superfluidity is called for by the data.  ii) Including superfluid gaps we see, in agreement with recent microscopic findings,  that $3P_2$ neutron gap should be as small as $\\lsim 10$~keV. Otherwise one cannot explain at least several old objects (see Figs. \\ref{fig19}, \\ref{fig19a}). Proton and $1S_0$ neutron gaps also might be suppressed by factors of order of several, as it is demonstrated by Figs. \\ref{fig20}, \\ref{fig21}, \\ref{fig15}, \\ref{fig15a}, but the suppression by factor $\\gsim 10$ is not already permitted.  iii) Medium effects associated with the pion softening are called for by the data. As the result of the pion softening the pion condensation may occur for $n\\geq n_c^{\\rm PU}$ ($n\\geq 3n_0$ in our model). Its appearance does not contradict to the data (see Fig. \\ref{fig16}) but also the data are well described, if the softening effect is rather saturated  (see Fig. \\ref{fig12}) with increase of the density (as demonstrated by curves 1a, 1b in Fig. \\ref{fig1}). At the same time the critical densities for the efficient DU-like exotic processes  (such as PU on charged pion) should not be as small as $<2.5~n_0$, if the given HHJ EoS is indeed correct. Otherwise one would get too rapid cooling of the NS of the typical $1.4~M_{\\odot}$ mass. This also means that the proper DU threshold density can't be too low that puts restrictions on the density dependence of the symmetry energy. Both  statements might be important in the discussion of the heavy ion collision experiments.  iv) We demonstrated a regular mass dependence: for the NS masses $M\\gsim 1.0~M_{\\odot}$ less massive NS cool slower, more massive NS cool faster.  As we have mentioned, for the sake of simplicity we did not include the possibility of the  hyperonization and other possibilities, like kaon condensation and fermion condensation, which may stimulate a more rapid cooling, being working in a line with the pion condensation. We did not include possible quark effects. The latter need a special treatment. The possibility of the color superconductivity in dense NS interiors opens a number interesting possibilities like the so called  two-flavor color superconductivity (2SC) phase, color-flavor-locking (CFL) phase, color-spin-locking (CSL) phase. Their cooling is essentially different. We will return to this discussion in the nearest future."
    },
    "0403/astro-ph0403616_arXiv.txt": {
        "abstract": "{ The currently available values and confidence limits for $\\Omega_{\\mathrm{m0}}$, $H_0$ and globular cluster ages still indicate that the dark energy that dominates the Universe could also be a form of quintessence or phantom energy. In fact, current cosmological values favor phantom energy. To increase the likelihood of a cosmological constant as dark energy instead of phantom energy, the possibilities seem to lie in reducing globular cluster ages, the Hubble constant, or both, and possibly advancing the epoch of globular cluster formation. For a set of possible dark energy equations of state that includes the cosmological constant, quintessence or phantom energy, age--redshift analytical expressions for null curvature universes that include ordinary matter are derived together with the corresponding ages for these universes. ``Cosmic coincidences'' are found for the currently accepted values of $\\Omega_{\\mathrm{m0}}$ when evaluating the age of the Universe with a cosmological constant or a phantom energy as dark energies. ",
        "introduction": "Curvature, age, the Hubble constant and energy density form a long-standing cosmological puzzle. Measurements of the cosmic microwave background radiation (CMBR) anisotropy have confirmed that the Universe is consistent with null, or nearly null, curvature (de Bernardis et al.\\ \\cite{bernardis}), as predicted by the inflation paradigm. However, a Euclidean universe composed of ordinary matter only has an age inconsistent with that derived from globular clusters (Bolte \\& Hogan \\cite{bolte}). Moreover, galaxy cluster data (Carlberg et al.\\ \\cite{carlberg}) allow us to obtain only about 20--30\\% of the amount of matter required for achieving null curvature. This puzzle seemed to be solved by the discovery in 1998 of the first evidence for an accelerating universe based on distant type Ia Supernova (SNeIa) observations (Riess et al.\\ \\cite{riess1}), which have been supported by additional data including more distant SNe (Riess et al.\\ \\cite{riess2}). This acceleration implies the presence of a significant fraction of dark energy with an equation of state of negative pressure. The addition of the required amount of this energy to the observed amount of matter allows us to obtain the total amount of energy required for achieving null curvature. It  also has the desired side effect of producing an older universe that can then be made consistent with independent age measurements.  However, the nature of this dark energy is still to be established. Moreover, its equation of state is still unknown. The first attempts rescued the cosmological constant as an ad hoc explanation, although their interpretation of the constant in terms of a vacuum energy is inconsistent by 124 orders of magnitude with respect to the required value, which leads us to assume that its value has to be zero for consistency. Other energy types, such as quintessence or phantom energy, have also been invoked, although their origin is no clearer than that of the cosmological constant.  Quintessences were introduced by Wetterich (\\cite{wetterich}), Caldwell et al. (\\cite{caldwell1}), Ratra \\& Peebles (\\cite{ratra}) to avoid the extreme fine-tuning needed to allow a cosmological constant to be significant only at recent epochs. Quintessences are characterized by a scalar field whose evolution depends on its potential. The most popular potentials are inverse power  and exponential laws (Kneller \\& Strigari \\cite{kneller} and references therein), although other possibilities have been considered (see, for example, Di Pietro \\& Claeskens\\ \\cite{dpc}). Both inverse power laws and exponential potentials converge to a unique solution for a broad range of different initial conditions.  Phanton energy (Caldwell \\cite{caldwell2}) violates the dominant-energy condition (Hawking \\& Ellis\\ \\cite{he}) that might allow the existence of wormholes. It also makes the Universe reach in a finite time a cosmic doomsday where all objects, from galaxies to nucleons, are ripped apart, a situation that has been termed the ``big rip'' (Caldwell \\cite{caldwell2}). Although some Chaplygin gas model generalization can avoid this big rip (Gonz\\'alez-D\\'\\i az \\cite{gd}), the violation of the dominant energy condition and the difficulties in obtaining a stable phantom model (Carroll et al.\\ \\cite{carroll}) renders this kind of energy more problematic  than quintessences.  In the case of either quintessences, the cosmological constant or phantom energy, it seems established that the introduction of negative pressure energies is required to fit and harmonize the existing CMBR anisotropy, SNeIa, clusters of galaxies, large scale structure, Big Bang nucleosynthesis and age estimator data.  The incorporation of these energies changes the age--redshift relation and estimates of the age of the Universe. The aim of this article is to show the difficulty in avoiding phantom energies with the current age estimator data and to provide analytical age expressions for the homogeneous and isotropic case for the most likely equations of state for our Universe, considered constant, as restricted by the cosmological parameters commonly assumed. These include the cosmological constant, quintessence and phantom energy.  Thomas \\& Kantowski (\\cite{thomas}) provide an age--redshift relation using elliptical integrals for a universe composed of matter and a cosmological constant, but without providing an easy ready-to-use expression. In general, no analytical expressions of age as a function of redshift seems to be readily available in the literature for the currently most widely accepted cosmologies: flat universes of matter plus a cosmological constant, quintessence or phantom energy.  Although numerical integration allows us to obtain ages in a relatively straightforward way (see, for example, Hogg \\cite{hogg}), analytical solutions are faster to evaluate and have the advantage of providing more precise and ready results than numerical approaches. Moreover, analytical expressions provide the explicit dependence on energy densities, thereby easing their study. Finally, they constitute a reference for checking numerical solutions.  In Section 2 the basic equations are given. In Section 3 possible values for the equation of state, considered constant, are reviewed according to recent figures for the cosmological parameters, in Section 4 the general case of a non-constant equation of state is outlined, and in Section 5 analytical age--redshift relations and ages are obtained by solving the equations on a case by case basis for the range of the constant equations of state assumed. ",
        "conclusions": "In the most accepted cosmological model the Universe is currently accelerated owing to a for{\\bf m} of dark energy of unknown origin, as found via distance--redshift fits to high redshift SNeIa. This energy complements the fraction of ordinary matter so that the curvature is null, according to the observed CMBR anisotropy and the inflation paradigm. Also, the dark energy increases the age of the Universe for the same values of $\\Omega_{\\mathrm{m0}}$ and $H_0$ with respect to a universe composed of non-relativistic matter only, thus alleviating possible age conflicts.  In this article a recent determination of $\\Omega_{\\mathrm{m0}}$, $t_0$ and $H_0$ has been selected from the literature with the criterion of providing a set of cosmological parameters as model-independent as possible. A simple analysis shows that the weighted mean equation of state of dark energy has upper values in the range $-0.23 > w > -1.21$. Hence, not only a cosmological constant but also quintessence or phantom energy are also viable dark energy candidates from an observational point of view. In fact, the mean values of $\\Omega_{\\mathrm{m0}}$, $t_0$ and $H_0$ favor phantom energy over other alternatives. Moreover, a time-dependent $w$ favors even more the existence of epochs with $w<-1$. It is qualitatively shown that to reduce the likelihood of phantom energy to favor a cosmological constant as dark energy, smaller globular cluster ages and/or smaller $H_0$ values with respect to those currently determined  are required. Pushing back globular cluster formation to earlier epochs cannot be ruled out, but only together with lower values of $t_{\\mathrm{GC}}$ and/or $H_0$. Otherwise, an age conflict might still be present.  The range of upper $w$ limits inferred is approximately covered by taking $w=-n/3$ for $n=1,2,3,4$. For each of these constant values of $w$, which include quintessence, the cosmological constant and phantom energy, analytical age--redshift solutions for Euclidean universes have been deduced and analytical age expressions derived. Some of these analytical solutions have not been derived before. Furthermore, no analytical expressions seem to be available in the literature for the range of the constant equations of state considered. Analytical equations are more precise, faster and readier to use than numerical calculations when evaluating the age of the Universe. Moreover, the explicit dependence on $\\Omega_{\\mathrm{m0}}$ and $\\Omega_{\\mathrm{w0}}$ renders their study easier. Finally, curious ``cosmic coincidences'' make 1/$H_0$ a good approximation for the age of the Universe for the currently accepted $\\Omega_{\\mathrm{m0}}$ values assuming null curvature and a cosmological constant or a phantom energy as dark energies."
    },
    "0403/astro-ph0403420_arXiv.txt": {
        "abstract": "We present detailed comparisons between high quality observational colour-magnitude diagrams (CMDs) of open star clusters and synthetic CMDs based on MonteCarlo numerical simulations.  The comparisons account for all of the main parameters which determine the shape of the CMD for a stellar population.  For the four clusters studied, NGC 6819, NGC 2099 (M37), NGC 2168 (M35) and NGC 2323 (M50), we derive reddening, distance, age, binary fraction, star formation rate and indicative metallicity by comparing the locations and density of points in the observed CMDs to the simulated CMDs. We estimate the uncertainties related to stellar evolution theories by adopting various sets of stellar models for all of the synthetic CMDs and discuss which stellar models provide the theoretical CMDs that best reproduce the observations. ",
        "introduction": "Theoretical isochrones are commonly fit to the major observational sequences of star clusters in order to both better understand the underlying physics of stellar evolution and to determine properties of the clusters, i.e., the age.  If the metallicity, reddening and distance to the cluster are well constrained from independent techniques, the comparisons typically involve matching the morphology of the turn-off and location of the red giant stars to predictions.  Recently, the newer method of using synthetic colour-magnitude diagrams to compare with observational data has proven to be much more informative and rewarding \\citep{aparicio,tosi,skillman}.  These MonteCarlo simulations allow modelling of several additional parameters which dictate the distribution of points in the CMD, such as stochastic star formation (SF) processes, binary fraction, photometric spread, main-sequence thickness, data incompleteness and small number statistics.  Consequently, the results not only provide a measure of the properties of the cluster, but can also constrain the star formation history (SFH) and the initial mass function (IMF). Furthermore, by comparing the simulations based on several different sets of evolutionary tracks, we can constrain which models use the best prescription of parameters (such as treatment of overshooting, mixing length, etc...).  Confronting the simulations with observations requires a large data set with accurate photometry.  For this, we use the deep $BV$ photometry presented in the CFHT Open Star Cluster Survey (\\cite{kalirai1}, hereafter JSKI).  JSKI observed 19 open star clusters in our Galaxy and have yet published results on the four richest clusters, NGC 6819 (\\cite{kalirai2}, hereafter JSKII), NGC 2099 (\\cite{kalirai3}, hereafter JSKIII), and NGC 2168 and NGC 2323 (\\cite{kalirai4}, hereafter JSKIV).  These data were reduced and calibrated in a homogenous manner as described in JSKI.  The resulting CMDs exhibit very tight main sequences showing several `kinks' and slope changes which are predicted by theory.  More importantly, the combination of very short and deep exposures, and the large aerial coverage of the detector (42$' \\times$ 28$'$) has allowed the measurement of stars from the brightest asymptotic giant branch (AGB) and red giant branch (RGB) phases down to very low-mass main-sequence phases ($\\sim$0.2 M$_\\odot$).  This allows our comparisons to yield evolutionary information over a wide mass range.   The reduced data set in the CFHT Open Star Cluster Survey has been requested by, and made available to, several investigators for additional science rewards outside our goals (e.g., astrometric studies, proper motions, variable stars, radial velocities, brown dwarfs, blue stragglers and Galactic disk star distributions).  The present study complements these efforts and analyses the four published clusters in a way that allows us to include them in a large homogeneous sample of open clusters aimed at studying the formation and evolution of the Galactic disk (Bragaglia 2003, and references therein). Galactic open clusters are indeed particularly well suited to this purpose, since they span a range of ages from a few million to several billion years and can be observed in various regions of the Galactic disk characterised by different star formation histories. They can be used to study both the present day disk structure and its temporal evolution (Janes \\& Phelps 1994, Friel 1995, Tosi 2000, Bellazzini et al. 2003). Old open clusters offer a unique opportunity to trace the whole kinematical and chemical history of our disk, if collected in populous and representative samples and accurately and homogeneously analysed (see e.g., Twarog, Ashman, \\& Anthony-Twarog 1997; Carraro, Ng, \\& Portinari 1998).  Here we apply the synthetic CMD method to NGC 6819, NGC 2099, NGC 2168 and NGC 2323 to derive their age, reddening, distance modulus and (approximate) metallicity homogeneously to the Bragaglia (2003) cluster sample. The method also allows us to determine other features of these clusters, such as the existence (or lack thereof) of a significant fraction of unresolved binary systems, the original total mass of formed stars and the possible evaporation of some of the lower mass stars.  The organisation of the paper is as follows, \\S \\ref{observations} briefly summarises the data and the reduction procedures.  Further details are given in JSKI.  In \\S \\ref{firstresults} we present a summary of our main results which relate to this work from the published papers in the CFHT Open Star Cluster Survey.  \\S \\ref{synthetic} sets up the numerical simulations and presents details on how the synthetic CMDs were created.  Next, we compare the CMDs and the corresponding luminosity functions from the observations with the simulations on a cluster-by-cluster basis (\\S \\ref{vs}).   Finally, we discuss the results in \\S \\ref{discussion} and conclude the study in \\S \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  We re-derive key parameters for the four very rich open star clusters, NGC 6819, NGC 2099, NGC 2168, and NGC 2323.  The parameters are measured by comparing high quality, deep empirical colour-magnitude diagrams with MonteCarlo simulations. The combination of comparing the morphology and number of stars in various evolutionary phases and cluster luminosity functions allows us to provide tight constraints on the reddening, distance, age and binary fraction for each cluster.  In cases where the cluster metallicity is not certain, simulations with different abundances are also compared.  In all cases the data are better reproduced when a fraction of unresolved binary systems between 20 and 30\\% is assumed. A summary of the results is given in Table 2.  The synthetic CMDs and LFs are generally found to be in excellent agreement with the observational data.  This circumstance, combined with the fact that different sets of stellar evolution tracks provide different values for the cluster parameters, confirms how important it is to use more than one set of models to estimate the theoretical uncertainties. It also shows that a homogeneous approach is crucial to derive reliable overall cluster properties, such as age-metallicity relations. In fact, cluster dating based on different stellar models may lead not only to different absolute ages, but also to different age ranking, with significant drawbacks on the interpretation of the cluster properties in terms of Galactic evolution.  Some discrepancies, such as the thickness of the main sequences {\\it vis \\`a vis} the size of the photometric errors, and the possibility that cluster stars may be contaminating our blank field (particularly in the case of NGC 2168) are identified and discussed.  For each cluster, we also measure the astration mass (i.e., the total mass that went in all of the stars formed in the cluster) according to a single slope Salpeter IMF.  From this, we calculate the star formation rate between 0.1--100 $M_{\\odot}$ to be $9.4 \\times 10^{-6}, 1.3 \\times 10^{-5}, 6.2 \\times 10^{-6}$, and $2.7 \\times 10^{-6} M_{\\odot}yr^{-1}pc^{-2}$ for NGC 6819, NGC 2099, NGC 2168, and NGC 2323 respectively. These rates would be a factor 1.6 lower if the IMF below 0.6 $M_{\\odot}$ has a slope of +0.44, as inferred by Gould et al. (1997) from HST data, rather than Salpeter.  The true value may in fact lie somewhere in between these two IMFs, as recently discussed e.g., by Chabrier (2003). In principle our data are deep enough to allow for a direct derivation of the cluster IMF; however, the four clusters are heavily contaminated by fore/background stars and the lack of appropriate decontamination fields has prevented us from a safe analysis of the star counts at the fainter magnitudes. The need of appropriate photometry in nearby fields is emphasised, also for the purpose of adequate studies of the clusters mass segregation/evaporation."
    },
    "0403/astro-ph0403566_arXiv.txt": {
        "abstract": "We place observational constraints on a recently proposed Galactic population, dubbed the {\\it shroud} (Gyuk \\& Gates 1999, Gates \\& Gyuk 2001). The shroud would be a very thick Galactic disk of low luminosity objects, most likely old white dwarfs, proposed to explain the optical depth seen in microlensing surveys towards the Magellanic clouds. The shroud is a simple alternative to the lenses being distributed in a classical, near-spherical dark halo; the advantage of the shroud is that it would compose only a fraction of a dark halo's total mass.  In this paper, we argue that stars of the Galactic shroud would be detectable in the recent proper motion survey of Oppenheimer et al. (2001) if their absolute luminosities were brighter than $M_{R_{59F}} = 19.4$ or approximately $M_V = 18.6$. We adopt a range of simple models of the shroud's kinematics and morphology, and the colours and luminosities of its white dwarfs; via Monte-Carlo simulations, we predict the numbers expected in the Oppenheimer et al. survey, which would be clearly separated from the numbers produced by white dwarfs of the disk, thick disk and halo.  The number of white dwarf detections in the proper motion survey (98) is found to be well explained by the disk, thick disk and halo. With {\\it the most conservative} kinematic and density parameters for the shroud, and an absolute luminosity of the white dwarfs of $M_{R_{59F}} = 17.6$, we find that the proper motion survey would detect over 100 WDs, just from the shroud.  For a $M_{R_{59F}} = 19.4$ shroud, the survey would find $5 \\pm 2$ peculiar objects, whereas only two white dwarfs with such characteristics are found in the original data. $M_{R_{59F}} = 19.4$ corresponds to $M_V = 18.6$ for WDs with $(V - I) = -1.030$. ",
        "introduction": "The microlensing surveys carried out in recent years (e.g. EROS, MACHO and OGLE) have reported on lensing population of dark objects seen towards the Magellanic clouds. The favoured mass for these objects is approximately half a solar mass, suggesting they are white dwarfs, since main sequence M stars of this mass are bright enough to be detected directly in surveys. To date, no directly detected counterpart for the dark population has been found. One scenario is a population of ancient white dwarfs (WDs) which would comprise a significant fraction of the Galactic dark halo (up to 20 per cent of its mass). This suggestion has later become disfavoured because even quite dim WDs would be directly detectable in the most recent proper motion surveys even if they comprised `only' 2 per cent of the total mass of the dark halo (Reyl\\'e, Robin \\& Cr\\'ez\\'e 2001; Flynn, Holopainen \\& Holmberg 2003).  To explain the MACHO microlensing results (Alcock et al. 2000), and at the same time to avoid some of the problems with a massive dark halo population, Gates and Gyuk (hereafter, G\\&G) (1999, 2001) proposed a new population in the Milky Way, dubbed the Galactic `shroud'. It would be a WD population in the form of a very thick disk with a scale height of 2.0 -- 3.0 kpc. It would produce the same microlensing optical depth as a dark halo WD population without having to be enormously massive.  We study here the implications of the shroud using the same techniques that we used in a previous study (Flynn et al. 2003; hereafter, Paper I). In Paper I, we constrained the luminosity and the number density of a dark halo WD population with the same simulation that we use for the current study.  In the previous study, we used two proper motion surveys independently for constraining the halo population. After testing the simulation for Paper I, we are now confident to use only the more recent one of those surveys (Oppenheimer et al. 2001) for constraining the new population.  Oppenheimer et al. proper motion survey was originally designed to find dark halo WDs. However, the follow up studies that have investigated the possibility of dark halo WDs in this survey have found that the survey is also very sensitive to the conventional thick disk WDs (e.g., Paper I; Reid, Sahu \\& Hawley 2001; Reyl\\'e et al. 2001). Because the shroud has a similar velocity structure to the thick disk, the survey is also very sensitive to the shroud. Thus, the Oppenheimer et al. survey is optimal for our purposes.  In section 2, we briefly introduce how to find nearby WDs and separate them by stellar population.  In section 3, we go through the parameters of the model in detail, and in section 4, we describe the Oppenheimer et al. proper motion survey and its findings. We illustrate the effect of proper motion to the predicted number counts in section 5 and present our results in section 6. Finally, we conclude in section 7. ",
        "conclusions": "Strong limits have been placed on the luminosities of white dwarfs which could make up the putative `shroud' of the Galaxy, proposed as a solution to the optical depth measurements seen in the microlensing surveys towards the Magellanic clouds. We use the Oppenheimer et al. (2001) proper motion survey of 4000 square degrees to $R_{59F} = 19.7$, containing 98 spectroscopically confirmed WDs. A range of shroud models are investigated, and the number of WDs with high reduced proper motions compared to the Oppenheimer et al. data. Most of the models produce significantly more WDs than are actually observed; in particular, models in which we probe the very highest reduced proper motion source (indicating very low luminosity and high space velocities, as expected for the shroud component) allow us to limit the luminosity of the WDs in the shroud to $M_{R_{59F}} = 19.4$ (the survey pass band), which corresponds to $M_V = 18.6$ or $M_I = 19.6$. If the Galaxy possesses a shroud of WDs which produces the microlensing signal, these WDs must be fainter than the above limits. Very few models of Hydrogen atmosphere WDs cool to such faint levels within the age of the Universe."
    },
    "0403/astro-ph0403085_arXiv.txt": {
        "abstract": "\\cbstart The Chandra X-ray Observatory grating spectrometers allow study of stellar spectra at resolutions on the order of 1000. Prior x-ray observatories' low resolution data have shown that nearly all classes of stars emit x-rays. Chandra reveals details of line and continuum contributions to the spectra which can be interpreted through application of plasma models based on atomic databases.  For cool stars with hot coronae interpreted in the Solar paradigm, assumption of collisional ionization equilibrium allows derivation of temperature distributions and elemental abundances. Densities can be derived from He-like ion's metastable transition lines.  Abundance trends are unlike the Sun, as are the very hot temperature distributions. For young stars, there is evidence of accretion driven x-ray emission, rather than magnetically confined plasma emission. For some hot stars, the expected emission mechanism of shocked winds has been challenged; there is now evidence for magnetically confined thermal plasmas. The helium-like line emission in hot stars is susceptible to photoexcitation, which can also be exploited to derive wind structure.  \\cbend ",
        "introduction": "X-ray emission is ubiquitous among late-type and pre-main sequence stars, as has been amply demonstrated by imaging and low-resolution x-ray observatories.\\citep{Feigelson:Montmerle:1999} With the advent of the Chandra transmission-grating and the XMM-Newton reflection-grating spectrometers, we can now probe the nature of the x-ray emission in detail through high-resolution diagnostics.  Early Chandra results confirmed some of the abundance anomalies derived from low-resolution imaging spectra and also unambiguously confirmed that few-component temperature models generally are not adequate.  Chandra spectroscopy has also challenged long-standing hot-star wind and x-ray production theories. Much effort is now being spent to survey and analyze stellar x-ray spectra over a range of evolutionary states, spectral types, rotational periods, and activity levels.  Here we will examine some results for a variety of stars, the ``active'' binaries; young, low-mass stars; and hot, high mass stars, with emphasis on Chandra grating spectrometers.  We will not discuss results from the XMM-Newton observatory, though they are complementary in many ways. ",
        "conclusions": ""
    },
    "0403/hep-ph0403115_arXiv.txt": {
        "abstract": "An analysis is carried out within mSUGRA of the estimated number of events originating from upward moving ultra-high energy neutralinos passing through Earth's crust that could be detected by the Extreme Universe Space Observatory (EUSO).  The analysis exploits a recently proposed technique that differentiates ultra-high energy neutralinos from ultra-high energy neutrinos using their different absorption lengths in the Earth's crust. It is shown that for the  part of the parameter space, where the neutralino is mostly a Bino and with squark mass $\\sim 1$ TeV, EUSO could see ultra-high energy neutralino events within mSUGRA models with essentially no background. In the energy range $10^{9}~{\\rm GeV} < E_{\\tilde \\chi} < 10^{11}~{\\rm GeV}$ the unprecedented aperture of EUSO makes the telescope sensitive, after 3~yr of observation, to neutralino fluxes as low as  $d\\Phi/dE_{\\tilde \\chi} > 1.1 \\times 10^{-6}\\ (E_{\\tilde \\chi}/{\\rm GeV})^{-1.3}$~GeV$^{-1}$ cm$^{-2}$\\ yr$^{-1}$\\ sr$^{-1},$ at the 95\\% CL. Such a hard spectrum is characteristic of supermassive particles' $N$-body hadronic decay. The case in which the flux of ultra-high energy neutralinos is produced via decay of metastable heavy ($m_X = 2 \\times 10^{12}$~GeV) particles with uniform distribution throughout the universe, and primary decay mode into 5 quarks + 5 squarks, is analyzed in detail. The normalization of the ratio of the relics' density to their lifetime has been fixed so that the baryon flux produced in the supermassive particle decays contributes to about 1/3 of the events reported by the AGASA Collaboration below $10^{11}$~GeV, and hence the associated GeV $\\gamma$-ray flux is in complete agreement with EGRET data. For this particular case, EUSO will collect between 4 and 5 neutralino events (with 0.3 of background) in $\\approx 3$~yr of running. NASA's planned  mission, the Orbiting Wide-angle Light-collectors (OWL), is also briefly discussed in this context. ",
        "introduction": "mSUGRA~\\cite{msugra}  and its extensions (generically called SUGRA models) are currently the leading candidates for physics beyond the standard model. These models contain a consistent mechanism for the breaking of supersymmetry softly by gravity mediation. An attractive feature of these models is that with R parity conservation the lightest neutralino is a possible candidate for cold dark matter~\\cite{goldberg} in a significant part of the mSUGRA parameter space~\\cite{scaling}. Further, over most of the parameter space the phenomenon of scaling occurs~\\cite{scaling} so that the light neutralino is mostly the supersymmetric partner of the $U(1)_Y$ gauge boson $B_{\\mu}$, i.e., it is mostly a $U(1)_Y$ gaugino or a Bino~\\cite{scaling,roberts}.  The parameter space of mSUGRA is characterized by the universal scalar mass, $m_0$, the universal gaugino mass, $m_{\\frac{1}{2}}$, the universal trilinear coupling, $A_0$ (all taken at the grand unification scale $M_G\\sim 2\\times 10^{16}$ GeV), $\\tan\\beta =<H_2>/<H_1>$ where $H_2$ gives mass to the up quark, and $H_1$ gives mass to the down quark and the lepton. In addition the model contains the Higgs mixing parameter $\\mu$ which enters in the superpotential in the form $\\mu H_1H_2$. The magnitude of $\\mu$ is determined by the constraint of radiative electro-weak symmetry breaking in the theory while $sign\\mu$ is arbitrary and must be constrained by experiment.  mSUGRA has been put to stringent test by the recent precision data from the  satellite experiment, the Wilkinson Microwave Anisotropy Probe (WMAP) which imposes a narrow range for cold dark matter (CDM) so that~\\cite{bennett,spergel} $  \\Omega_{\\rm CDM} h^2 =0.1126^{+0.008}_{-0.009}$. The candidacy of neutralinos as the dark matter of the universe is based on relic densities surviving annihilation processes of non-relativistic particle. Detailed analyses show that mSUGRA allows for a small amount of the parameter space in agreement with WMAP observations~\\cite{wmap1}.  As a consequence, the applicability of mSUGRA demands that the contribution of other sources of CDM to the dark matter mix are negligible.  In this paper, we will be interested in a flux of ultrarelativistic neutralinos  resulting from decays of a population of CDM metastable superheavy particles~\\cite{Berezinsky:1997hy,Kuzmin:1997cm}. In concert with the previous statement, these  particles should contribute negligibly to the dark matter density.  The weak couplings of neutralinos imply an interaction length in air which is greater than the atmospheric depth, even at horizontal incidence. The interaction probability is then roughly  uniform throughout the atmosphere. As with neutrinos, showers initiated by neutralino primaries can be distinguished from hadronic events by restricting the zenith angle space to near horizontal -- this maximizes the probability to detect showers of weakly interacting primaries, while screening out the electromagnetic component of hadronic showers which are initiated high in the atmosphere. However, deeply developing neutralino cascades cannot be isolated from neutrino induced air showers.  In this paper we show that the part of the parameter space where the neutralino is mostly a Bino and the mass $m_{\\tilde q}$ of the first and second generation squarks is $\\sim$ 1 TeV can lead to ultra-high energy neutralino signals that may be seen by the Extreme Universe Space Observatory (EUSO)~\\cite{Catalano:mm,Scarsi:fy}. These two conditions are fully compatible with the WMAP constraint, and the neutralino as lightest supersymmetric particle. Further, with appropriate cuts the background events arising from ultra-high energy neutrinos are essentially negligible. We discuss now the details of the analysis.   The problem of discriminating between neutralino and neutrino induced showers with space-based experiments has been examined recently~\\cite{Barbot:2002et}. The method makes use of the Earth as a filter. Specifically, in the  region of the mSUGRA parameter space under consideration the cross section for neutralino-nucleon interaction is smaller than that for neutrino-nucleon scattering processes. Thus, by restricting the angular bin for arrival of upward going showers to a region where neutrinos are largely absorbed during traversal in the Earth, it may be possible to obtain a clean signal~\\cite{Anchordoqui:2001cg}. In Ref.~\\cite{Barbot:2002et}, the discussion was presented in terms of neutralino-nucleon cross sections parameterized as a series of constant fractions of the neutrino-nucleon cross section. In this paper, we first calculate the neutralino cross section in the squark-resonance approximation.  We then proceed to estimate the sensitivity of EUSO  to neutralino-induced air showers. The sensitivity will be characterized by a lower bound on the neutralino flux, which is then related to some particular models of $X$-particle decay. ",
        "conclusions": "Using a technique that exploits the different absorption lengths of neutrinos and neutralinos in the Earth's crust, we have estimated the sensitivity of EUSO to isolate upward coming showers of ultra-high energy $\\tilde\\chi$. The neutralino-nucleon interaction  has been approximated by resonant squark production, with the neutralino being largely Bino in composition, and $m_{\\tilde q}\\simeq$ 1 TeV. We have shown that, during the complete mission lifetime, the telescope will be sensitive to $E_{\\tilde \\chi}^2\\ d\\Phi/dE_{\\tilde \\chi} > 1.1 \\times 10^{-6}\\ (E_{\\tilde \\chi}/{\\rm GeV})^{0.7}$ GeV cm$^{-2}$\\ yr$^{-1}$ sr$^{-1}$ at the 95\\% CL, for $10^{9}~{\\rm GeV} < E_{\\tilde \\chi} < 10^{11}~{\\rm GeV},$ and for the region $m_{\\tilde q}=1.0\\pm 0.2$ TeV. A hard spectrum $\\propto E_{\\tilde\\chi}^{-1.3}$ is typical of super heavy relic $N$-body decays that are purely hadronic. This is a conservative estimate, since regeneration effects have been only considered in computing the neutrino background. We have explicitly analyzed the case in which the flux of ultra-high energy neutralinos is produced via decay of metastable heavy ($m_X = 2 \\times 10^{12}$~GeV) particles with uniform distribution throughout the universe, and primary decay mode into 5 quarks + 5 squarks. The normalization of $n_X/\\tau_X$  has been fixed to contribute about 1/3 of the events reported by the AGASA Collaboration below $10^{11}$~GeV~\\cite{Takeda:2002at}. For this particular case, EUSO will collect between 4 and 5 neutralino events (with 0.3 of background) in $\\approx 3$~yr of running.  Existing limits on the diffuse photon flux in the GeV region strongly limit the sensitivity of EUSO for primary 2-body decays of hadronic nature. This is because, for the normalization of the baryonic contribution to the ultra-high energy cosmic ray flux assumed in Ref.~\\cite{Barbot:2002et}, which is marginally consistent with new EGRET bounds, the accompanying neutralino flux produced for $q \\bar q$ and $q \\tilde q$  is about an order of magnitude below the flux generated in the primary 10-body decay discussed above. On the other hand, the telescope can still be sensitive to the leptonic mode $X \\rightarrow l \\tilde l.$ In this case, by reducing $n_X/\\tau_X$ by a factor of $\\sim$ 2, one lessens the problem with EGRET data and still leaves a window open for neutralino detection at EUSO.  We note that for a Bino-like neutralino, the primary decay mode (90\\% branching fraction) of the squark is $\\tilde q\\rightarrow q \\tilde g,$ with a subsequent decay $\\tilde g\\rightarrow q\\bar q \\ \\tilde\\chi.$ Thus, the neutralino energy of the decay is about 1/6 of the primary energy. In the remaining 10\\% of the decays, $\\tilde q\\rightarrow q \\tilde\\chi.$ In either case, the shower energies are far above the $\\sim$ 1 PeV threshold for the detector. If more detailed considerations are warranted in the future, regeneration effects during passage through Earth can be assessed, taking into account the energy losses of the decay modes. These effects will lead to some enhancement of $P$, and consequently of the event rate.  We turn now to a brief discussion on the potential of the planned NASA mission Orbiting Wide-angle Light-collectors (OWL)~\\cite{Stecker:2000ek}. This mission will involve photo detectors mounted on 2 satellites in low equatorial orbit (600 - 1200 km). The eyes of the OWL will stereoscopically image a geometric area of $\\sim 9 \\times 10^5$~km$^{2},$ yielding $A \\sim 3 \\times 10^{6}$~km$^2$ sr. With its superior effective aperture ($\\epsilon_{_{\\rm DC}} \\approx 10\\%$), a 10 yr mission lifetime will allow one to discern on contributions of metastable relics to the upper end of the cosmic ray spectrum at the level of 1 part in $10^2.$  Consequently, the data from this mission will allow one to probe more deeply the parameter space of mSUGRA and its extensions."
    },
    "0403/astro-ph0403017_arXiv.txt": {
        "abstract": "We present an analysis of the {\\em Chandra} High Energy Transmission Grating Spectrometer observation of the rapidly rotating ($P_{\\rm rot}=0.94\\,{\\rm d}$) post T~Tauri ($\\sim20 $\\,Myr old) star PZ~Telescopii, in the Tucana association. Using two different methods we have derived the coronal emission measure distribution, $em(T)$, and chemical abundances. The $em(T)$ peaks at $\\log T = 6.9$ and exhibits a significant emission measure at temperatures $\\log T > 7$. The coronal abundances are generally $\\sim 0.5$ times the solar photospheric values that are presumed fairly representative of the composition of the underlying star. A minimum in abundance is seen at a first ionization potential (FIP) of 7-8\\,eV, with evidence for higher abundances at both lower and higher FIP, similar to patterns seen in other active stars. From an analysis of the He-like triplet of \\ion{Mg}{11} we have estimated electron densities of $\\sim 10^{12}-10^{13}\\,{\\rm cm^{-3}}$. All the coronal properties found for PZ~Tel are much more similar to those of AB~Dor, which is slightly older than PZ~Tel, than to those of the younger T~Tauri star TW~Hya. These results support earlier conclusions that the soft X-ray emission of TW~Hya is likely dominated by accretion activity rather than by a magnetically-heated corona. Our results also suggest that the coronae of pre-main sequence stars rapidly become similar to those of older active main-sequence stars soon after the accretion stage has ended. ",
        "introduction": "\\label{intro}  One of the primary characteristics of low mass pre-main sequence (PMS) stars is their intense X-ray activity. This X-ray emission therefore represents an important means for investigating the properties and evolution of young stellar objects. X-ray activity is present during the evolution of PMS stars both in the initial evolutionary stages of a Classic T~Tauri Star (CTTS, Class I and II sources), during which the star has an accretion disk that surrounds it, and in the subsequent Weak-Line T~Tauri Star (WTTS, Class III sources) phase in which the star has no accreting material and is approaching the zero-age main sequence \\citep[ZAMS, ][]{FeigelsonMontmerle1999}. How the presence of accreting material influences the X-ray emission of PMS stars, and for how long, remain questions of debate.  The high resolution X-ray spectra now available with the {\\em Chandra} and {\\em XMM-Newton} satellites offer the possibility to perform detailed studies of stellar coronae because key emission lines diagnostics can now be resolved. These diagnostics can be used to derive elemental abundances, temperature and density structure of the emitting plasmas.  It is worth noting that very few young stars in star forming regions or associations are sufficiently X-ray bright to allow high resolution X-ray spectroscopy with current instrumentation. In particular, among the CTTSs, \\objectname{TW~Hya} is the best studied case to date because it is the nearest ($\\sim56$\\,pc) known CTTS. TW~Hya shows spectral characteristics very different from those of young but otherwise \\emph{normal} active stars \\citep{KastnerHuenemoerder2002,StelzerSchmitt2004}: very low plasma temperature ($\\log T \\sim 6.5$), high density ($\\log N_{\\rm e} \\sim 13$), very low Fe abundance ($A_{\\rm Fe}/A_{\\rm Fe\\sun} \\sim 0.2$). \\citet{KastnerHuenemoerder2002} and \\citet{StelzerSchmitt2004} attributed these characteristics to an accretion shock rather than coronal activity. Without the benefit of high resolution spectra of similar stars for comparison, however, the nature of the peculiarity of TW~Hya remains uncertain.  The situation with slightly more evolved stars is more clear. Studies of \\objectname{AB~Dor}, a young active star that has nearly arrived at the ZAMS and which has been observed in the past with several space-borne X-ray observatories, have shown that it is characterized by a hot corona ($T \\sim 10^7$\\,K) with plasma densities ranging from $6\\times10^{10}\\,{\\rm cm^{-3}}$ at $2\\times10^{6}$\\,K, to $3\\times10^{12}\\,{\\rm cm^{-3}}$ at $10^{7}$\\,K, and a moderately low Fe abundance ($A_{\\rm Fe}/A_{\\rm Fe\\sun} \\sim 0.25$, \\citealt{MeweKaastra1996,GudelAudard2001}; \\citealt*{Sanz-ForcadaMaggio2003}; P. Testa, in preparation;  D. Garc{\\'{\\i}}a-Alvarez, in preparation). In many respects, AB~Dor can be considered the prototype of very active single stars.  High resolution X-ray spectroscopy of other PMS stars with ages between few $10^{6}$\\,yr, typical of CTTSs like TW~Hya, and $10^{8}$\\,yr, the ZAMS of solar-type stars like AB~Dor, is crucial for understanding how the characteristics of stellar X-ray activity change during these early evolutionary phases.  In this work, we present a {\\em Chandra} High Energy Transmission Grating Spectrometer (HETGS) observation of \\objectname{PZ~Telescopii} (HD~174429, HIP~92680). Classified as a K0V star by \\citet{Houk1978}, PZ~Tel was determined by \\citet{ZuckermanWebb2000} to be a member of the Tucana association, a nearby star forming region about 45\\,pc away. From observations with the ROSAT PSPC, \\citet{StelzerNeuhauser2000} deduced a PZ~Tel X-ray luminosity of $L_{\\rm X}\\sim(2.88\\pm0.0 8)\\times10^{30}\\,{\\rm erg\\,s^{-1}}$. At a Hipparcos parallax distance of $49.7\\pm2.9$\\,pc \\citep{PerrymanLindegren1997}, PZ~Tel is a single star with a rotational period of 0.94\\,d (\\citealt{CoatesHalprin1980}; \\citealt*{InnisCoates1984,InnisThompson1986}). \\citet{FavataMicela1998} have estimated a mass of $1.1\\,M_{\\sun}$ for PZ~Tel and an age of approximately 20\\,Myr. Its youth is also confirmed by prominent H$\\alpha$ emission and a relatively undepleted Li abundance \\citep*{SoderblomKing1998}. Further evidence of its PMS status has been pointed out by \\citet{BarnesCollierCameron2000} who deduced from $v \\sin i$ and $P_{\\rm rot}$ that the minimum radius of PZ~Tel is $R\\,\\sin i \\sim 1.27\\,R_{\\sun}$: this value is larger than the radius of a main sequence star with the same mass as PZ~Tel.  The analysis of {\\em Chandra} HETGS high resolution spectra of PZ~Tel offers us the opportunity to study the coronal properties of a single star which has dissipated its accretion disk and is approaching the ZAMS. This study also allows us to compare the coronal properties of PZ~Tel with those of both younger and older stars, providing a glimpse of the evolution of stellar coronae through the PMS phase. In particular, the CTTSs have disks from which active accretion is still taking place. These stars could have quite different magnetospheric geometries, possibly involving magnetic connections between star and disk \\citep[e.g.][]{Montmerle2002}. In the case of coeval and older stars, it appears that coronal activity of PMS stars without an accreting disk can be explained on the same basis as that of main-sequence stars \\citep*{FlaccomioMicela2003} --- stellar rotation and convection. In this domain, PZ~Tel provides a new window on phenomena such as the chemical fractionation of elements that is seen to occur in coronae over a wide range of activity level. In the solar case, elements with low first ionization potential (FIP) such as Mg, Fe and Si are seen to be enhanced relative to elements with high FIP, such as O, Ne and Ar \\citep[e.g.][]{Feldman1992}. In more active stars, the situation appears somewhat reversed, with elements such as Ne appearing enhanced relative to those with lower FIP \\citep[e.g.][and references therein]{Drake2002}. The case of the active but very young post-T Tauri stars remains unexplored at high spectral resolution.  In \\S~\\ref{obs} and \\ref{analysis} we describe the {\\em Chandra} observation and  the techniques used in its analysis. Section~\\ref{results} presents the resulting coronal temperature structure and abundances. These are discussed in \\S~\\ref{disc}, in which we also present a comparison of the coronal properties of PZ~Tel with those of AB~Dor and TW~Hya.  \\begin{figure*}[t] \\centering \\includegraphics[width=13cm] {figure1.ps} \\caption{Light curve of PZ~Tel obtained from the HEG and MEG spectra (excluding zero-order events) with bin size of 1000\\,s.} \\label{fig:pztel_lightcurve} \\end{figure*} ",
        "conclusions": ""
    },
    "0403/astro-ph0403221_arXiv.txt": {
        "abstract": " ",
        "introduction": "My visits to Japan have occurred always in crucial times in the development of our research. In 1975 I visited the University of Kyoto and gave the lectures which were then co-authored in the Japanese book with Humitaka Sato \\cite{bookhs}. The focus then was on three major topics: a) the basic paradigm for the identification of a black hole I had just established and which had found a very significant application in Cygnus X-1 through the splendid data obtained by Riccardo Giacconi and Minoru Oda \\cite{bookhs}; b) the Cristodoulou-Ruffini \\cite{CR71} mass-energy formula for black holes: \\begin{equation} E_{BH}^2=M^2c^4=\\left(M_{\\rm ir}c^2 + \\frac{Q^2}{\\rho_+}\\right)^2 + \\frac{L^2c^2}{\\rho_+^2}\\, , \\label{em} \\end{equation} where $M_{\\rm ir}$ is the irreducible mass, $\\rho_+ = 2(G/c^2) M_{\\rm ir}$ is the quasi-spheroidal cylindrical coordinate of the horizon evaluated at the equatorial plane and $Q$ and $L$ are respectively the charge and angular momentum of the black hole; this mass-energy formula allows to estimate the maximum energy extractable from a process of gravitational collapse; c) a specific energy extraction process from the black hole by pair creation due to supercritical electric fields, first introduced by Sauter \\cite{S31}, Heinsenberg \\& Euler \\cite{HE35}, Schwinger \\cite{S51}, I developed with T. Damour \\cite{DR75}. In that paper we had also pointed out that such process could be the source of the then newly discovered Gamma-Ray Bursts (GRBs). Our model had a very distinct signature, which differentiates it from all the other models: the characteristic energy of the GRBs should be of the order of $10^{54}$ ergs (see Fig. \\ref{fig1a}--\\ref{fig1b}).  \\begin{figure}[t] \\begin{center} \\includegraphics[width=8cm,clip]{foto} \\end{center} \\caption{Princeton 1971.} \\label{fig1a} \\end{figure}  \\begin{figure}[t] \\begin{center} \\includegraphics[width=\\hsize,clip]{dia3a} \\end{center} \\caption{The basic components of the Damour \\& Ruffini black hole vacuum polarization.} \\label{fig1b} \\end{figure}  The strategy we had followed, both in the case of Cygnus X1 and the GRBs, was not to try to understand the astrophysical aspects of the phenomenon evidencing the black hole formation. On the contrary we had capitalized on the physics of the black hole and on specific properties of the solutions of Einstein Maxwell Equations in order to infer specific signatures to be expected in the astrophysical scenario in order to obtain the observational evidence for a black hole in a realistic astrophysical setting. Indeed the paradigm for the identification of the black hole in Cygnus X1 (Leach \\& Ruffini \\cite{lr73}) was based mainly on three general relativistic considerations: a) the comprehension of the gravitational binding energies around a Kerr black hole, which I found with Wheeler in 1969 \\cite{ll}, clearly pointing to the possibility of having accretion energy as the origin of the observed enormous luminosities in X-ray observed in Cygnus X1, $L=10^4L_\\odot$; b) the uniqueness theorem of black hole (see e.g. Ruffini \\& Wheeler \\cite{rw71}) endowed only of charge mass and angular momentum, clearly pointing to the impossibility of having periodic signals out of a black hole; c) the existence of an absolute maximum mass of a neutron star, again derived out of first principles, from the equation of equilibrium in the Einstein theory of gravity, the principle of causality implying speed of sound not exceeding the speed of light, and the existence of a fiducial density (Rhoades \\& Ruffini \\cite{rr74}). All these points were later summarized in the proceedings of the Varenna School organized by Riccardo Giacconi and myself \\cite{gr75,gr78} and in the Solvay conference \\cite{r74}. Riccardo Giacconi, in his splendid lecture \\cite{nobel} recalls the significance of this theoretical work for the understanding of binary X-ray sources.  In the case of GRBs our approach was similar: priority was given to the identification of the energy source of GRBs. That nuclear energy is the energy source of main sequence stars has been credibly proved \\cite{sch}, that accretion and gravitational energy release around neutron stars and black holes was the Energy sources of binary x-ray sources had been demonstrated \\cite{nobel}, we decided to look in the possibility of having a new energy source as powering the GRBs: the extractable energy of a black hole \\cite{CR71,DR75}. The mechanism I had conceived with T. Damour was indeed viable, supported by the very basic and well established physical principles on which it was grounded.  My second scientific visit to Japan occurred in occasion of the sixtieth birthday of Humitaka Sato: again I reported \\cite{rukyoto} some new progress in our research: a) The situation with GRBs had dramatically modified by the observations of the Italian-Dutch satellite BeppoSAX (Costa \\cite{ca97}) which gave origin to an unprecedented collaboration between X- and $\\gamma$-ray, optical and radio astronomy. This observational effort had lead to the determination of the distances of GRBs had unequivocally established the cosmological nature of their source: indeed energetics of the order of $10^{54}$ ergs were implied as predicted by our model with Damour \\cite{DR75}. b) I clarified some basic conceptual issues on the energy extraction process from a black hole endowed with electromagnetic structure, introducing the novel concept of ``dyadosphere'' of a black hole, as the region surrounding the black hole horizon where the electron-positron pairs created in the process of vacuum polarization are localized. c) I finally pointed out that the very process of thermalization of such an electron-positron plasma created in the dyadosphere is the main mechanism originating the GRB expansion and the engine of the entire GRB phenomenon \\cite{RSWX99,RSWX00} (see Fig. \\ref{fig2a}--\\ref{fig2b}).  \\begin{figure}[t] \\begin{center} \\includegraphics[width=\\hsize,clip]{dia6a} \\end{center} \\caption{The basic parameters of the dyadosphere.} \\label{fig2a} \\end{figure}  \\begin{figure}[t] \\begin{center} \\includegraphics[width=8cm,clip]{cl3d} \\end{center} \\caption{The theoretically predicted luminosity and spectral distribution of a short GRB. Details in Ruffini et al. \\cite{RFVX03}.} \\label{fig2b} \\end{figure}  In this visit I like to report new results connected with: a) the physics of the dyadosphere and a possible future verification derived from the observation of short GRBs; b) the three fundamental paradigms for the theoretical interpretation of GRBs; c) our understanding of the long bursts and the GRBs afterglows. ",
        "conclusions": ""
    },
    "0403/astro-ph0403198_arXiv.txt": {
        "abstract": " ",
        "introduction": "A number of observations of BHC's do not show consistent correlations of low QPO frequency with disk parameters [1,2]. On the other hand, strong consistent correlations between power law spectral index and low frequency QPO's have been recently observed [3]. In addition, there is mounting observational evidence that the large number of  spectral \"states\" formerly developed by to explain the wide variability of BHC's such as GRS1915+105, can be reduced to a few canonical states, i.e. a hard state with spectral power law index $\\Gamma \\sim1.6\\pm 0.1$ , a soft or \"extended\" power-law state characterized by $\\Gamma \\sim 2.7\\pm0.2$, and a thermal state [4]. \\par These data have prompted us to introduce a model, i.e. the Transition Layer (TL) model [5] to explain the correlations observed. The main feature of the TL model is a hot compact region near the BH which serves as the primary region for Compton upscattering of soft disk photons. The TL model shows how the QPO's are related to the size, optical depth, temperature and spectral index and predicts the correlation between index and QPO frequency. ",
        "conclusions": ""
    },
    "0403/astro-ph0403151_arXiv.txt": {
        "abstract": "We present Very Large Array observations at 7 mm of the sources IRAS 2A, IRAS 2B, MMS2, MMS3 and SVS 13, in the NGC1333 region. SVS 13 is a young close binary system whose components are separated by 65 AU in projection. Our high angular resolution observations reveal that only one of the components of the SVS 13 system (VLA 4B) is associated with detectable circumstellar dust emission.  This result is in contrast with the well known case of L1551~IRS5, a binary system of two protostars separated by 45 AU, where each component is associated with a disk of dust. Both in SVS 13 and in L1551~IRS5 the emission apparently arises from compact accretion disks, smaller than those observed around single stars, but still massive enough to form planetary systems like the solar one. These observational results confirm that the formation of planets can occur in close binary systems, either in one or in both components of the system, depending on the specific angular momentum of the infalling material. ",
        "introduction": "SVS~13, in the NGC1333 region, was discovered as a 2.2 $\\mu$m source by Strom, Vrba, \\& Strom (1976), and since the source is roughly aligned with the chain of Herbig-Haro objects 7-11 (Herbig 1974; Strom, Grasdalen, \\& Strom 1974), it was assumed to be the exciting source of this classical HH system. Later, Goodrich (1986) detected a faint visible counterpart of SVS~13. However, the star SVS~13 presents a number of peculiar properties. The source exhibited a significant increase of its brightness at optical ($\\sim$3 mag), and IR ($\\sim$1 mag) wavelengths in 1988-1990 (Eisl\\\"offel et al. 1991; Liseau, Lorenzetti, \\& Molinari 1992; Harvey et al. 1998), and since then, the flux has remained almost steady (Aspin \\& Sandell 1994; Khanzadyan et al. 2003). In addition, despite being optically visible, indicating that it is a relatively evolved young object, SVS 13 is a strong millimeter source, known as MMS1 (e.g., Grossman et al. 1987; Looney, Mundy, \\& Welch 2000), and presents other characteristics, such as the presence of an extremely high velocity CO outflow, that suggest it is in a much earlier evolutionary stage (a Class 0/I object; Bachiller et al. 2000).  Radio continuum emission from SVS 13 was first reported by Snell \\& Bally (1986). Rodr\\'{\\i}guez, Anglada, \\& Curiel (1997, 1999) mapped SVS 13 with the Very Large Array (VLA) at 3.6 and 6 cm (their source VLA 4). Anglada, Rodr\\'{\\i}guez, \\& Torrelles (2000), through VLA observations at 3.6 cm of higher angular resolution and sensitivity, discovered that SVS~13 is, in fact, a close binary system. The two components of the binary (VLA 4A and VLA 4B) are separated by $0\\rlap.''3$, corresponding to 65 AU in projection (assuming a distance of 220 pc; \\v{C}ernis 1990), and have similar flux densities at 3.6 cm.  The water masers associated with SVS~13 appear segregated in position and velocity, supporting the binary hypothesis (Rodr\\'{\\i}guez et al. 2002).  Anglada et al.  (2000) noted that the optical position for SVS~13 (as measured by Rodr\\'{\\i}guez et al. 1997) is closer to VLA~4A, while the millimeter position (Looney et al. 2000) is closer to VLA~4B.  Although the precision of the astrometry available at that time did not allow an unambiguous association, this result led these authors to suggest that the strong millimeter emission reported for SVS~13 could arise from only one of the components of the binary (VLA~4B, the eastern component), while the optical emission would come from the other component (likely from VLA~4A, the western component). In the interpretation proposed by Anglada et al. (2000), only one of the stars (VLA 4B) is surrounded by a dusty envelope or disk, while the other (VLA 4A, the visible star) is not. A similar interpretation has been proposed recently by Loinard et al. (2002) for IRAS 04368+2557 in L1527, on the basis of the morphology of the two obscured sources observed at 7 mm.  A confirmation of this interpretation for SVS 13 requires a precise comparison of the positions of the sources observed at different wavelengths, in order to identify the individual contribution of each component of the binary. Since the angular separation between VLA 4A and VLA 4B is only $0\\rlap.''3$, absolute astrometry down to $<0\\rlap.''1$ is required in order to obtain an accurate enough registration of the positions. Although the accuracy of the absolute astrometry of the optical observations ($\\pm0\\rlap.''3$) is difficult to improve due to the lack of a large enough number of suitable reference stars in the field, it is possible to improve the accuracy of the registration between the centimeter and millimeter positions by using the same instrument and calibration procedures in both wavelength ranges.  In this Letter, we present VLA observations at 7 mm of the region near SVS 13, carried out in the D and B configurations, using the same phase calibrator and procedures as in the previous VLA observations at 3.6 cm. The B configuration observations at 7 mm provide an angular resolution of $\\sim0\\rlap.''2$, similar to that of the A configuration at 3.6 cm, and an expected accuracy in the registration between both images down to $<0\\rlap.''05$, allowing a precise comparison of the emission at both wavelengths, necessary to test the single disk hypothesis for the SVS 13 binary. These observations also provide 7 mm data on other sources in NGC1333:  VLA 2 (MMS3), VLA 7 (IRAS 2A), VLA 10 (IRAS 2B), and VLA 17 (MMS2). ",
        "conclusions": "We detect at 7 mm the sources VLA 2 (MMS3), VLA 7 (IRAS 2A), VLA 10 (IRAS 2B), and VLA 17 (MMS2=SVS 13B) (see Table 1).  These sources were previously observed at 3.6 and 6 cm by Rodr\\'\\i guez et al. (1999), where a discussion on their properties and counterparts can be found. VLA 7, VLA 10, and VLA 2 were further observed at 3.6 cm with higher angular resolution by Reipurth et al. (2002). Interferometric observations at 3 mm of VLA 7 and VLA 10 have been recently reported by J{\\o}rgensen et al. (2004).  We also detect SVS 13 at 7 mm, as an unresolved source in the D configuration, and resolving it in its two components (VLA 4A and VLA 4 B) in the B configuration observation (see Table 1). In Figure 1 we compare the 3.6 cm map observed with the A configuration (Anglada et al. 2000) with the 7 mm map observed with the B configuration (this paper). Both maps were obtained with a similar angular resolution of $\\sim 0\\rlap.''2$. The positions of the sources in the two maps are in agreement within $0\\rlap.''02$, eliminating, thus, any ambiguity in their identification at different wavelengths. As can be seen in the figure, at 3.6 cm both sources present a similar flux density, while at 7 mm VLA 4B is much stronger than VLA 4A, suggesting that VLA 4B is the dominant source in the millimeter wavelength range, with a negligible contribution from VLA 4A.  In order to quantitatively confirm this hypothesis it should be verified that the increase of flux density of VLA 4B over VLA 4A at 7 mm is caused by dust emission, and not by free-free emission with a steep spectral index. Usually, the emission at wavelengths longer than a few cm is dominated by the free-free emission from ionized gas, while the emission at wavelengths shorter than a few mm is dominated by thermal emission from dust. Since the wavelength of 7 mm falls in between the centimeter and millimeter regimes, it should be checked, for each particular object, which is the nature of the dominant emission at this wavelength. To do that, we have plotted in Figure 2 the highest angular resolution data available on SVS 13 in the centimeter and millimeter ranges. The flux densities of VLA 4B at 3.6 cm and 1.3 cm give a spectral index of $\\alpha=1.36\\pm0.26$, which is typical of a thermal ionized jet or a partially optically thick ionized region. An extrapolation at 7 mm of this free-free emission yields a flux density of $\\sim$1 mJy, much smaller than the observed value. On the other hand, a fit to the observed data from 3.4 mm to 1.3 mm (that are supposed to trace dust emission) gives a spectral index $\\alpha=2.55\\pm0.05$ in the millimeter range, resulting in an extrapolated flux density at 7 mm in good agreement with the value observed for VLA 4B (see Table 1 and Fig. 2). Then, we conclude that free-free emission cannot account for the observed flux density of VLA 4B at 7 mm, and that a contribution of a different nature, namely, thermal dust emission, likely from a circumstellar disk surrounding this object, is required.  In the case of VLA 4A, the overall 3.6 cm, 1.3 cm, and 7 mm data points can be fitted together with a single spectral index $\\alpha=1.25\\pm0.20$, indicating that the observed flux density at 7 mm can be explained as free-free emission (see Fig. 2). However, we cannot discard a small contribution from dust emission at 7 mm, since an extrapolation of the 3.6 and 1.3 cm data alone gives a free-free contribution at 7 mm slightly below the observed flux density of VLA 4A. The difference ($\\la$0.8 mJy) could correspond to a contribution from dust, being it about five times smaller than in the case of VLA 4B.  In summary, our results indicate that of the two components of the SVS 13 binary system, the source VLA 4B is associated with a much larger amount of dust than the source VLA 4A. This strongly supports the proposal by Anglada et al. (2000), who suggested that the millimeter source MMS1 is the counterpart of the centimeter source VLA 4B. On the basis of the available optical astrometry, these authors also proposed that VLA 4A is the counterpart of the visible star SVS 13.  VLA 4B appears as a compact source in our 7 mm map obtained with the B configuration with an angular resolution of $\\sim0\\rlap.''2$ (Fig. 1b), suggesting that the dust emission traced by this source originates in a compact circumstellar structure, likely a disk, with radius $\\la 30$ AU. The size of the disk is smaller than that of the typical accretion disks observed around single T Tauri stars, with radii of 100-150 AU (e.g., Dutrey et al.  1996; Wilner et al.  2000), and is more similar to that of the compact disks observed in the L1551~IRS5 binary system (10 AU; Rodr\\'\\i guez et al. 1998), suggesting that they are truncated by the tidal effects of the companion star. The 7 mm flux density of VLA 4B, together with the millimeter data shown in Figure 2, can be fitted by a simple disk model. We assume a geometrically thin, vertically isothermal disk, characterized by power-law radial dependences of temperature and surface density, and we adopt the opacity law given by D'Alessio, Calvet, \\& Hartmann (2001).  The data are fitted with a distribution of dust grains with a maximum size of 1 mm, a temperature distribution $T(r)=470~(r/\\rm AU)^{-0.5}$~K, a surface density $\\Sigma(r)=2800~(r/\\rm AU)^{-1}$~g~cm$^{-2}$, implying a disk mass of 0.06 $M_\\odot$, for a radius of 30 AU. Thus, it seems plausible that the emission of VLA 4B is tracing a protoplanetary disk, since the mass obtained exceeds the minimum mass required to form a planetary system like the solar one ($\\sim 0.01~M_\\odot$). Since any dust emission associated with VLA 4A is at least five times weaker than that of VLA 4B, if a disk was associated with VLA 4A, we expect its mass to be at least five times smaller than that of the disk associated with VLA 4B. We note that the flux density observed in the D configuration is larger than the total flux observed in the B configuration, indicating an additional contribution from an extended component, likely the infalling envelope.  A fit to the overall spectral energy distribution, taking into account simultaneously the contributions of different components (disks+envelope), similarly to what has been done for L1551 IRS5 (Osorio et al. 2003), would be useful to better constrain the properties of each component.  Thus, the observational results obtained for SVS 13 imply that it is feasible that the development of a protoplanetary disk occurs preferentially in only one of the components of a young close binary system, with the disk absent or much less significant in the other component.  In this respect, the case of the SVS 13 binary system appears to be opposite to the L1551~IRS5 case, where both components of the binary system are associated with circumstellar disks of dust of comparable characteristics (Rodr\\'\\i guez et al. 1998).  These results are in agreement with theoretical simulations (e.g., Bate \\& Bonnell 1997) that show that, depending on the mass ratio of the components and the specific angular momentum of the system, the development of circumstellar disks can occur either around a single component or around both components of the binary. According to Bate \\& Bonnell (1997), a circumstellar disk forms around one of the components of the binary only if the specific angular momentum of the infalling gas is greater than the specific orbital angular momentum of that component about the center of mass of the binary. Thus, if a binary system grows to its final mass mainly via the accretion of material with low specific angular momentum, the primary may have a large circumstellar disk, while the secondary is essentially naked. The systems formed by this method are expected to be binaries with separations of the order of $\\sim 100$ AU, and they should not be developing a significant circumbinary disk. For infall with high angular momentum, both components can develop a circumstellar disk and even a circumbinary disk can be formed.  Under this simplified scheme, the low angular momentum scenario, with a disk in only one star, appears to correspond to the case of the SVS 13 binary, where the two components are separated by 65 AU in projection, and where VLA 4A would be the secondary, whereas VLA 4B (the dominant source at millimeter wavelengths) would be the primary viewed through its circumstellar disk. On the other hand, L1551~IRS5 apparently corresponds to a case of accretion of material with higher angular momentum, with both components associated with a circumstellar disk (Rodr\\'\\i guez et al. 1998) and probably surrounded by a circumbinary disk (Osorio et al. 2003).  In order to confirm these suggestions, it would be interesting to obtain the orbital parameters of these systems and to identify which component is the primary through accurate measurements of absolute proper motions. Relative proper motions between the two components of the L1551~IRS5 system have been obtained recently (Rodr\\'\\i guez et al. 2003), and accurate absolute proper motions have been measured for the T Tau Sa/Sb system (Loinard, Rodr\\'\\i guez \\& Rodr\\'\\i guez 2003). At present, no proper motions are available for the SVS 13 system, but these promising results obtained for other sources suggest that this goal could be attainable in a relatively near future, providing a complete test for the properties of the SVS 13 binary system."
    },
    "0403/astro-ph0403367_arXiv.txt": {
        "abstract": "A semianalytic method to estimate the angular resolution of tracks, that have been reconstructed by a likelihood approach, is presented. The optimal choice of coordinate systems and resolution parameters, as well as tests of the method are discussed based on an application for a neutrino telescope. ",
        "introduction": "This paper describes a statistical procedure to extract resolution estimates on a track by track basis. The method was developed for the AMANDA Neutrino Telescope at the South Pole~\\cite{tillpub:Andres:2001ty}, which uses a 3-dimensional grid of photosensors imbedded in highly transparent ice to provide spatial and time resolution of Cherenkov photons, that e.g. arise from long muon tracks.  The knowledge of track resolutions is of particular importance in the search for localized sources, such as distant galaxies. The resolution information can in addition be used to suppress mis-reconstructed tracks, that typically are less well defined. The method is not limited to muon reconstruction in neutrino telescopes and can be applied to any experiment in which tracks have been reconstructed with a likelihood approach. ",
        "conclusions": "\\label{section:tests}  The above algorithms have been tested within the framework of the AMANDA neutrino telescope. The estimates perform in a stable and reliable way and produce sensible results.  To ensure the correctness of the estimations, several testing procedures both in data and in Monte Carlo have been carried out:  \\begin{enumerate} \\item For $\\sigma_{\\theta}$ and $\\sigma_{\\phi}$ its {\\em pull\\/} has been studied. It is defined as the ratio of the difference of true and reconstructed direction over the resolution estimator. If the estimation is good, the pull should be Gaussian distributed, centered at zero and with unit width. This investigation can only be done in a Monte Carlo simulation. \\item For the spatial angle resolution the pull is not a sensible quantity, because angles in space can only be positive. Hence one can study ensembles of events with the same $\\sigma_{\\mathrm{a}}$ or $\\sigma_{\\mathrm{a}}^{\\epsilon}$, respectively. The median of the true deviation can then be compared to the corresponding $\\sigma_{\\mathrm{a}}$. Again this can only be done in the simulation. \\item In data the true directions are not known. Thus another approach is being followed. Each event is split into two subevents by assigning every second hit to subevent 1 and the remaining hits to subevent 2. Each subevent undergoes reconstruction and is subjected to the resolution estimation algorithm. That leads to directions\\footnote{Note that the reconstructed directions are expressed in standard detector coordinates, whereas the error estimates are obtained in their respective rotated systems. That is due to technical reasons only.} $\\vartheta^i$, $\\varphi^i$ and errors $\\sigma_{\\theta}^i$ und $\\sigma_{\\phi}^i$ with $i\\in\\{1,2\\}$.  The pull \\begin{equation} \\mathrm{P}_{\\vartheta} = \\frac{\\mathrm{D}_{\\vartheta}}{\\sigma_{\\theta}^{\\mathrm{D}}} = \\frac {\\vartheta^1 - \\vartheta^2} {\\sqrt{(\\sigma_{\\theta}^1)^2 + (\\sigma_{\\theta}^2)^2}} \\end{equation} should also be a Gaussian distribution centered at zero with unit width. For $\\mathrm{P}_{\\varphi}$ the difference in azimuth angles must be multiplied by $\\sin\\vartheta$ to make up for the differences in the rotated and not rotated coordinate systems: \\begin{equation} \\mathrm{P}_{\\varphi} = \\frac{\\mathrm{D}_{\\varphi}}{\\sigma_{\\phi}^{\\mathrm{D}}} = \\frac {(\\varphi^1 - \\varphi^2)\\cdot\\sin\\vartheta} {\\sqrt{(\\sigma_{\\phi}^1)^2 + (\\sigma_{\\phi}^2)^2}} \\qquad . \\end{equation} Of course this last check can also be done in the simulation.  \\end{enumerate}  In all tests the results were satisfactory. Further details about the tests and their results can be found in \\cite{phd:neunhoeffer_eng:2003}.%  \\subsection{Concluding remarks}  \\begin{enumerate} \\item Often one applies quality cuts to the data set, which rely on additional information that is not reflected in the construction of the likelihood function itself. Hence this information is also not used in obtaining the resolution parameters. In this case one expects the selected tracks to be on average closer to the true direction than from the likelihood analysis alone. Consequently one observes a pull with a width smaller than 1. \\item The resolution parameters - such as $\\sigma_{\\mathrm{a}}$ - can efficiently be used as cut variables, as misreconstructed events on average have a worse resolution than correctly reconstructed ones. \\item The resolution estimation on an event-per-event basis can be used as additional input in point source searches. The information is hard to be included in the commonly used binned search algorithms \\cite{tillpub:ahrens:2003xy3}. An appropriate procedure based on maximum likelihood methods that can integrate the new information in a natural way will be presented in a forthcoming publication. \\end{enumerate}   \\begin{ack}  I wish to thank the AMANDA collaboration for their support in obtaining the method described herein. I thank Lutz K{\\\"o}pke and Alexander Holfter for many an illuminative and productive discussion. I would also like to thank the German Research Foundation (DFG) and the German Ministry of Research and Education (BMBF) for financial support of the AMANDA project. \\end{ack}"
    },
    "0403/astro-ph0403684_arXiv.txt": {
        "abstract": "We present an analysis of candidate members of the $\\eta$~Cha and MBM~12A young associations. For an area of 0.7~deg$^2$ toward $\\eta$~Cha, we have performed a search for members of the association by combining $JHK_s$ photometry from 2MASS and $i$ photometry from DENIS with followup optical spectroscopy at Magellan Observatory. We report the discovery of three new members with spectral types of M5.25-M5.75, corresponding to masses of 0.13-0.08~$M_{\\odot}$ by theoretical evolutionary models. Two and three of these members were found independently by Lyo and coworkers and Song and coworkers, respectively. Meanwhile, no brown dwarfs were detected in $\\eta$~Cha down to the completeness limit of 0.015~$M_{\\odot}$. For MBM~12A, we have obtained spectra of three of the remaining candidate members that lacked spectroscopy at the end of the survey by Luhman, all of which are found to be field M dwarfs. Ogura and coworkers have recently presented four ``probable\" members of MBM~12A. However, two of these objects were previously classified as field dwarfs by the spectroscopy of Luhman. In this work, we find that the other two objects are field dwarfs as well. ",
        "introduction": "The environs of the B8 star $\\eta$ Cha and the dark cloud MBM~12 have been revealed as sites of small associations of newly formed stars. Through deep {\\it ROSAT} observations of 0.35~deg$^2$ toward $\\eta$ Cha, \\citet{mam99} discovered X-ray emission from 12 sources, including $\\eta$ Cha itself. From {\\it Hipparcos} parallactic distances and proper motions, the two earliest stellar counterparts and an early-type star without X-ray emission were found to be comoving at a distance of $\\sim100$~pc. \\citet{mam99} confirmed the youth of the remaining 10 late-type stars through measurements of H$\\alpha$ emission and Li absorption, which indicated their probable membership in an association with the early-type stars. An age of 4-10 Myr has been inferred for these stars through comparisons of their positions on the Hertzsprung-Russell (H-R) diagram to theoretical isochrones \\citep{mam99,law01}. Recently, five additional low-mass members of the association (0.08-0.3~$M_{\\odot}$) have been identified through spectroscopy of candidates appearing in color-magnitude diagrams generated from photometry at $V$ and $I$ \\citep{law02,lyo04} and from data in the USNO-B1 and Two-Micron All-Sky Survey (2MASS) catalogs \\citep{sz04}. Although not precisely determined, the completeness limits of these latest surveys appeared to be near 0.1~$M_{\\odot}$.  Surveys for sources of H$\\alpha$ and X-ray emission toward cloud 12 from \\citet{mbm85} (MBM~12) have resulted in the discovery of seven T Tauri stars (e.g., \\citet{hb88}). MBM~12 was long believed to be the nearest molecular cloud at a distance of 50-100 pc \\citep{hbm86,hea00}, making the associated group of young stars (MBM~12A) a particularly attractive site for studies of star formation. However, during a search for new members, \\citet{luh01} noticed that the members of MBM~12A exhibited anomalously old ages ($\\sim100$~Myr) when placed on the H-R diagram with the published distances. Through new distance estimates in that work and subsequent studies \\citep{str02,and02}, the distance of MBM~12 cloud is now firmly established near a value of $\\sim300$~pc. Using this distance, theoretical isochrones produce an age of $\\sim2$~Myr, which is consistent with the evolutionary state of the association implied by various signatures of newly formed stars \\citep{luh01}. The magnitude-limited survey by \\citet{luh01} achieved a completeness limit of 0.03~$M_{\\odot}$ and uncovered five new members (0.1-0.4~$M_{\\odot}$), bringing the total known membership to a dozen sources at a resolution of $\\sim1\\arcsec$. \\citet{ogu03} have since identified four ``probable\" members of MBM~12A via H$\\alpha$ emission detected with slitless grism spectroscopy.  In this paper, we present spectroscopy of candidate members of the $\\eta$~Cha and MBM~12 associations. For $\\eta$~Cha, we select candidate members of the association by combining $i$ photometry from the Deep Near-Infrared Survey of the Southern Sky (DENIS) and $JHK_s$ data from the Two-Micron All-Sky Survey (2MASS) (\\S~\\ref{sec:ident1}), measure their spectral types and determine their status as field stars or association members (\\S~\\ref{sec:class1}), evaluate the completeness of the survey (\\S~\\ref{sec:complete}), and place the known members of $\\eta$~Cha on the H-R diagram (\\S~\\ref{sec:hr}). For MBM~12A, we select candidate members from \\citet{luh01} and \\citet{ogu03} for spectroscopy (\\S~\\ref{sec:ident2}) and use these data to measure spectral types and assess membership (\\S~\\ref{sec:class2}). ",
        "conclusions": "We have performed spectroscopic studies of candidate members of the $\\eta$~Cha and MBM~12A young associations, the conclusions for which are summarized as follows:   For $\\eta$~Cha, we have used $JHK_s$ photometry from 2MASS and $i$ photometry from DENIS to construct color-color and color-magnitude diagrams for an area of 0.7~deg$^2$. Through spectroscopy of the candidate members appearing in these data, we have discovered three new members of the association with spectral types of M5.25-M5.75, corresponding to masses of 0.13-0.08~$M_{\\odot}$ according to evolutionary models. These sources were independently found in recent work by \\citet{sz04} and \\citet{lyo04}. No brown dwarfs were detected in $\\eta$~Cha down to the completeness limit of 0.015~$M_{\\odot}$ for our survey, which is roughly consistent with the yield of $\\sim2$ substellar members expected if the relative numbers of stars and brown dwarfs in $\\eta$~Cha are similar to those of the Taurus and IC~348 star-forming regions. On the other hand, the fact that the three least massive members of the association are the outermost members may indicate that the substellar members are preferentially located at large distances from the center of the association and outside of our survey field. For the 18 known members of $\\eta$~Cha, we have estimated bolometric luminosities and effective temperatures and placed the members on the H-R diagram, from which we infer an age of $6^{+2}_{-1}$~Myr for the association with the evolutionary models of \\citet{bar98}.  For MBM~12A, we have presented spectra of three of the remaining candidate members that lacked spectroscopy at the end of the survey by \\citet{luh01} and of the four ``probable\" members from \\citet{ogu03}. We classify all of these sources as field dwarfs based on the absence of Li absorption in their spectra."
    },
    "0403/astro-ph0403401_arXiv.txt": {
        "abstract": "{ We present a statistical analysis of the {\\it Chandra} observation of the source field around the 3C 295 galaxy cluster (z=0.46) aimed at the search for clustering of X-ray sources. We applied three different methods of analysis, all suggesting a strong clustering in the field on scales of a few arcmin. In particular 1) the logN-logS computed separately for the four ACIS-I chips reveals that there is a significant ($3.2\\; \\sigma$ in the $0.5-2$ keV, $3.3\\; \\sigma$ in the $2-10$ keV and $4.0\\; \\sigma$ in the $0.5-10$ keV band) excess of sources to the North-North East and a void to the South of the central cluster.  2) the two point, two-dimensional Kolmogorov-Smirnov (KS) test, shows the probability that the sources are uniformly distributed is only a few percent. 3) a strong spatial correlation emerges from the study of the angular correlation function of the field: the angular correlation function (ACF) shows a clear signal on scales of $0.5\\div 5$ arcmin, correlation angle in the $0.5-7$ keV band $\\theta_0=8.5^{+6.5}_{-4.5}$, $90$\\% confidence limit (assuming a power law ACF with slope $\\gamma=1.8$). This correlation angle is $2$ times higher than that of a sample of $8$ ACIS-I field at the $2.5 \\; \\sigma$ confidence level. The above scales translate to 0.2$\\div$2 Mpc at the cluster redshift, higher than the typical cluster core radius, and more similar to the size of a ``filament'' of the large scale structure. ",
        "introduction": "N-body and hydrodynamical simulations show that clusters of galaxies lie at the nexus of several filaments of galaxies (see e.g. Peacock 1999, Dav\\`e et al. 2001 and references therein).  Such filaments map out the ``cosmic web'' of voids and filaments of the large scale structure (LSS) of the Universe. According to the same simulations these filaments contain a large fraction ($30-40\\%$) of the baryons in the Universe at z$<1$ (the remainder ending up in the hot gas in clusters of galaxies on one side, and in stars and cold gas clouds on the other side).  Despite its larger total mass, observations of the intergalactic matter in filaments have yielded so far only limited information, mostly due to its low density (most of the baryons in this phase should be at densities only 10-100 times higher than the average density in the Universe).  The most direct ways to detect a filament at low redshift is through its soft X-ray diffuse emission (see e.g. Zappacosta et al. 2002, Soltan, Freyberg \\& Hasinger 2002), or through soft X-ray and UV absorption line studies (see e.g. Fiore et al. 2000 and references therein, Nicastro et al. 2002, Nicastro et al. 2003, Mathur et al. 2003).  Both methods require very difficult observations, at the limit of the present generation of X-ray and UV facilities. Alternatively, filaments could be mapped out by galaxies (Daddi et al. 2001, Giavalisco \\& Dickinson 2001) and by the much more luminous Active Galactic Nuclei (AGNs), assuming that AGNs trace galaxies. Since rich clusters of galaxies are good indicators of regions of sky where filaments converge, numerous AGN searches around clusters of galaxies have been performed in the past. Several of these studies suggest overdensities of AGNs around distant clusters of galaxies (Molnar et al. 2002 for the cluster Abell 1995; Best et al. 2002 for MS1054-03; Martini et al. 2002 for Abell 2104; Pentericci et al. 2002 for the protocluster at $z \\sim 2.16$ around the radio galaxy MRC 1138-206; see also Almaini et al. 2003 for the ELAIS North field). Many of these studies have been performed in X-rays, since extragalactic X-ray sources, which are mostly AGNs, have a space density $\\sim 10$ times higher than optically selected AGNs (see Yang et al. 2003), and therefore provide denser tracing of LSS.  One of the first studies of the X-ray source population in cluster fields was performed by Cappi et al. (2001) who studied the {\\it Chandra} $8\\times 8$ arcmin ACIS fields around the distant cluster of galaxies RX J003033.2+261819 ($z=0.5$) and 3C 295 ($z=0.46$).  Cappi et al. (2001) reported the tentative detection of an overdensity of faint X-ray sources in a region of a few arcmin around both clusters, with respect to the average X-ray source density at the same flux limit.  However, the observations were too short ($\\sim 30$ ks and $\\sim 18$ ks, respectively), and the source samples consequently too small, to derive any more detailed conclusion.  In addition to counting sources in selected sky areas, the clustering of X-ray sources can be studied using other observational and statistical tools.  Gilli et al. (2003) report narrow spikes in the redshift distribution of the sources in the {\\it Chandra} Deep Field South (CDFS), indicating strong clustering of sources in these narrow redshift ranges.  If the source samples are sufficiently large (of the order of 100 sources or more) one derives more detailed and quantitative information on the source clustering by studying the angular correlation function (ACF) of the sources in the field. Vikhlinin \\& Forman (1995) were among the first to study the ACF in the X-ray band; they evaluated an average ACF from a large set of deep ROSAT observations, covering in total 40 deg$^2$ of sky. They found positive correlation on scales from a fraction of arcmin to tens of arcmin.  However, their best fit ``correlation angle'' $\\theta_0 \\sim 10$ arcsec, is smaller than the ROSAT PSPC Point Spread Function ($\\sim 25$ arcsec FWHM on-axis); as noted by the authors, this leads to an ``amplification bias'' and their measured ACF is consequently somewhat overestimated.  This effect can be greatly mitigated by using a telescope like {\\it Chandra}, whose on-axis PSF is only 0.5 arcsec FWHM, 50 times better than the PSPC on-axis PSF. So far a correlation analysis of {\\it Chandra} X-ray sources have been published only by Giacconi et al. (2000), who used the first 100 ks of observation of the CDFS, and by Yang et al. (2003) who analyzed a mosaic of 9 moderately deep (30ks) {\\it Chandra} pointings of the Lockman Hole area.  Yang et al. could study the source angular correlation on scales of several tens of arcmin, being limited on smaller scales by the small number of detected sources per unit area.  Our main goal in this paper is to push this kind of analysis toward a) smaller scales, of order of a few arcmin; and b) targeting a field where the likelihood of observing LSS is high, i.e. a field around a cluster of galaxies.  For these reasons we observed again with {\\it Chandra} for about 100ks the 3C295 field, where Cappi et al. (2001) found tentative evidence of an overdensity of faint X-ray sources in a region of a few arcmin around the central cluster.   The paper is organized as follows: Section 2 presents the observations and data reduction; Section 3 presents the results of our analysis; Section 4 discusses such results and draws our conclusions. ",
        "conclusions": "A $92$ ks {\\it Chandra} observation of the source field around the $z =0.46 $ 3C 295 cluster shows an excess of sources visible in the NE corner of the {\\it Chandra} observation (fig. 1). This is clear in the density contour plot of fig. 11. This figure shows that the denser region of the putative filament has a sources density more than $5$ times higher than the average source density of the field ($\\sim 0.5$ sources per arcmin$^2$. A chip by chip logN-logS analysis demonstrates that the NE chip has a $4.0 \\sigma$ excess of sources over the SW chip in the total ($0.5-7$ keV) {\\it Chanda} band. [In the standard soft ($0.5-2$ keV) band the excess is $3.2 \\sigma$, and in the standard hard ($2 - 7$ keV) band the excess is $3.3 \\sigma$.]  This result confirms the basic result of Cappi et al. (2001) and extends it to deeper fluxes, larger field-of-view and the $2-10$ keV band.   \\begin{figure} \\centering \\includegraphics[angle=0,width=9cm]{cont3.ps}% \\caption{The density profile of the 3C 295 field, computed for the whole $0.5 - 7$ keV band. The linear smoothing factor is $1.5$ arcmin, and the four contour levels indicate source densities of $1.3$, $1.8$, $2.2$ and $2.7$ sources per arcmin$^2$. The clustering of sources in the NE corner is clearly visible; the bright central spot is the 3C 295 cluster. } \\label{spe1} \\end{figure}     The asymmetric distribution of the 3C 295 field is confirmed by two analyses: (1) the two dimensional Kolmogorov-Smirnov test, which show the probability that the sources are uniformly distributed is $0.632 \\%$ in the total {\\it Chandra} band [ $3.09\\,\\%$ for the soft band and $3.84\\;\\%$ for the hard band).  (2) the two point angular correlation function (ACF) strongly indicate positive correlation, for scales of $0.5-5$ arcmins. This strong correlation  has not been found in a sample of eight ACIS-I fields with a similar exposure time, for which the correlation angle $\\theta_0$ is smaller by a factor of $\\gs2$ than that of 3C 295 at the $2.5\\;\\sigma$ confidence level, suggesting that the 3C 295 overdensity of sources in the NE chip is truly peculiar. An intriguing explanation, to be confirmed when the redshifts of the sources are measured, is that the overdensity of sources is actually related to a cosmic filament of the  LSS.   Cappi et al. (2001) discussed four possible causes for this `surplus' of sources: (1) gravitationally lensed very faint sources; (2) rapid evolution of cluster AGN or starburst galaxies; (3) cosmic variance of background sources; (4) LSS associated with the clusters.  Since the surplus sources are not symmetrically place around the cluster, our results rule out an enhanced AGN population in the 3C~295 cluster itself, and lensing by the cluster potential. Since N-body and hydrodynamical simulations and galaxy surveys (e.g 2dF and Sloan) lead us to expect that clusters of galaxies lie at the nexus of several filament, we believe that this excess is likely to represent a filament of the LSS of the Universe converging onto the 3C 295 cluster.  The redshifts of the sources making up the excess are not yet known. However, if we assume that the sources are associated with the 3C 295 cluster, then we can use the cluster redshift to estimate some of their properties.  Adopting a concordance cosmology ($H_0 =65$ km/s Mpc, $\\Omega_M = 0.3$ and $\\Omega_{\\Lambda} = 0.7$, Spergel et al. 2003) we obtain luminosities in the range $7.5\\times 10^{41} \\div 1.1 \\times 10^{44}$ ergs s$^{-1}$ (median $= 3.1 \\times 10^{42}$ ergs s$^{-1}$) for the total {\\it Chandra} band.  [$2.8\\times 10^{41} \\div 3.8\\times10^{43}$ ergs s$^{-1}$ (median $= 1.1 \\times 10^{42}$ ergs s$^{-1}$) for the soft band; $1.5\\times 10^{42} \\div 5.0 \\times 10^{43}$ ergs s$^{-1}$ (median $= 4.5 \\times 10^{42}$ ergs s$^{-1}$) for the hard band.] These are moderate, Seyfert galaxy, luminosities, extending down to the range of luminous starburst galaxies at the faint end of the soft band.  An association with 3C 295 also defines a spatial scale of $2$ Mpc (5 arcmin) for the adopted cosmology.  For an, admittedly unlikely, spherical region of volume of $\\sim 4 $ Mpc$^{3}$ containing the excess sources, the implied space density of the excess sources (i.e. subtracting the CDF-S defined logN-logS contribution) is $0.9$ Mpc$^{-3}$ ($0.5-2$ keV) and $0.8$ Mpc$^{-3}$ ($2-10$ keV). This is well above the normal maximum AGN space density at $z=0.5$ ($\\sim 10^{-4}$ Mpc$^{-3}$, for luminosities down to $10^{42}$ ergs s$^{-1}$ in the 0.5-2keV band, Hasinger 2003, and $10^{43}$ ergs s$^{-1}$ in the 2-10 keV band Fiore et al. 2003). However, since the X-ray source luminosities are as low as $\\sim 2 \\times 10^{41}$ ergs s$^{-1}$ (at the redshift of 3C 295), the little studied lower end of the XLF is being probed, and contributions from starburst and even normal galaxies can be important.  In fact, the integral of the field galaxy luminosity function at z=0.5 (e.g. Poli et al. 2001) gives $\\sim 0.13$ galaxies Mpc$^{-3}$ for $M_B < -17$  (a resonable faint end optical luminosity, corresponding to our lower X-ray luminosities). If we assume that roughly one tenth of the galaxies are active X-ray sources of $L_X > 3 \\times 10^{41}$ ergs s$^{-1}$, then we would expect $\\sim0.013$ X-ray sources Mpc$^{-3}$. Since we count $\\sim0.9$ sources Mpc$^{-3}$, this implies a galaxy overdensity of $\\approx70$, with of course a large (factor of 2-4) positive and negative uncertainty, because of the uncertainties in our space densities and assumptions. Still, this is intriguingly close to the expected galaxy overdensity of filaments $\\sim 10 \\div 10^2$, and much smaller than the overdensities of clusters of galaxies ($\\sim 10^3 \\div  10^4$), (the density contrast at the virial radius being $\\sim200$). Furthermore, in a filamentary structure the implied space density will depend strongly on the orientation of the filament to our line of sight.  We stress again that these arguments hold only under the assumption that the redshift of the excess sources is the same of that of the 3C 295 cluster. In order to determine if this excess is associated or not to the 3C 295 cluster, we need to know the redshift of the sources via optical identifications.  For this reason, in order to study this candidate filament, we are pursuing new {\\it Chandra} and optical observations to map  out the 3C~295 region and delimiting the filament properties up to scales of $\\sim 24$ arcmin (i.e. $\\sim 6$ Mpc) from the 3C 295 cluster. These studies are important because they may open-up a new way to map high-density peaks of LSS at high redshifts with high efficiency."
    },
    "0403/astro-ph0403637_arXiv.txt": {
        "abstract": "{$uvby$(--$\\beta$) photometry has been obtained for an additional 411 very metal-poor stars selected from the HK survey, and used to derive basic parameters such as interstellar reddenings, metallicities, photometric classifications, distances, and relative ages.  Interstellar reddenings adopted from the Schlegel et al.~(1998) maps agree well with those from the intrinsic-color calibration of Schuster \\& Nissen (1989).  [Fe/H] values are obtained from the CaII K line index of the HK survey combined with the $uvby$ and $UBV$ photometry.  The $c_{\\rm 0},(b-y)_{\\rm 0}$ diagram is seen to be very useful for classifying these very metal-poor field stars into categories similar to those derived from globular cluster color-magnitude diagrams; it is found that the HK survey has detected metal-poor candidates extending from the red-giant to the blue-horizontal branch, and from the horizontal branch to subluminous stars. Distances derived from $UBV$ photometry agree reasonably well with those from $uvby$, considering the paucity of good calibrating stars and the extrapolations required for the most metal-poor stars.  These very metal-poor stars are compared to M92 in the $c_{\\rm 0},(b-y)_{\\rm 0}$ diagram, and evidence is seen for field stars 1--3 Gyrs younger than this globular cluster; uncertainties in the [Fe/H] scale for M92 would only tend to increase this age difference, and significant reddening uncertainties for M92 are unlikely but might decrease this difference.  The significance of these younger very metal-poor stars is discussed in the context of Galactic evolution, mentioning such possibilities as hierarchical star-formation/mass-infall of very metal-poor material and/or accretion events whereby this material has been acquired from other (dwarf) galaxies with different formation and chemical-enrichment histories. ",
        "introduction": "Over the past two decades, our collective knowledge of the nature of the thick disk and halo of the Galaxy has expanded enormously, due primarily to the impact of several ongoing large-scale survey efforts carried out to detect and analyze metal-poor stars. These include the HK survey of Beers and collaborators (Beers, Preston, \\& Shectman 1992; Beers 1999) and the Hamburg/ESO stellar survey of Christlieb \\& collaborators (Christlieb 2003), both of which select stars with objective-prism techniques, and hence introduce no kinematic bias into their samples. Such biases are present (and must be corrected for) in proper-motion selected survey samples, such as the exhaustive previous studies of, e.g., Ryan \\& Norris (1991) and Carney et al.~(1996). The prism-survey selected samples are hence well suited for studies of the kinematics and dynamics of the old stellar populations of the Milky Way, in particular because of the burgeoning databases of proper motion information that are presently being assembled from a variety of sources (e.g., UCAC2: Zacharias 2002; SPM: Girard et al.~2003). In order to make optimal use of the proper motions for kinematic analyses, accurate stellar classifications, and photometrically determined distances, are crucial.  The $uvby$--$\\beta$ photometric system is particularly suited for the study of very-metal-poor (hereafter, VMP) F- and G-type stars, as has already been pointed out in Paper VIII by Schuster et al.~(1996; hereafter S96).  Briefly, intrinsic-color calibrations, $(b-y)_{\\rm 0}$--$\\beta$, exist which allow accurate and precise, $\\pm 0\\fm01$, measures of interstellar reddening excesses, $E(b-y)$, for individual field stars; such a calibration has been given by Schuster \\& Nissen (1989).  Photometric absolute magnitudes and distances can be calibrated and used effectively, as shown in the papers by Olsen (1984) and Nissen \\& Schuster (1991).  This photometric system has the great advantage that it permits us to obtain accurate stellar distances even for evolving main-sequence and subgiant stars due to the gravity sensitivity of the $c_{\\rm 0}$ index.  Also, importantly, theoretical isochrones in the $M_{\\rm bol}$, $T_{\\rm eff}$ diagram can be transformed to the $M_{\\rm V}$, $(b-y)_{\\rm 0}$ or $c_{\\rm 0}$, $(b-y)_{\\rm 0}$ diagrams for the estimation of relative and/or absolute ages of evolving field stars which are near their respective turn-offs, and in several of the previous papers of this series the isochrones of VandenBerg et al.~have been used for such purposes, to study the Galactic halo population and to make comparative analyses between the relative ages of the halo and thick-disk stellar populations.  Most recently the isochrones of Bergbusch \\& VandenBerg (2001) have been transformed to the $uvby$ photometric system using the color--$T_{\\rm eff}$ relations of Clem et al.~(2003).  Also, the $uvby$--$\\beta$ photometry can provide basic stellar atmospheric parameters as a prelude to detailed chemical abundance studies making use of high-resolution spectroscopy and model atmospheres.  Several empirical calibrations already exist in the literature for the conversion of $(b-y)_{\\rm 0}$ or H$\\beta$ to $T_{\\rm eff}$; these calibrations include appropriate metallicity dependences.  Index diagrams, such as $c_{\\rm 0}$, $(b-y)_{\\rm 0}$, or the reddening-free $[c_{\\rm 1}]$, $[m_{\\rm 1}]$, or $[c_{\\rm 1}]$, $\\beta$, allow the classification of field stars according to their evolutionary status, permiting us to estimate the stellar surface gravities, also for input into the model-atmosphere analyses.  In this paper, $uvby$--$\\beta$ photometry is presented for an additional 411 VMP stars from the HK survey, providing a total database of such photometry for 497 VMP stars, when combined with the data of S96.  For the present sample the stars have been selected with [Fe/H] $\\la -1.5$, and 243 were observed in M\\'exico using classical photometric (photoelectric) techniques and 177 in Chile using DFOSC (CCD) techniques.  In Sect.~2 the observing and reduction techniques are described briefly, the catalogues of new $uvby$--$\\beta$ data presented, and the $V$ magnitudes and $(b-y)$ colors from the $uvby$ observations compared to magnitudes and $(B-V)$ from the HK survey.  In Sect.~3, the photometry is dereddened using a modification of the Schlegel et al.~(1998) reddening maps and also the intrinsic-color calibration of Schuster \\& Nissen (1989); reddenings from the two methods, $E(B-V)$ and $E(b-y)$, are compared.  In Sect.~4, [Fe/H] values are derived for the VMP stars using the techniques developed in the HK survey, and probable carbon-enhanced stars are identified based on a comparison of the GP and KP indices.  Photometric classifications are derived for the VMP stars in Sect.~5 using the $c_{\\rm 0}$, $(b-y)_{\\rm 0}$ diagram.  Stars are found covering a wide range of stellar types from the horizontal branch (HB) to subluminous stars (SL), and from the red giant stars (RG) to the blue horizontal branch (BHB), and other categories include main-sequence (MS), turn-off (TO), subgiant (SG), blue-straggler (BS), and red-horizontal-branch-asymptotic-giant-branch (RHB-AGB) stars. Possible abundance anomalies for some VMP stars have been identified from the $uvby$ photometric indices and diagrams, such as the $[c_{\\rm 1}]$, $[m_{\\rm 1}]$; for example, ten probable Am stars have been found and also a number of possible AGB stars with unusual chemical abundance ratios or binary companions.  Distance estimates are made for the VMP stars in Sect.~6 using $uvby$ photometry plus various methods and new calibrations, and also using the $UBV$ photometry and techniques developed in the HK survey. Comparisons of these photometric distances show reasonably good agreement, considering the paucity of calibrating stars and extrapolations required for the more VMP stars.  In Sect.~7, the VMP field stars are compared to the globular cluster M92 in the $c_{\\rm 0}$, $(b-y)_{\\rm 0}$ diagram, using the isochrones of Bergbusch \\& VandenBerg (2001), as transformed to $uvby$ by Clem et al.~(2003), to interpolate relative and absolute ages.  A number of VMP stars apparently 1--3 Gyrs younger than M92 are noted, and their importance for understanding the formation and evolution of the Galactic halo discussed. ",
        "conclusions": "\\begin{enumerate} \\item The overall VMP HK-survey sample contains a wide range of stellar types, ranging from horizontal branch stars to subluminous, and from red giant stars to the blue horizontal branch. \\item The dereddened $c_{\\rm 0},(b-y)_{\\rm 0}$ diagram has been shown to be quite useful for providing photometric classifications of the VMP stars analogous to types derived from GC color-magnitude diagrams, such as Turn-Off stars (TO), SubGiants (SG), Red Giants (RG), Horizontal Branch stars (HB), Blue Horizontal Branch stars (BHB), Blue Stragglers (BS), SubLuminous stars (SL), and so forth (see Fig.~6). \\item The intrinsic-color calibration of Schuster \\& Nissen (1989), as modified slightly by Nissen (1994), is shown to provide reddening excesses, $E(b-y)$ or $E(B-V)$, very similar to the adopted reddening estimates derived in this publication from the maps of Schlegel, Finkbeiner, \\& Davis (1998) (see Eq.~1).  No significant systematic offsets between these two dereddening techniques are noted (see Fig.~3). \\item A number of VMP stars have been noted with probable anomalous photometric traits, especially from the $m_{\\rm 1}$ and $[m_{\\rm 1}]$ indices; two such groups stand out.  First, there are several stars with $(b$--$y)_{\\rm 0} \\la 0\\fm45$ and with $m_{\\rm 0} \\ga 0\\fm17$, much larger than would be expected for VMP stars with [Fe/H] $\\la -1.5$.  Most of these have been classified SG, and some show clear evidence of photometric variability.  These are perhaps analogous to stars discussed in S96 with larger than expected $[m_{\\rm 1}]$ values.  We suggest here that these are misclassified AGB stars with unusual chemical-abundance ratios, photometric variability, and/or binary companions. \\item The second group of anomalous stars are those ten classified BS and having $m_{\\rm 1}$, $[m_{\\rm 1}]$, and $(U-B)_{\\rm 0}$ values indicating nearly solar [Fe/H] values.  There is a clear discrepancy here between these photometric indices and the KP index used to derive [Fe/H] for the HK survey.  These stars are very similar to the Am stars identified by Wilhelm et al.~(1999a, 1999b) and have been noted as ``BS (Am)'' in Table 5. \\item The photometric distances from the $UBV$ and $uvby$ photometries agree reasonably well considering the problems, lack of calibrating stars, and extrapolations needed for the more VMP stars.  Our Hipparcos-based, photometric calibration for $M_{\\rm v}$ seems to work quite well for the turn-off, main-sequence, and subgiant VMP stars, as suggested in Figs.~7 and 8. \\item In the $c_{\\rm 0},(b-y)_{\\rm 0}$ diagram, the youngest VMP stars appear to have ages 1--3 Gyrs younger than the GC M92.  Uncertainties in the [Fe/H] scale for M92 would tend to increase this age difference even more.  (The interstellar reddening of M92 seems to be well determined but might be as uncertain as $\\pm 0\\fm01$).  Such younger VMP stars are showing evidence for important details upon the overall formation and evolution of the Galaxy, such as possible hierarchical star-formation/mass-infall for the VMP material, and/or accretion processes from other (dwarf) galaxies with different formation and chemical-enrichment histories.  \\end{enumerate}"
    },
    "0403/astro-ph0403547_arXiv.txt": {
        "abstract": "{ We considered the structure of steady--state plane--parallel radiative shock waves propagating through the partially ionized hydrogen gas of temperature $T_1 = 3000$~K and density $10^{-12}~\\gcc\\le\\rho_1\\le 10^{-9}~\\gcc$. The upstream Mach numbers range within $6\\le M_1\\le 14$. In frequency intervals of hydrogen lines the radiation field was treated using the transfer equation in the frame of the observer for the moving medium, whereas the continuum radiation was calculated for the static medium. Doppler shifts in Balmer emission lines of the radiation flux emerging from the upstream boundary of the shock wave model were found to be roughly one--third of the shock wave velocity: $-\\dv\\approx \\frac{1}{3}U_1$. The gas emitting the Balmer line radiation is located at the rear of the shock wave in the hydrogen recombination zone where the gas flow velocity in the frame of the observer is approximately one--half of the shock wave velocity: $-V^*\\approx\\frac{1}{2} U_1$. The ratio of the Doppler shift to the gas flow velocity of $\\dv/V^* \\approx 0.7$ results both from the small optical thickness of the shock wave in line frequencies and the anisotropy of the radiation field typical for the slab geometry. In the ambient gas with density of $\\rho_1\\ga 10^{-11}~\\gcc$ the flux in the $\\Ha$ frequency interval reveals the double structure of the profile. A weaker $\\Hb$ profile doubling was found for $\\rho_1\\gtrsim 10^{-10}~\\gcc$ and $U_1\\lesssim 50~\\kms$. The unshifted redward component of the double profile is due to photodeexcitation accompanying the rapid growth of collisional ionization in the narrow layer in front of the discontinuous jump. ",
        "introduction": "It is now a well--established fact that hydrogen emission lines are a characteristic feature of radially pulsating stars of various types. Strong hydrogen emission is observed in Mira type \\citep{Joy:1947,Joy:1954}, W~Vir \\citep{Abt:1954,Wallerstein:1959} and RV~Tau \\citep{Preston:1962} pulsating variables. Moreover, in the spectra of Mira stars Balmer emission lines persist during the major part of the pulsation period \\citep{Joy:1947,Richter:Wood:2001}. The intensity of the hydrogen emission seems to correlate with the amplitude of the pulsation since RR~Lyr variables exhibit only weak emission lines \\citep{Preston:1964} and in classical Cepheids (for example in $\\beta$~Dor) hydrogen emission lines are scarcely detected \\citep{Hutchinson:1975}. The hydrogen emission in the spectra of pulsating stars is thought to be due to radiative cooling of the gas compressed by the shock wave propagating through the stellar atmosphere in each pulsation cycle \\citep{Kraft:1959,Wallerstein:1959,Abt:Hardie:1960,Gorbatskii:1961}.  High resolution spectroscopy reveals the doubling of $\\Ha$ and $\\Hb$ emission profiles, whereas profiles of higher Balmer lines exhibit only the asymmetry. This feature is observed not only in Mira stars \\citep{Bidelman:Ratcliffe:1954,Fox:Wood:Dopita:1984,Gillet:1988,Woodsworth:1995} but also in W~Vir and RV~Tau variables \\citep{Lebre:Gillet:1991,Lebre:Gillet:1992}. \\citet{Bidelman:Ratcliffe:1954} explained the $\\Ha$ profile doubling observed in the Mira type star T~Cen as an absorption reversal rather than a real duplicity, the absorption resulting from the cool hydrogen gas above the propagating shock wave. \\citet{Willson:1976} interpreted the double structure of the emission profiles in terms of a spherically symmetric shock wave with a radial distance from the center of the star appreciably larger than the radius of the photosphere. According to this model the flux of the redshifted component emerges from the back side of the shock wave moving outward from the observer. \\citet{Woodsworth:1995} modelled double $\\Ha$ profiles as three emission components of equal width, two of which are blended. Thus, different phenomenological models demonstrate the ambiguity existing so far in our understanding of the origin of the Balmer emission lines produced by the shock waves in stellar atmospheres.  This paper is the fifth in our series on the structure of radiative shock waves. In our previous Papers I--IV \\citep{Fadeyev:Gillet:1998,Fadeyev:Gillet:2000,Fadeyev:Gillet:2001,Fadeyev:2002} we presented the method of computation and described results for the structure of radiative shock waves propagating through partially ionized  hydrogen gas with temperature and density typical for atmospheres of pulsating late--type stars. An advantage of this approach is that the gas dynamics, radiation field and atomic level populations are considered self--consistently for the whole shock wave model. However, in all our previous papers the radiative transfer was treated in a static medium approximation, so that we were unable to compare the calculated monochromatic radiation flux with observed emission profiles. Below we describe the shock wave models with Doppler shifts in the line profiles computed from the transfer equation in the frame of the observer. ",
        "conclusions": "In our attempts to solve the transfer equation for the shock wave structure in the co--moving frame we encountered a severe difficulty because of the numerical instability arising at the velocity discontinuity. In the present study this difficulty could be circumvented because the role of the radiation field in spectral lines is quite small in comparison with that of the continuum. According to our estimates the Doppler shifts in the hydrogen lines do not affect perceptibly either the structure or the radiative losses of the shock wave. In particular, the emergent flux integrated over the line frequency interval was found to be the same within $\\lesssim 1$\\% for both the static and the moving medium. This allowed us to leave out the effects of the Doppler shifts from the global iteration procedure and to solve the transfer equation in the frame of the observer only in the final iteration.  The most remarkable result of our study is that the shock wave models show the double emission structure in the $\\Ha$ and $\\Hb$ profiles of the emergent radiation flux which is well known from high resolution high signal--to--noise ratio spectroscopy. We showed that the redward emission feature results from the narrow layer just ahead of the discontinuous jump within which the growth of collisional ionization is accompanied by photodeexcitation onto the second atomic level. The contribution of preshock photodeexcitation decreases with decreasing density of the ambient gas, therefore the double emission feature can be considered as a tool for diagnostics of stellar atmospheres with propagating shock waves. Here one should bear in mind, however, that in the framework of our model the unpertubed gas is at rest with respect to the observer, whereas in pulsating stars the gas ahead of the shock wave falls down onto the star with velocity in the range of one to a few dozen $\\kms$. Thus, the preshock emission feature in $\\Ha$ and $\\Hb$ profiles should be observed as a redshifted component.  Another important conclusion is that the velocity inferred from Doppler shifts of Balmer lines is roughly one--third of the shock wave velocity: $\\dv\\approx \\frac{1}{3} U_1$. This is due to the fact that the gas layers emitting the Balmer line radiation are located at the rear of the shock wave in the hydrogen recombination zone where the velocity in the frame of the observer is roughly one half of the shock wave velocity: $-V^*\\approx\\frac{1}{2}U_1$. The ratio of the Doppler shift to the gas flow velocity of $\\dv/V^*\\approx 0.7$ results from the small optical thickness of the shock wave model and the anisotropy of the radiation field produced by the shock wave.  $\\Ha$ is the broadest emission line with FWHM comparable to the velocity of the shock wave $U_1$. This, as well as the significant contribution from the preshock photodeexcitation zone make this emission line less appropriate for inferring the velocity of the shock wave from observationally measured Doppler shifts. However the width of the emission profile is proportional to the temperature of the hydrogen atoms behind the discontinuous jump $\\Ta^+$ which is related to the shock wave velocity $U_1$ via the Rankine--Hugoniot relations. Thus, the width of the Balmer line is a function of two general quantities: the shock wave velocity $U_1$ and the ambient gas density $\\rho_1$. In particular, the FWHM of $\\Ha$ decreases by a factor of two with decreasing gas density within $10^{-9}~\\gcc\\le\\rho_1\\le 10^{-12}~\\gcc$.  Though our results are consistent with the observations, there is a number of other parameters that determine the propagation of the shock wave in the atmospheres of pulsating stars. In particular, our model is confined to a flat finite slab and therefore does not take into account the cool hydrogen gas of the outer stellar atmosphere above the propagating shock wave. Thus, we cannot exclude the role of absorption in the formation of the double emission structure in Balmer lines. The presence of such absorption in emission profiles observed in RV~Tau and W~Vir stars was pointed out, for example, by \\cite{Lebre:Gillet:1991,Lebre:Gillet:1992}."
    },
    "0403/hep-ph0403134_arXiv.txt": {
        "abstract": "We show that the non-standard neutrino interactions can play a role as sub-leading effect on the solar neutrino oscillations. We observe that very small flavor universality violations of order of 0.1-0.2 $G_F$ is sufficient to induce two phenomena: suppression of the $\\nu_e$-earth regeneration and a shift of the resonance layer in the sun. We obtain these phenomena even in the absence of any flavor changing interactions. We discuss their consequences and confront with a global analysis of solar+KamLAND results. We conclude that a new compatibility region in the $\\Delta m^2 \\times \\tan^2 \\theta_{\\odot}$ , which we call very low Large Mixing Angle region is found  for $\\Delta m^2 \\sim 10^{-5}$ eV$^2$ and $\\tan^2 \\theta_{\\odot}= 0.45$. ",
        "introduction": "In the last years, the discovery of neutrino oscillation in solar and reactor experiments selected as a more probable explanation to the solar neutrino problem the so called Large Mixing Angle (LMA) MSW solution. The SNO~\\cite{sno,sno-all} and the KamLAND~\\cite{Kam} experiments confirm and refine the trend of the evidences of neutrino oscillations due the solar neutrino observations, as measured by Homestake~\\cite{Cl}, SAGE~\\cite{sage}, GALLEX~\\cite{gallex}, GNO~\\cite{gno} and Super-Kamiokande~\\cite{SK,SK2}. As a result, the solar oscillation parameters have pinned down to $ 6\\times 10^{-5}$ eV$^2 <\\Delta m^2 < 1\\times 10^{-4}$ eV$^2$ and $0.3 < \\tan^2 \\theta_{\\odot} < 0.55$ at $2\\sigma$~\\cite{us:msw2}. Several analyzes have arrived to same conclusions~\\cite{balan,fogli,valle2,alia,crem,choubey}.  In a more general context, sub-leading effects can change this picture, which motivate us to investigate the robustness of the determination of the solar parameters. In this letter, we assume that non-standard neutrino interactions, which we parameterized by two parameters $\\epsilon'$ and $\\epsilon$, are  present, relaxing the allowed region of the parameters. In the presence of non-standard neutrino interactions, we have found that the allowed interval for $\\Delta m^2$ increases, rescuing the very low LMA region, $\\Delta m^2 \\sim 1\\times 10^{-5}$ eV$^2$, and the high part of LMA region, $\\Delta m^2 \\sim 2\\times 10^{-4}$ eV$^2$, respectively due the suppression of earth matter and due to a $\\Delta m^2$ shift induced by a non-zero $\\epsilon'$. ",
        "conclusions": "We showed that NSNI will affect the fit in the LMA region of the MSW solution to the solar neutrino anomaly. When one takes into account the KamLAND results, positive values of the $\\epsilon'$  push the allowed region of the neutrino parameters $\\Delta m^2$ and $\\tan^2\\theta_{\\odot}$ at 95\\% C.L. from pure MSW low-LMA and high-LMA to a completely new region in which $\\Delta m^2$ is lower than the previous two ones, which we call very-low-LMA. If one chooses $\\epsilon'<0$, the preferred allowed region tends to higher values of $\\Delta m^2$.  Almost all our conclusions below are independent of specific sources of the non-standard neutrino interactions, that could be present in interactions with d-quarks, u-quarks or electrons.  We have found that the main effects of the presence of the NSNI interactions are:  \\begin{itemize}  \\item{displacement of low-LMA region to lower (higher) values of $\\Delta m^2$, for a positive (negative) value of $\\epsilon'$.}  \\item{suppression of Earth regeneration at $\\Delta m^2\\sim 10^{-5}$ eV$^2$ for positive values of $\\epsilon'$.}  \\item{Due to suppression of Earth regeneration, appearance of a new region of compatibility between solar and KamLAND data around $\\Delta m^2\\sim 10^{-5}$ eV$^2$,  with no spectrum distortion for the low-energy SK and SNO data.}  \\item{improvement of high-LMA fit quality for positive values of $\\epsilon'$}  \\item a 1 kton-yr of KamLAND can make a strong statement about the existence of non-standard neutrino interactions. The striking signal of this NSNI would be the location of the prefered oscillation parameters in the very low or in the high LMA region. \\end{itemize}  {\\em Note added:} When we were finishing our paper, an article by Friedland, Lunardini and Pe\\~na-Garay (hep-ph/0402266) appeared, which discusses topics similar to the ones discussed in our paper, where we discuss not only the non-standard neutrino interaction induced by d-quarks case as well the u-quarks and electrons. Also we made a quantitative statement about the role of more statistics on KamLAND experiment, combined with the present solar neutrino data, to put more restrictive bounds on non-standard neutrino interactions."
    },
    "0403/hep-ph0403299_arXiv.txt": {
        "abstract": "\\noindent We analyze the uncertainties involved in obtaining the injection spectra of UHECR particles in the top-down scenario of their origin. We show that the DGLAP $Q^2$ evolution of  fragmentation functions (FF) to $Q=M_X$ (mass of the X particle) from their initial values at low $Q$ is subject to considerable uncertainties. We therefore argue that, for $x\\lsim 0.1$ (the $x$ region of interest for most large $M_X$ values of interest, $x\\equiv 2E/M_X$ being the scaled energy variable), the FF obtained from DGLAP evolution is no more reliable than that provided, for example, by a simple Gaussian form (in the variable $\\ln(1/x)$) obtained under the coherent branching approach to parton shower development process to lowest order in perturbative QCD. Additionally, we find that for $x\\gsim0.1$, the evolution in $Q^2$ of the singlet FF, which determines the injection spectrum, is ``minimal'' --- the singlet FF changes by barely a factor of 2 after evolving it over $\\sim$ 14 orders of magnitude in $Q\\sim M_X$. We, therefore, argue that as long as the measurement of the UHECR spectrum above $\\sim10^{20}\\ev$ is going to remain uncertain by a factor of 2 or larger, it is good enough for most practical purposes to directly use any one of the available initial parametrisations of the FFs in the $x$ region $x\\gsim0.1$ based on low energy data, without evolving them to the requisite $Q^2$ value. ",
        "introduction": "One of the main problems in understanding the origin of the observed Ultra-High Energy Cosmic Ray (UHECR) events with energy $E\\gsim10^{20}\\ev$\\cite{uhecr_obs} --- below we will sometimes refer to these as Extreme Energy Cosmic Ray (EECR) events --- is the difficulty of producing such enormously energetic particles in astrophysical environments by means of known acceleration mechanisms. There are but a few astrophysical objects --- among which are, perhaps, Gamma Ray Burst (GRB) sources and a class of powerful radio galaxies --- where protons can in principle be accelerated to requisite energies (at source) of $\\gsim10^{21}\\ev$ by the standard diffusive shock acceleration mechanism albeit with optimistic assumptions on the values of the relevant parameters. However, even for these objects, their locations and spatial distributions are not easy to reconcile with the observed spectrum and large-scale isotropy of the UHECR particles. (For recent reviews on astrophysical source origin of EECR see, for example, Refs.~\\cite{springer_book,torres_rev}).  \\para An alternative mechanism of producing the EECR particles is provided by the so-called ``top-down'' (TD) scenario (see \\cite{physrep} for a review) in which the EECR particles are envisaged to result from {\\it decay} of some sufficiently massive particles, generically called ``X'' particles, of mass $M_X\\gg10^{20}\\ev$, which could originate from processes in the early Universe. This is in contrast to the conventional ``bottom-up'' scenario in which {\\it all} cosmic ray particles including the EECRs are thought to be produced through processes that accelerate particles from low energies to the requisite high energies in suitable astrophysical environments.  \\para The X particles of the TD scenario, if at all they exist in Nature, are most likely to be associated with some kind of new physics at some sufficiently high energy scale that could have been realized in an appropriately early stage of the Universe. Two possibilities for the origin of the X particles have been discussed in the literature: They could be short-lived particles released in the Universe today from cosmic topological defects such as cosmic strings, magnetic monopoles, etc.~\\cite{td_book} formed in a symmetry-breaking phase transition in the early Universe. Alternatively, they could be some metastable (and currently decaying) particle species with lifetime larger than or of the order of the age of the Universe.  \\para Since the mass scale $M_X$ of the hypothesized X particle is well above the energy scale currently available in accelerators, its primary decay modes are unknown and likely to involve elementary particles and interactions that belong to unknown physics beyond the Standard Model (SM). However, irrespective of the primary decay products of the X particle, the observed UHECR particles must eventually result largely from ``fragmentation'' of the Standard Model quarks and gluons, that come from the primary decay products of the X particles, into hadrons. The most abundant final observable particle species in the TD scenario are expected to be photons and neutrinos from the decay of the neutral and charged pions, respectively, created in the parton fragmentation process, together with a few percent baryons (nucleons). The injection- or the source spectra of various species of UHECR particles (nucleons, photons and neutrinos) in this TD scenario are thus ultimately determined by the physics of the parton fragmentation process. The final observable UHECR particle spectra are determined by further processing of these injection spectra due to extragalactic and/or Galactic propagation effects depending on where the X particle decay takes place. Clearly, in order to test the predictions of the TD scenario against UHECR experimental data, it is crucial to be able to reliably calculate the injection spectra of various UHECR particles in this scenario. This is the subject we concern ourselves with in this paper.  \\para The problem at hand is essentially the same as determining the single-particle inclusive spectrum of hadrons produced, for instance, in the process $e^+e^-\\to \\gamma/Z\\to q\\bar{q}\\to {\\rm hadrons}$ (see, for example, \\cite{ellis_book}). The primary quarks produced in the collision would in general not be on-shell and would have large time-like virtuality $Q\\sim \\sqrt{s}$, the center-of-mass energy of the process. Each quark would, therefore, reduce its virtuality by radiating a gluon, the latter in turn splitting into a $q\\bar{q}$ pair or into two gluons, and so on. This process gives rise to a parton shower whereby at each stage a virtual parton splits into two other partons of reduced virtualities. This process of parton shower development is well-described by perturbative QCD until the virtuality reduces to $Q=Q_{\\rm hadron}\\sim 1\\gev$ when non-perturbative effects come into play binding partons into colorless hadrons. In the end, the link between partons and hadrons is quantitatively described in terms of fragmentation functions (FFs) $D_a^h(x,Q)$, which give the probability that a parton $a$ produced with an initial virtuality $Q=\\sqrt{s}$ produces the hadron $h$ carrying a fraction $x\\equiv 2E/\\sqrt{s}$ of the energy of $a$ ($E$ being the energy of the hadron)\\footnote{At high energies $E$ of our interest throughout this paper we shall assume $E\\simeq p$, the momentum of the particle.}. The final single particle inclusive spectrum of hadrons is given by a convolution of these FFs with the production probabilities of the primary partons (see next section).  \\para In the same way, the problem of determining the injection spectrum of UHECR particles from the decay of X particles essentially reduces to determining the FFs $D_a^h(x,M_X)$ for various hadron species $h$ (pions, nucleons) where $a$ represents the primary partons to which the X particle decays. (Actually, in our present case, we will be interested only in the so-called ``singlet'' FF corresponding to a sum over all partons $a$ as explained later).  \\para Clearly, the FFs themselves cannot be directly calculated from first principles entirely within perturbative QCD without extra assumptions about the nature of the non-perturbative process of formation of hadrons from partons. Several different approaches have been taken in the recent literature for evaluating the relevant FFs, which are discussed below.  \\para In this paper, we critically examine one of the approaches of evaluating the relevant FFs, namely, the DGLAP evolution equation method~\\cite{fodor, sarkar-toldra,barbot-drees,barbot_thesis,aloisio}, that has been widely used in recent calculations of the UHECR injection spectra in the TD scenario. We discuss the inherent uncertainties involved in this approach in calculating the relevant FFs over the ranges of $x$ and $M_X$ of interest. We also compare the FFs so obtained with those given by a simple analytical expression (given by a Gaussian in the variable $\\ln(1/x)$ as discussed later) obtained within the context of an analytical approach, namely, the coherent branching formalism, to lowest order in perturbative QCD~\\cite{ellis_book}, this analytical approach being valid only under ``small\" $x$ and ``large\" $Q$ approximation. We show that except for ``large\" $x\\gsim 0.1$, the uncertainties involved in obtaining the relevant FFs by numerical solution of the DGLAP evolution equation do not allow much significant advantage of using this numerical method over the simple analytical (but approximate) formula for FFs provided by the coherent branching approach. At the same time, we also find that, in the region $x\\gsim 0.1$, the evolution (in $Q$) of the {\\it singlet} FFs (which is what we are interested in) is very little --- the singlet FF changes by only a factor of 2 or so after evolving it over $\\sim$ 14 orders of magnitude in $Q\\sim M_X$. We explain the reason for this, and argue that, as long as the measurement of the EECR spectrum is going to remain uncertain within a factor of 2 or larger (which is likely to be the case in the foreseeable future), it is good enough for most practical purposes to directly use any one of the available parametrisations of the FFs in the $x$ region $x\\gsim0.1$ based on low energy (say at the Z-pole) data from $e^+e^-\\to {\\rm hadrons}$ experiments even without evolving them in $Q$ by means of DGLAP evolution equation.  \\para As mentioned above, the X particle decay process may involve particles and interactions belonging to possible new physics beyond SM. Most of the recent studies using DGLAP evolution equation method have been done in the context of a particular model of the possible new physics beyond SM, namely, the Minimal Supersymmetric Standard Model (MSSM). While these studies are certainly useful, there exists, however, no direct evidence yet of Supersymmetry in general and the MSSM in particular. Indeed, the unknown nature of the physics beyond SM introduces additional uncertainties in the whole problem over and above the intrinsic uncertainties associated with the DGLAP evolution method itself which is fundamentally based on standard QCD. In order to analyze these uncertainties associated with the DGLAP evolution method itself, we restrict our analysis here to the standard DGLAP evolution equations for FFs based on QCD. Also, to keep our analysis simple, we shall illustrate our main results by considering the behavior of the FF for only one of the hadron species, namely, pions; our general conclusion, however, apply to nucleons as well as to other mesons like the K meson, too.  \\para The rest of this paper is organized as follows: In the following section we set our notations and express the energy spectrum of hadrons resulting from the decay of the X particle in terms of the singlet fragmentation function (FF). In section 3, we review the various methods of evaluating the FF. Our main results are presented and discussed in section 4, and brief conclusions are presented in section 5. ",
        "conclusions": "In this paper we have analyzed the uncertainties involved in obtaining the injection spectra of UHECR particles in the top-down scenario of their origin. We have demonstrated that evaluating the relevant FFs at the values of $M_X$ and $x$ of interest by evolving them (in $Q=M_X$) from their initial (parametrised) values at low $Q$ by numerically solving the DGLAP evolution equation for FF is subject to considerable uncertainties. Indeed, we find that for $x\\lsim 0.1$ (the $x$ region of interest for most large values of $M_X$ of interest), the FF obtained from DGLAP evolution cannot be said to be any more reliable than that provided by the simple Gaussian form (in the variable $\\xi$) based on coherent branching approach to parton shower development. At the same time, we also find that for $x\\gsim0.1$, the evolution of the singlet FF, which determines the injection spectrum, is ``minimal'' --- the singlet FF changes by barely a factor of 2 after evolving over $\\sim$ 14 orders of magnitude in $Q\\sim M_X$. We, therefore, argue that as long as the measurement of the EECR spectrum is going to remain uncertain by a factor of 2 or larger (which is likely to be the case in the foreseeable future), it is good enough for most practical purposes to directly use any one of the available {\\em initial} parametrisations of the FFs in the $x$ region $x\\gsim0.1$ based on low energy (say at the Z-pole) data from $e^+e^-\\to {\\rm hadrons}$ experiments, without any need for evolving them to the required EECR $Q^2$ value.  \\para {\\bf Acknowledgments}\\\\ This work was begun at the Seventh Workshop on High Energy Physics Phenomenology (WHEPP-7) held at Harish-Chandra Research Institute, Allahabad, India, January 4--15, 2002. We thank all the organizers and participants of that Workshop for providing a stimulating workshop environment. The work of PB is partially supported by a NSF US-India cooperative research grant. RB would like to thank D.~Indumathi for useful discussions."
    },
    "0403/astro-ph0403286_arXiv.txt": {
        "abstract": "s{Recent developments in solar, reactor, and accelerator neutrino physics are reviewed. Implications for neutrino physics, solar physics, nuclear two-body physics, and r-process nucleosynthesis are briefly discussed.} ",
        "introduction": "Solar neutrino experiments, especially with the announcement of recent results from the Sudbury Neutrino Observatory (SNO) \\cite{Ahmed:2003kj}, have reached the precision stage. An analysis of the data from SNO as well as data from other solar neutrino experiments (Super-Kamiokande [SK] \\cite{Fukuda:2002pe}, Chlorine \\cite{Cleveland:nv}, and Gallium \\cite{Abdurashitov:2002nt,Hampel:1998xg,Altmann:2000ft}), combined with the data from the reactor experiment KAMLAND \\cite{Eguchi:2002dm}, place severe constraints on the neutrino parameters, especially mixing between first and second generations \\cite{Balantekin:2003dc,deHolanda:2003nj,Balantekin:2003jm}. The neutrino parameter space obtained from such a  global analysis, including the neutral-current results from the SNO salt phase, is shown in Fig. \\ref{fig:1} \\cite{Balantekin:2003jm}.  \\begin{figure} \\includegraphics[scale=0.35]{salt-kl} \\vspace*{0cm} \\caption{ \\label{fig:1} Allowed confidence levels from the joint analysis of all available solar neutrino data (chlorine, average gallium, SNO and SK spectra and SNO salt phase) and KamLAND reactor data The isolines are the ratio of the shifted $^8$B flux to the SSM value. At best fit (marked by a cross) the value of this ratio is determined to be $1.02$ (from Reference 10).} \\end{figure}  The mixing angle between first and second generations of the neutrinos dominates the solar neutrino oscillations whereas the mixing angle between second and third generations dominates the oscillations of atmospheric neutrinos. There are several puzzles in the data. Both mixing angles seem to be close to maximum, very unlike the mixing between quarks. Also the third mixing angle, between first and third generations, seems to be very small, even possibly zero. It is especially important to find out if this mixing angle is indeed different from zero since in the mixing matrix it multiplies a CP-violating phase. Such a CP-violation may have far reaching consequences. To explain the baryon excess (over antibaryons) in the Universe, Sakharov pointed out that it may be sufficient to satisfy three conditions: i) Baryon number non-conservation (which is readily satisfied by the grand unified theories), ii) CP-violation, and iii) Non-equilibrium conditions. It is entirely possible that the CP-violation necessary for the baryogenesis is hidden in the neutrino sector.  \\begin{figure}[t] \\includegraphics[scale=0.29]{all-2} \\vspace*{0cm} \\caption{ \\label{fig:2}  Allowed regions of the neutrino parameter space with solar-density fluctuations when the data from the solar neutrino and KamLAND experiments are used. The SSM density profile of Reference 14 and the correlation length of 10 km are used. The case with no fluctuations ($\\beta=0$) are compared with results obtained with the indicated fractional fluctuation.  The shaded area is the 70 \\% confidence level region. 90 \\% (solid line), 95 \\% (dashed line), and 99 \\% (dotted line) confidence levels are also shown (From Reference 15).} \\end{figure}  It is worth pointing out that high-precision solar-neutrino data have potential beyond exploring neutrino parameter space. Here we discuss two such applications to solar physics and to nuclear physics. ",
        "conclusions": ""
    },
    "0403/astro-ph0403553_arXiv.txt": {
        "abstract": "We report the discovery of a new binary pulsar,  PSR~J1829+2456, found during a mid-latitude drift-scan survey with the Arecibo telescope. Our initial timing observations show the 41-ms pulsar to be in a 28-hr, slightly eccentric, binary orbit. The advance of periastron $\\dot{\\omega}=0.28\\pm0.01$~deg yr$^{-1}$ is derived from our timing observations spanning 200 days. Assuming that the advance of periastron is purely relativistic and a reasonable range of neutron star masses for PSR~J1829+2456 we constrain the companion mass to be between 1.22~M$_\\odot$ and 1.38~M$_\\odot$, making it likely to be another neutron star. We also place a firm upper limit on the pulsar mass of 1.38~M$_\\odot$.  The expected coalescence time due to gravitational-wave emission is long ($\\sim$~60~Gyr) and this system will not significantly impact upon calculations of merger rates that are relevant to upcoming instruments such as LIGO. ",
        "introduction": "The first binary pulsar B1913+16 was discovered by Hulse \\& Taylor (1975). This double neutron star (DNS) system with its 7.75-hr orbital period and large eccentricity ($e=0.6$) has since become a wonderful laboratory for testing general relativity in the strong-field regime \\nocite{tw89} (Taylor \\& Weisberg 1989). Perhaps most importantly, it has provided the first evidence for the existence of gravitational radiation \\citep{tw82}. DNS binaries start life as binary systems with main-sequence stars of mass~$>$~6~M$_{\\odot}$. Eventually the more massive of the two stars undergoes a supernova explosion, leaving a neutron star, sometimes seen as a pulsar. As the less-massive star evolves it increases in size until it overfills its Roche lobe. At this point, matter starts to accrete onto the pulsar causing it to spin up to periods as short as a few milliseconds \\nocite{acrs82} (Alpar et al.~1982), a process known as recycling. In most cases the outer layers of the companion are blown away after accretion, exposing the core of the companion and producing a white dwarf-millisecond pulsar binary. For some recycled systems, however, the companion is massive enough to explode as a supernova. In most cases, this violent event will disrupt the binary system. The DNS binaries are those systems fortunate enough to survive. For further details see \\cite{bv91}.  Our understanding of the DNS binary population is currently hampered by small-number statistics. Large-scale pulsar surveys are being carried out by a number of groups in order to improve this situation by increasing the sample of objects.  In this {\\it Letter} we report on the results of a drift-scan survey using the Arecibo telescope. The survey was conducted at 430 MHz and covered mostly intermediate Galactic latitudes ($|b|<60^{\\circ}$) to optimise the likelihood of finding millisecond and binary pulsars (see \\nocite{cc97} e.g.~Cordes \\& Chernoff 1997).  The known pulsars detected are described along with preliminary parameters for a new 41-ms pulsar J1829+2456 which is likely to be a DNS binary. We compare the properties of this new system with the known DNS and other relativistic binaries. Finally, we outline the future observational prospects for this system. ",
        "conclusions": "\\subsection{Mass determination} The orbital parameters in Table~\\ref{binpars} can be used to constrain the mass of the binary system. The Keplerian mass function relates the orbital period, $P_{b}$, and the projected semi-major axis, $x$, to the masses of the binary components. For this system we calculate a mass function \\begin{equation} f(m_1,m_2) = \\frac{4\\pi^{2} x^3}{P_b^2 T_{\\odot}} = \\frac{(m_{2}\\sin i)^{3}}{(m_{1} + m_{2})^{2}} = 0.294\\ \\rm{M}_{\\odot}, \\label{cmass} \\end{equation} where $T_{\\odot} = GM_{\\odot}c^{-3} = 4.925490947 \\mu$s, $i$ is the inclination between the plane of the orbit and the line of sight, $x$ is in light-seconds, $P_b$ is in seconds and the pulsar and companion masses $m_1$ and $m_2$ are in Solar masses. Neutron stars are observed to have a narrow range of masses. Given the data in \\cite{tc99}, and the low measured mass of PSR~J0737$-$3039B (Lyne et al. 2004), we expect the pulsar mass to lie between 1.25~M$_\\odot$ and 1.47~M$_\\odot$. For an orbit viewed edge-on ($i=90^{\\circ}$), and a minimum neutron star mass for PSR~J1829+2456 of $m_{1}=1.25\\ \\rm{M}_\\odot$, we calculate a minimum companion mass $m_2 \\simeq$ 1.22 M$_\\odot$. As shown in the following equation, the measurement of the advance of periastron, $\\dot{\\omega}$, if assumed to be purely relativistic, allows a measurement of the sum of the masses: \\begin{equation} \\dot\\omega=3\\left(\\frac{2\\pi}{P_{b}}\\right)^{5/3}T_{\\odot}^{2/3}(m_{1}+m_{2})^{2/3}(1-e^{2})^{-1}. \\label{omegadot} \\end{equation} When the measured $\\dot \\omega$ is used in combination with the constraint that $\\sin i<1$, we obtain an upper limit on the maximum pulsar mass~$m_1<1.38$ M$_{\\odot}$. Similarly, assuming $m_1>1.25$ M$_{\\odot}$, the maximum value of $\\dot{\\omega}$ implies $i>66^{\\circ}$. These constraints are shown in a mass-mass diagram in Fig.~\\ref{massdiag}. Although we cannot currently rule out a massive white dwarf or main-sequence star, given the similar mass functions, spin and orbital parameters of J1829+2456 to other DNS binaries (see Table~\\ref{DNSsystems}) it seems most likely that the companion is another neutron star.  \\begin{figure} \\includegraphics[width=8cm, angle=0]{fig3.ps} \\caption{The constraints on the masses of the pulsar and companion. The vertical lines are the mass limits for neutron stars described in Thorsett \\& Chakrabarty (1999), and the low measured mass of PSR~J0737$-$3039B (Lyne et al. 2004). The lower companion mass limit is given by the mass function assuming a maximum inclination of $90^{\\circ}$. The advance of periastron, $\\dot{\\omega}$, if assumed to be purely relativistic, provides the constraint for the maximum companion mass and maximum pulsar mass.} \\label{massdiag} \\end{figure}  \\subsection{A search for the companion} The accuracy of the position of PSR~J1829+2456 inferred through radio timing is sufficient to allow for a search for an optical companion. No optical companion is present in the uncalibrated plates of the Digitised Sky Survey. Preparations are underway for a deep optical search for the companion star.  Given the possibility that the companion of J1829+2456 is a neutron star, a search was made of the available radio data for a periodic signal which could be coming from the companion if it were active as a radio pulsar. This search was given greater significance following the recent detection of 2.8-s pulsations from the companion of PSR~J0737$-$3039A (Lyne et \\nocite{lbk+04} al.~2004). Initial searches of the data (at the DM of PSR~J1829+2456) looking for other periodic signals showed that, if the companion was an active radio pulsar, it was too weak to be easily detectable.  To carry out a more sensitive search for the companion, we made use of the fact that the orbital parameters of the system are known and that the pulsations from the companion star would be Doppler shifted in the opposite sense to PSR~J1829+2456 due to its orbital motion about their common centre of mass. To correct for this effect, we applied the first-order Doppler formula to calculate the effective sampling interval in the rest-frame of the companion: \\begin{equation} t_{\\rm companion} = t_{\\rm samp} (1 + v(t)/c), \\end{equation} where $t_{\\rm samp}$ is the sampling interval at the observatory and $v(t)$ is the effective line-of-sight radial velocity of the companion as a function of time, $t$. We used TEMPO to calculate $v(t)$ to account for the effects of orbital motion in the binary system (using the orbital parameters given in Table~\\ref{binpars}, assuming a pulsar mass of 1.35~M$_\\odot$ and changing only the longitude of periastron by $180^{\\circ}$) as well as contributions from the Earth's rotation and motion about the Sun. All available data were dedispersed and resampled in this way before being passed through the standard search analysis. No significant candidates were found down to a S/N of 5. This corresponds to a 430-MHz flux limit of $\\sim$~0.1~mJy.  While the lack of detection of the companion suggests that it is too weak to be seen as a radio pulsar, or unfavourably beamed, it should be noted that the companion to PSR~J0737$-$3039A is only clearly detectable over certain phases of the orbit (Lyne et al.~2004). If any companion to PSR~J1829+2456 was behaving in a similar fashion it may not be visible in the current data set. Further, more sensitive, searches will be carried out as more data are taken.  \\subsection{Expected relativistic parameters} Using the constrained range of masses for the pulsar and companion the expected orbital period derivative and period for geodetic precession due to misalignment of the spin axis of the pulsar with the orbital angular momentum vector \\citep{bo75b} can be determined. Assuming a combined mass of 2.53~M$_\\odot$, the expected values for $\\dot P_{b}$ and the expected timescale for geodetic precession are given in Table~\\ref{binpars}. The expected $\\gamma$ should be measurable after 18 months of timing data. The expected $\\dot P_{b}$ is too small to be measurable without a long-term (5--10 yr) timing campaign. The expected timescale for geodetic precession is $\\sim$~4200~yrs, too long to be measurable through pulse profile or timing variations.  \\begin{table} \\caption{Measured, derived, and expected parameters for PSR~J1829+2456 from 29 TOAs from MJDs 52786 to 52986.} \\begin{center} \\begin{tabular}{l l} \\hline \\multicolumn{2}{c}{Measured Parameters}                                         \\\\ \\hline Right ascension (J2000) (h:m:s)                                & 18:29:34.6 (1) \\\\ Declination (J2000) ($^{\\circ}$:':'')                          & 24:56:19 (2)   \\\\ Period (ms)                                                    & 41.00982358 (1)\\\\ Epoch of period (MJD)                                          & 52887          \\\\ Projected semi-major axis (s)                                  & 7.2360 (1)     \\\\ Dispersion measure (pc cm$^{-3}$)                              & 13.9 (5)       \\\\ Binary period (days)                                           & 1.176028 (1)   \\\\ Eccentricity                                                   & 0.13914 (4)    \\\\ Longitude of periastron (deg)                                  & 229.94 (1)     \\\\ Epoch of periastron (MJD)                                      & 52848.57977 (3)\\\\ 430-MHz flux density (mJy)                                     & 0.3 (1)        \\\\ Pulse width at 50\\% of peak $w_{50}$ (ms)                      & 1.07 (16)      \\\\ Pulse width at 10\\% of peak $w_{10}$ (ms)                      & 2.07 (16)      \\\\ Rate of advance of periastron (deg yr$^{-1}$)                  & 0.28 (1)       \\\\ Mass function (M$_\\odot$)                                      & 0.29413 (1)    \\\\ RMS residual to fit ($\\mu$s)                                   & 19             \\\\ \\hline \\multicolumn{2}{c}{Derived Parameters}                                          \\\\ \\hline Galactic longitude (J2000) (deg)                               & 53.343 (1)     \\\\ Galactic latitude (J2000) (deg)                                & 15.612 (1)     \\\\ Distance (kpc)$^{a}$                                           & $\\sim$ 1.2     \\\\ Mean orbital speed (km s$^{-1}$)                               & 134            \\\\ Total system mass (M$_\\odot$)                                  & 2.5 (2)        \\\\ Minimum companion mass (M$_\\odot$)                             & 1.22           \\\\ Maximum pulsar mass (M$_\\odot$)                                & 1.38           \\\\ \\hline \\multicolumn{2}{c}{Expected Relativistic Parameters}                            \\\\ \\hline Orbital period derivative ($\\times10^{-12}$)                   & $-$0.02        \\\\ Coalescence time (Gyr)                                         & 60             \\\\ Geodetic precession rate (deg yr$^{-1}$)                       & 0.075          \\\\ Relativistic time-dilation\\\\ and gravitational redshift $\\gamma$ (ms)   & 1.3   \\\\ \\hline \\end{tabular} \\end{center} a: Distance inferred from DM \\citep{cl02}.\\\\ The numbers in parentheses are the 1--$\\sigma$ uncertainties in the least significant digit quoted. \\label{binpars} \\end{table}  \\subsection{Comparison with other relativistic binary systems} If our future measurements show PSR~J1829+2456 to be a double neutron star system it will be the seventh such system to be discovered. Table~\\ref{DNSsystems} contains a list of the most relativistic binary pulsar systems and their orbital parameters. We also list the coalescence time $\\tau_{\\rm GW}$ due to the emission of gravitational radiation of each system. Following Lorimer (2001)\\footnote{We reproduce the formula here to correct a typographic error in the original citation.}, we can approximate the detailed calculations of $\\tau_{\\rm GW}$ by Peters \\nocite{pet64} (1964) via \\begin{equation} \\label{coaltime} \\tau_{\\rm GW}     \\simeq  10^7 \\, {\\rm yr} \\, \\left(\\frac{P_{b}}{\\rm hr}\\right)^{8/3} \\left(\\frac{\\mu}{{\\rm M}_{\\odot}}\\right)^{-1} \\left(\\frac{m_1+m_2}{{\\rm M}_{\\odot}}\\right)^{-2/3} (1-e^{2})^{7/2}, \\end{equation} where the reduced mass $\\mu = m_1 m_2 / (m_1+m_2)$. For any reasonable range of $m_1$ and $m_2$ discussed above, we find $\\tau_{\\rm GW} \\sim$~60~Gyr.  The systems listed in Table~\\ref{DNSsystems} span a wide range of orbital parameters but can be broadly split into systems that will and will not coalesce within a Hubble time, with PSR~J1829+2456 lying toward the edge of the non-coalescing group. The relatively long coalescence time for J1829+2456 means that it will not affect the DNS merger rate calculations (e.g. Kalogera et al. 2004) for gravitational wave detectors such as LIGO.  \\begin{table*} \\caption{The orbital parameters of various eccentric binary systems.} \\begin{center} \\begin{tabular}{l r@{.}l r@{.}l r@{.}l r@{.}l r@{.}l r@{.}l c r@{.}l c c} \\hline PSR &\\multicolumn{2}{c}{$P$} &\\multicolumn{2}{c}{$P_{b}$} &\\multicolumn{2}{c}{$a_{1}\\sin~i$} &\\multicolumn{2}{c}{$e$} &\\multicolumn{2}{c}{$\\dot \\omega$} &\\multicolumn{2}{c}{$\\dot P _{b}$} & $f(m)$ &\\multicolumn{2}{c}{$m_{1} + m_{2}$} &$\\tau_{\\rm GW}^{\\dag}$ & References\\\\ &\\multicolumn{2}{c}{(ms)}  &\\multicolumn{2}{c}{(days)} &\\multicolumn{2}{c}{(lt-s)} &\\multicolumn{2}{c}{} &\\multicolumn{2}{c}{(deg yr$^{-1}$)} &\\multicolumn{2}{c}{($\\times 10^{-12}$)} & \\multicolumn{2}{c}{(M$_\\odot$)} &(M$_\\odot$) &(Gyr) & \\\\ \\hline &\\multicolumn{17}{c}{Double neutron star binaries}\\\\ \\hline B1913+16     &  59&03 &  0&323 &  2&34 & 0&617 &  4&227                        & $-$2&428                 & 0.13 & 2&83                   & 0.31  & 1  \\\\ B1534+12     &  37&90 &  0&421 &  3&73 & 0&274 &  1&756                        & $-$0&138                 & 0.31 & 2&75                   & 2.69  & 2  \\\\ B2127+11C    &  30&53 &  0&335 &  2&52 & 0&681 &  4&457                        & $-$3&937                 & 0.15 & 2&71                   & 0.22  & 3  \\\\ J1518+4904   &  40&93 &  8&634 & 20&04 & 0&249 &  0&011                        &\\multicolumn{2}{c}{--}    & 0.12 & 2&62                   & 9600  & 4  \\\\ J1811$-$1736 & 104&18 & 18&779 & 34&78 & 0&828 &  0&009                        &\\multicolumn{2}{c}{$<$30} & 0.13 & 2&6                    & 1700  & 5  \\\\ J0737$-$3039A&  22&70 &  0&102 &  1&42 & 0&088 & 16&88                         & $-$1&24$^{\\ast}$         & 0.29 & 2&58                   & 0.087 & 6  \\\\ \\hline &\\multicolumn{17}{c}{White dwarf binaries}\\\\ \\hline B2303+46     &1066&37 & 12&34  & 32&69 & 0&66  &  0&010                        &\\multicolumn{2}{c}{--}    & 0.25 & 2&53                   & 4500  & 7  \\\\ J1141$-$6545 & 393&90 &  0&20  &  1&86 & 0&17  &  5&33                         &\\multicolumn{2}{c}{$<$50} & 0.18 & 2&30                   & 0.59  & 8  \\\\ \\hline &\\multicolumn{17}{c}{Unknown companion}\\\\ \\hline B1820$-$11   & 279&83 &357&76  &200&67 & 0&79  &\\multicolumn{2}{c}{$<10^{-4}$} &\\multicolumn{2}{c}{--}    & 0.07 & \\multicolumn{2}{l}{--} & --    & 9  \\\\ J1829+2456   &  41&00 &  1&17  &  7&24 & 0&14  &  0&28                         & $-$0&02$^{\\ast}$         & 0.29 & 2&53                   & 60    & -- \\\\ \\hline \\end{tabular} \\begin{tabular}{l l} 1: \\cite{ht75a, wt03a, tw89} &6: \\cite{bdp+03}                                \\\\ 2: \\cite{wol90z, sttw02}     &7: \\cite{tamt93, arz95}                         \\\\ 3: \\cite{agk+90, and92}      &8: \\cite{klm+00a}                               \\\\ \\cite{pakw91, dk96}          &9: \\cite{lm89, tamt93}                          \\\\ 4: \\cite{nst96, hlk+03}      &$\\ast$: Predicted value.                        \\\\ 5: \\cite{lcm+00}             &$\\dag$: Calculated using formula \\ref{coaltime}.\\\\ \\end{tabular} \\end{center} \\label{DNSsystems} \\end{table*}"
    },
    "0403/astro-ph0403079_arXiv.txt": {
        "abstract": "I review cosmological simulations of X-ray clusters. Simulations have increased in resolution dramatically and the effects of radiative cooling, star formation feedback, and chemical enrichment on the ICM are being simulated. The structure and evolution of non-radiatve X-ray clusters is now well characterized. Such models fail to reproduce the observed $L_x-T$ relation, implying the need for additional physics. Simulations adding radiative cooling produce too much cool gas and unreasonably high X-ray luminosities. Simulations including star formation and feedback appear more promising, but need further refinement. New observations should help in this regard. ",
        "introduction": "As the largest gravitationally bound objects in the universe, clusters of galaxies have attracted the attention of observers and numerical simulators alike. For over a decade, beginning the with pioneering hydrodynamic simulations of Evrard (1990), numerical simulations have been used to understand the physics of X-ray cluster formation and to predict their abundance at high redshift which is a sensitive probe of cosmology (see review Henry in these proceedings.) Observationally, clusters of galaxies have historically been studied in the optical and X-ray portions of the EM spectrum (Forman \\& Jones 1982). X-rays in particular provide an unambiguous method for detecting clusters at low and intermediate redshift, and many surveys have been conducted (Henry, these proceedings.)  A number of groups have simulated the formation of statistical ensembles of X-ray clusters (Kang et al. 1994; Bryan et al. 1994a,b; Bryan \\& Norman 1998; Eke, Navarro \\& Frenk 1998; Yoshikawa, Jing \\& Sato 2000) in order to provide a theoretical bridge between what is observed--the X-ray luminosity function (XLF) and X-ray temperature function (XTF)--and the cluster mass function (CMF). The CMF in turn is directly related to the matter fluctuation power spectrum $P(k)$--one of the holy grails of observational cosmology. In so doing, simulators have discovered that X-ray clusters are not the simple gas-bags they were once thought to be. It has been found that quite high resolution is required to converge on the predicted properties of non-radiative clusters (Anninos \\& Norman 1996; Frenk et al. 1999), and that the inclusion of radiative cooling and other non-adiabatic effects strongly affects the clusters' emission and structural properties (e.g., Pearce et al. 2000).  This review will follow the development of cosmological simulations of X-ray clusters primarily from a historical perspective, starting with the non-radiative simulations of Evrard (1990) and others and concluding with current models incorporating cooling, star formation, supernova feedback and chemical enrichment. The field has been enlivened by the arrival of new observations and new questions. I will attempt to keep the questions at the forefront of this review, for while some have been convincingly answered, many are still open. Simulations of cluster mergers done outside the framework of CDM-driven structure formation are not reviewed here for space reasons. ",
        "conclusions": ""
    },
    "0403/astro-ph0403623_arXiv.txt": {
        "abstract": "s{The Combined EIS-NVSS Survey Of Radio Sources (CENSORS) has been produced with the primary goal of investigating the cosmological evolution of the radio luminosity function. This 1.4GHz sample, complete to the $7.2$mJy level, contains 150 radio sources. Host galaxies are almost entirely identified in optical and near-IR bands and the sample is now approaching 70\\% spectroscopic completeness. We show preliminary results demonstrating how CENSORS will improve upon previous work and how it is applicable to other projects.} ",
        "introduction": "Radio galaxies and radio loud quasars are among the most powerful objects in the Universe. Observable at high redshifts, they trace large scale structure and are associated with the most massive black holes[\\refcite{rad_bhmass}], they therefore have great potential for revealing the evolution of massive galaxies and investigating the relation between active galactic nuclei (AGN) and galaxy formation. However there are still many questions facing radio-loud AGN astrophysicists. For example, the high redshift cosmological evolution of strongly radio active sources is a property that has not been resolved and so provides a route by which we may gain a better understanding of these objects. This is of far-reaching importance: given their association with the most massive black holes, we would also learn about the evolution of the upper part of the black hole mass distribution.  In 1990 Dunlop and Peacock[\\refcite{dp90}] undertook a  study of cosmic evolution of the radio luminosity function (RLF) presenting the first evidence of a decline in the comoving number density of powerful radio sources beyond z $\\squig$ ~2.5 (the redshift cut-off). Since then numerous advances have provided a good consensus in the determination of the low redshift RLF (eg.[\\refcite{mob1999}]), but the high redshift evolution and the reality of the redshift cut-off in the radio source population remain areas of controversy. For example, the deep sample of Waddington et al.[\\refcite{wad2001}] shows evidence of a cut-off in number density of low luminosity sources beyond z $>$ 2, but has insufficient sky coverage to investigate the most luminous sources. Conversely, the sample of Jarvis et al.[\\refcite{jarvis2001}] proved too shallow to resolve the issue. Shaver et al.[\\refcite{shaver}] claimed evidence for a sharp decline in number density of flat spectrum sources between $z \\sim 2.5$ and $z \\sim 5$. However Jarvis and Rawlings[\\refcite{jar_raw}] showed that this result was erroneous due to a failure to properly account for the  spectral index of the sources. So high redshift RLF evolution of radio sources remains ill-understood. ",
        "conclusions": ""
    },
    "0403/astro-ph0403309_arXiv.txt": {
        "abstract": "{ Four Bok globules were studied in the Near-Infrared, through narrow-band filters, centered at the 1.644~$\\mu$m line of [FeII], the H$_2$-line at 2.122~$\\mu$m, and the adjacent continuum. \\hfill\\break\\noindent We report the discovery of \\FE\\ and H$_2$ protostellar jets and knots in the globules CB3 and CB230. The [FeII]-jet in CB230 is defined by a continuous elongated emission feature, superimposed on which two knots are seen; the brighter one lies at the tip of the jet. The jet is oriented in the same direction as the large-scale CO outflow, and emerges from the nebulosity in which a Young Stellar Object is embedded. The H$_2$ emission associated with this jet is fainter and wider than the [FeII] emission, and is likely coming from the walls of the jet-channel. \\hfill\\break\\noindent In CB3 four H$_2$ emission knots are found, all towards the blue-shifted lobe of the large-scale outflow. There is a good correspondence between the location of the knots and the blue-shifted SiO(5$-$4) emission, confirming that SiO emission is tracing the jet-like flow rather well. \\noindent No line emission is found in the other two targets, CB188 and CB205, although in CB205 faint line emission may have been hidden in the diffuse nebulosity near the IRAS position. Around this position a small group of ($\\geq 10$) stars is found, embedded in the nebula. A diffuse jet-like feature near this group, previously reported in the literature, has been resolved into individual stars. ",
        "introduction": "\\label{int} Outflow and infall are inextricably associated with the very earliest stages of star formation. Even while still accreting matter, a newborn star generates a fast, well-collimated stellar wind that forms jets which sweep up the ambient molecular gas, creating bipolar molecular outflows. The high-velocity winds create shocks, which heat the gas to thousands of K when breaking up into the ambient gas. Understanding the details of the mechanism producing the acceleration of the outflow is fundamental to the understanding of the star formation process itself.  \\noindent Emission of millimetre molecular lines (usually CO) at excitation temperatures T$_{\\rm ex} \\approx 10-20$~K, is used to study the outflow's large-scale morphology. Whereas this component consists mostly of swept-up cloud material, and as such offers a time-averaged picture, the {\\it present} flow activity is represented by a fast, hot component, traced by H$_2$ ro-vibrational lines at $T_{\\rm ex} \\sim 2000$~K. Where this jet-like component interacts with the slower flowing gas or the ambient cloud, bow-shaped shock fronts are visible. This component, at T$_{\\rm ex} \\sim 100$~K, can be traced by mm-emission of molecular species that are produced only in a shock-driven chemistry (e.g Bachiller~\\cite{bach}).  The link between the 10-100~K gas and the hot jet-component is not well understood. In particular, it is still an open question how the efficiency of the processes leading to chemical anomalies depends on the shock-type (J/C) and -characteristics. The comparison between the shocked gas components at different temperatures is thus fundamental to the study of the jet/outflow system and the energetics of the star forming process.  \\begin{figure} \\resizebox{7.5cm}{!}{\\rotatebox{270}{\\includegraphics{0931fig1.eps}}} \\caption[]{TNG/NICS observations of CB3: The H$_2$ emission (still including the contribution from the continuum). The 4 knots (indicated by arrows) of H$_2$ line-emission are not at all visible in the K-continuum image (not shown). } \\label{cb3h2kcont} \\end{figure}  \\begin{figure} \\resizebox{8.5cm}{!}{\\rotatebox{270}{ \\includegraphics{0931fig2.eps}}} \\caption[]{{\\bf a}\\ A zoom-in on the four H$_2$ knots in CB3. To better outline k2 and k3, contours for them are drawn with smaller increments than for the other knots in the panel; lowest contour is at $\\sim 2.5\\sigma$. The dashed lines indicate the alignment of pairs of knots with the outflow centre. The centre of the outflow (Codella \\& Bachiller~\\cite{codbach}) lies at offset (0,0) and is indicated by a cross. The 1.3-mm peak (Launhardt \\& Henning~\\cite{launhen}) is indicated by the star. {\\bf b}\\ As a, but overlaid with the contours of integrated SiO(5--4) emission at $-42.5 (\\pm 0.5)$~km s$^{-1}$ (Codella \\& Bachiller~\\cite{codbach}). The angular resolution of the SiO map is 11$\\arcsec$. } \\label{cb3h2} \\end{figure}  \\smallskip\\noindent A very good place to study the jet-component, the large-scale outflow, and their interaction are the Bok globules: cold (10 K) and relatively isolated molecular clouds associated with star formation. A catalogue of such objects (at $\\delta > -30^{\\circ}$) was compiled by Clemens \\& Barvainis (\\cite{clemens}; hereafter CB). Because of relatively simple structure, globules form mainly low-mass stars in small numbers, and are therefore without the observational confusion that one encounters in regions like Orion or Ophiuchus.  With the italian TNG (Telescopio Nazionale Galileo) we have searched in four globules for the jet-component in the Near-InfraRed (NIR) through narrow-band H$_2$ (2.122~$\\mu$m) and [FeII] (1.644~$\\mu$m) filters. These two lines are particularly useful, as [FeII] traces (J-) shocks with velocities of a few 100~km s$^{-1}$ and is therefore expected to outline the inner jet-channel, closest to the driving source of the flow. Molecular hydrogen, which dissociates at shock velocities $>25-45$~km\\,s$^{-1}$, traces slower (C-) shocks and is excited in bow shocks and in shock-wakes (e.g. Allen \\& Burton~\\cite{allen}) and is therefore a good probe for the region of interaction between jet and ambient material. All four globules have been observed in the broad-band NIR (Yun \\& Clemens~\\cite{yc94a}) and detected at 1.3~mm continuum (Launhardt \\& Henning~\\cite{launhen}) and were found to contain embedded objects. Molecular outflows were also found in all four objects (Yun \\& Clemens~\\cite{yc94b}; Codella \\& Bachiller~\\cite{codbach}). Together, these findings identify these globules as good candidates in a search for protostellar jets. In this paper we report the detection of H2 and [FeII] line emission knots and jets in two of the four globules. ",
        "conclusions": "\\label{res}  \\subsection{Removing the continuum}  Observations in the H$_2$ and [FeII] filters contain, apart from the emission in the lines themselves (if present) also a contribution from the K- (in the case of H$_2$) and H-continuum (in the case of [FeII]). Therefore in order to detect the line emission the continuum has to be subtracted. Because the central wavelength of the continuum filters is different from those of the line filters, a proper subtraction requires a careful scaling of the emission detected in the continuum filters. Field stars are used for this, because their emission is expected to be continuum only.  \\noindent We have assumed that the wavelength-dependence of the emission in the NIR is proportional to $\\lambda^x$. The value of the slope $x$ is derived from a comparison of the flux of a number of field stars in the line- and continuum filters.  \\noindent In CB188 and CB205, no jet-like features are detected after subtraction, whereas in CB3 and CB230 pure H$_2$ and [FeII] emission knots and protostellar jets have been clearly revealed. \\hfill\\break\\noindent Photometry of the line-emission knots has been performed on the subtracted images using the task {\\tt POLYPHOT} in IRAF, roughly including the emission down to a $\\sim 1 \\sigma$ limit. The sensitivity limit (of integrated line flux) is $2-3 \\times 10^{-16}$ erg cm$^{-2}$ s$^{-1}$ arcsec$^{-1}$, both in the H$_{2}$ and in the [FeII] images; the limiting flux increases towards regions of diffuse emission.  \\begin{figure*} \\resizebox{10cm}{!}{\\rotatebox{270}{\\includegraphics{0931fg4a.eps}}} \\resizebox{6cm}{!}{\\rotatebox{270}{\\includegraphics{0931fg4b.eps}}} \\caption[]{ Results of the TNG/NICS observations of CB230 of the H$_2$ and [FeII]-emission (still including continuum emission): (Central panel) A zoom-in on the objects detected in H$_2$ in the core of CB230, with contours to better outline the more intense emission. The YSO is indicated by ``A''. (Left) A close-up of the smaller emission object, which is clearly seen to contain two sources (B1 and B2). The right-hand panel shows the contoured [FeII] emission (still including the continuum). A clear difference is seen in the diffuse emission just North of the YSO, with respect to the H$_2$ image (central panel); this is the signature of a jet (see Fig.~\\ref{feiih2closeup}). } \\label{cb230h2} \\end{figure*}  \\subsection{CB3}  This globule is associated with different generations of star formation, as pointed out by the presence of a NIR Young Stellar Object (YSO; Yun \\& Clemens~\\cite{yc95}, \\cite{yc94a}), a 1.3 mm object (Launhardt \\& Henning \\cite{launhen}), and a sub-mm source (Huard et al. \\cite{huard}). This globule also contains an IRAS point source (IRAS00259+5625), which lies among these objects; its flux is probably derived (in part) from contributions by these various objects. A molecular outflow was detected by Yun \\& Clemens (\\cite{yc92}, \\cite{yc94b}) in CO and mapped in various molecular lines by Codella \\& Bachiller (\\cite{codbach}), who found that the outflow is centered at or very near the 1.3~mm source, which is 17$\\farcs$4 W, 1$\\arcsec$ S of the IRAS position.  \\noindent Our NIR images are centered on the centre position of the outflow. Fig.~\\ref{cb3h2kcont} shows the image obtained in the H$_2$ filter. Even though this image still contains a contribution from the continuum, four regions of H$_2$ emission, hereafter called knots k1, k2, k3, and k4 (North to South), are clearly distinguishable; they are not present in the K$_{\\rm cont}$-image (not shown). Fig.~\\ref{cb3h2}a shows a close-up of the central region of the narrow-band line image; with the dashed lines we have tried to suggest how the spots can be aligned in pairs and how they may be traced back to the centre of the outflow (which, within the accuracy of the positions is presumably coincident with the 1.3~mm source). The lines show that the centers of k1 and k3 are aligned with the center of the outflow, while the line connecting the origin of the flow and the centers of k2 and k4 is slightly offset from that. This may suggest a change of direction of the outflow axis with time.  \\noindent Though unbeknownst to us at the time of the observations, the features k1 and k4 in Fig.~\\ref{cb3h2}a were already noted in deep NIR images of this region by Launhardt et al. (\\cite{launISO}), who remarked that these two diffuse, non-stellar objects ``might have H$_2$-line emission contributing to their K-band luminosity''. The fact that these features are not visible in our K$_{\\rm cont}$ image confirms that they arise from pure H$_2$ line emission. The absolute coordinates of the four H$_2$ knots, as well as their integrated flux, are given in Table~\\ref{knotparams}.  \\noindent Knot k4 lies at a distance of $\\sim 40\\arcsec$ from the centre of the outflow; for an assumed distance of 2.5~kpc (Launhardt \\& Henning~\\cite{launhen}) this corresponds to $\\sim 0.5$~pc. Knot k1 is the nearest to the outflow centre, at a distance of 10$\\arcsec \\approx 0.12$~pc. All knots are found South of the outflow centre, towards the blue lobe of the molecular outflow (Codella \\& Bachiller \\cite{codbach}). The fast outflow component is usually traced by H$_2$ line emission, which originates behind the shocks created where this jet-component interacts with the ambient cloud material. That we do not see any H$_2$ spots associated with the red (northern) lobe of the outflow is likely because it is much more embedded in the globule, and the emission is too much obscured to be visible even at the NIR wavelength of H$_2$.  \\noindent Closer inspection of the brightest knots k1 and k4 shows that the emission contours are slightly convex. Together with the knot alignment, and the location of the knots along the outflow axis, this argues in favour of the bow shock model for the interaction of a wind from the YSO and the ambient medium.  \\noindent Downstream from the shocks caused by the interaction of the YSO-wind and the ambient medium, in high-density and -temperature gas, one expects to find an enhancement of several molecular species, which are liberated from the mantles of dust grains and injected into the gas phase through endothermic or gas-grain reactions as well as through sputtering. Codella \\& Bachiller (\\cite{codbach}) have shown that the high-velocity component of the outflow in CB3 is well-traced by SiO, which is enhanced in the presence of shocks and which is only present along the main flow axis. In Fig.~\\ref{cb3h2}b we compare the location of the H$_2$ spots with (blue-shifted) channel maps of the emission of SiO(5--4). There is a close correspondence between the location of (at least two of) the H$_2$ spots and the peaks of the SiO distribution. This confirms that in CB3 the SiO emission is tracing the jet-like flow rather well, and indicates that along the outflow axis different temperature regimes coexist: the warm component at $\\sim$100~K (Codella \\& Bachiller 1999) traced by SiO and a hot ($\\ge$2000~K) component revealed by H$_2$.  \\noindent In Fig.~\\ref{cb3allobjects} we show the locations of the various tracers of star formation found in this globule, projected on the narrow-band H$_2$ image. An evolutionary sequence can be distinguished, in that we encounter progressively younger objects when going from the NIR-detected object YC1 (Yun \\& Clemens~\\cite{yc95}), via the sub-mm peaks (Huard et al. \\cite{huard}), to the 1.3~mm peak (Launhardt \\& Henning~\\cite{launhen}), which is probably coincident with the centre of the outflow (Codella \\& Bachiller~\\cite{codbach}). This implies that star formation in this globule has been going on for some time. Assuming the IRAS source and the sub-mm peaks refer to the same object, the mutual distance between the 3 star formation indicators in Fig.~\\ref{cb3allobjects} is about 0.2~pc.  \\noindent We take this opportunity to point out some inconsistencies in the original papers regarding the relative positions of the various objects plotted in Fig.~\\ref{cb3allobjects}. First, we note that the object identified as YSO (based on its NIR excess) in Yun \\& Clemens (\\cite{yc95}) is different from the one they identified in CB3 in Yun \\& Clemens~\\cite{yc94a}, which is about 100$\\arcsec$ to the North. And even so, although YC1-I is a bright source (K = 8.24~mag.; Yun \\& Clemens~\\cite{yc95}), it does not coincide with any of the sources in our image -- the nearest object (see Fig.~\\ref{cb3allobjects}) is $\\sim 4\\arcsec$ to the South, much more than the estimated positional accuracy of both our objects and YC1-I ($\\sim 2\\arcsec$). \\hfill\\break\\noindent The relative positions of the Huard et al. (\\cite{huard}) sub-mm peak and the IRAS source are incorrect in their Fig.~1: based on the coordinates given by Huard et al. the sub-mm peak should be much closer to the IRAS position ($\\sim 5\\arcsec$ rather than 15$\\arcsec$). Although Launhardt \\& Henning (\\cite{launhen}) do not give a positional accuracy for their 1.3~mm peak, it is probably similar to that of the sub-mm peak (i.e. $\\sim 3\\arcsec$), suggesting that these are different objects.  \\subsection{CB230}  This globule also hosts a single IRAS source (IRAS21169+6804), which is associated with the YSO found by Yun \\& Clemens (\\cite{yc94a}). The YSO is located at the apex of a cone-shaped nebulosity (see Fig.~\\ref{cb230h2}) and is at the centre of a CO outflow (Yun \\& Clemens~\\cite{yc92}; \\cite{yc94b}). The outflow is bipolar, but while the red component is quite compact, the blue lobe is more elongated, and extends northwards beyond the boundaries of the map made by Yun \\& Clemens (\\cite{yc94b}) (i.e. its length $>$3\\arcmin). Yun \\& Clemens (\\cite{yc94a}) found a second NIR object, \"possibly forming a binary system\" together with the YSO at the base of the outflow. In our observations we find this second NIR object to consist of two nuclei, which would make this a triple system (Fig.~\\ref{cb230h2}). The conical nebula seen in Fig.~\\ref{cb230h2} has a steep intensity-gradient towards the South, while towards the North the nebulous light fans out and decreases in intensity more gradually. The nebula opens up to the North, which is also the direction of the blue lobe of the outflow, the axis of which is North-South (Yun \\& Clemens~\\cite{yc94b}). The approaching component of the bipolar outflow may have cleared out a cavity; the more extended emission seen in the northern part of the nebula could then be scattered light from walls of this cavity (see Yun \\& Clemens~\\cite{yc94a}). Likewise, the diffuse emission associated with the embedded binary object $\\sim 9\\arcsec$ W of the main nebula is extended towards the SE (see Fig.~\\ref{cb230h2}). This might be the signature of an as yet undetected (or more likely: unresolved) outflow associated with the objects embedded herein.  \\noindent Because we want to distinguish the (pure) line emission from the nebular emission, and because the spectral slope of the nebulosity differs from that of the field stars, we have determined the slope $x$ of the wavelength dependence of the emission ($\\propto \\lambda^x$) by using the integrated flux of the northern part of the nebula, excluding the region with the embedded star. As we shall see (e.g. Fig.~\\ref{feiih2closeup}) the consequence of this is that in the continuum-subtracted [FeII]-image the YSO is still visible. Its spectral index differs from that of the nebula (which is radiation scattered by the dust), which we have used to scale the H$_{\\rm cont}$-image; the YSO cannot therefore be correctly subtracted.  \\noindent The final reduction, with careful scaling of the H$_{\\rm cont}$- and K$_{\\rm cont}$ images of the diffuse emission, reveals the pure line emission that remains after subtraction of the continuum, and is shown in Fig.~\\ref{feiih2closeup}: a jet is seen in [FeII], primarily defined by two knots along its length: a bright one at the tip (hereafter called k1) and a fainter one (k2) about half-way between k1 and the YSO. The knots are superimposed on a fainter, but clearly visible elongated emission feature. We stress that this detection is independent of potential inaccuracies in the continuum subtraction, as the knots are visible in the contours of the [FeII] image before continuum subtraction (cf. Fig.~\\ref{cb230h2}), but not at all in the (H$_{\\rm cont}$) continuum image alone. \\hfill\\break\\noindent The jet is oriented in the N-S direction, and lies at the base of the large-scale molecular outflow. This jet is also traced by H$_2$ emission (Fig.~\\ref{feiih2closeup}b), which is weaker and slightly less narrow than that of [FeII]. Moreover, the H$_2$ emission appears to be anti-correlated with [FeII] peaking next to, rather than on the secondary knot in [FeII], and we conclude that in H$_2$ we probably see the {\\sl walls} of the jet-channel. The detection of strong [FeII]- and weaker H$_2$ emission suggests the presence of fast, dissociative J-shocks.  \\begin{figure*} \\resizebox{11cm}{!}{\\rotatebox{270}{ \\includegraphics{0931fg5a.eps}}} \\hspace{0.5cm} \\resizebox{6cm}{!}{\\rotatebox{270}{ \\includegraphics{0931fg5b.eps}}} \\caption[]{A close-up of the jets and knots in CB230, after subtraction of the continuum. {\\bf a}\\ [FeII]-emission, and {\\bf b}\\ the same but overlaid with H$_2$ contours (white). {\\bf c}\\ Comparison of cuts through centre of the knots in the [FeII]-jet at constant Dec-offset from the continuum-subtracted image. The profile for knot k1 is drawn as a thick histogram, that for k2 as a thin line. Most of the elevated emission plateau between about $-8\\arcsec\\ {\\rm and} +9\\arcsec$ is an artifact of the mosaicing together of the different frames constituting the final image. } \\label{feiih2closeup} \\end{figure*}  \\noindent  The total length of the jet, from its base at the location of the embedded YSO to the edge of knot k1 is about 9\\pas5, corresponding to $\\sim 0.02$~pc at the assumed distance of 450~pc (Launhardt \\& Henning~\\cite{launhen}). Profiles of the knots in the [FeII]-jet, along cuts made at constant declination offsets and passing through the location of peak intensity of the knots, are shown in Fig.~\\ref{feiih2closeup}c. Both features are resolved in the NIR image. The width (FWHM) of k1 is 1\\pas2\\, corresponding to $\\sim 540$~AU. In the profile of k2 one can also identify the two intensity enhancements to the East and West, seen in Fig.~\\ref{feiih2closeup}a, which deliniate the edges of the conical nebula. Both knots are superimposed on a broad ($\\sim 17\\arcsec \\approx 0.04$~pc) plateau of emission; this is an artifact due to imperfect mosaicing, rendering the middle section of the image (in which the nebula and the jet are located) more noisy than the edges. The absolute coordinates of the three stellar objects (A, B1, and B2; Fig.~\\ref{cb230h2}) and of the two knots (k1, k2; Fig.~\\ref{feiih2closeup}a) in the [FeII]-jet are given in table~\\ref{knotparams}, where we also list the integrated fluxes of the two [FeII]-knots and the entire H$_2$ jet. \\hfill\\break\\noindent In addition to the main jet, a second [FeII]-feature is seen, emanating from the bright nebulosity and oriented in a N-NW-direction; this may indicate the existence of a second (unresolved) outflow, or it is a shock at the surface of the cavity, and thus related to the main jet.  \\begin{table} \\caption[] {List of coordinates and fluxes of the detected objects} \\label{knotparams} \\begin{tabular}{lcccc} \\hline \\multicolumn{1}{c}{Object} & \\multicolumn{1}{c}{$\\alpha_{\\rm 2000}$} & \\multicolumn{1}{c}{$\\delta_{\\rm 2000}$} & \\multicolumn{1}{c}{$F_{\\rm [FeII]}$} & \\multicolumn{1}{c}{$F_{\\rm H_2}$} \\\\ \\multicolumn{1}{c}{CB} & \\multicolumn{1}{c}{($^h$ $^m$ $^s$)} & \\multicolumn{1}{c}{($\\degr$ $\\arcmin$ $\\arcsec$)} & \\multicolumn{2}{c}{(erg cm$^{-2}$ s$^{-1}$)$^{\\dagger}$} \\\\ \\hline 3-k1  &  00 28 42.17 & +56 41 57.36 & $<$ 2.4$-$16 & 3.5$-$14 \\\\ 3-k2  &  00 28 42.12 & +56 41 37.17 & $<$ 2.4$-$16 & 1.4$-$14 \\\\ 3-k3  &  00 28 41.75 & +56 41 34.20 & $<$ 2.4$-$16 & 5.6$-$15 \\\\ 3-k4  &  00 28 42.12 & +56 41 28.56 & $<$ 2.4$-$16 & 2.1$-$14 \\\\ 205-A  & 19 45 24.15 & +27 50 57.18 & -- & --  \\\\ 205-B  & 19 45 23.99 & +27 50 59.25 & -- & --  \\\\ 230-A  &  21 17 38.36 & +68 17 32.87 &  -- & --  \\\\ 230-B1 &  21 17 40.09 & +68 17 32.15 &  -- & --  \\\\ 230-B2 &  21 17 39.96 & +68 17 31.91 &  -- & --  \\\\ 230-H$_2$ &     --       &    --        &  --  & 1.1$-$14 \\\\ 230-k1 &  21 17 38.11 & +68 17 41.04 & 6.4$-$15 &  --  \\\\ 230-k2 &  21 17 38.12 & +68 17 38.08 & 3.2$-$15 &  --  \\\\ \\hline \\multicolumn{5}{l}{$\\dagger$ a$-$b means a$\\times 10^{-b}$} \\end{tabular} \\begin{center} \\end{center} \\end{table}   \\subsection{CB188}  \\begin{figure*} \\resizebox{15cm}{!}{\\rotatebox{270}{ \\includegraphics{0931fig6.eps}}} \\caption[]{Results of the TNG/NICS observations showing the H$_2$ emission, still including continuum emission, towards CB188 (left panel) and CB205 (right panel). The locations of the IRAS point sources are indicated by their uncertainty ellipses. The upper small panel on the far right shows a blow-up of the region around the IRAS source position in CB205. The lower small panel on the far right shows the alleged jet-like feature (Yun et al.~\\cite{yetal93}) being resolved into stars. } \\label{cb188205} \\end{figure*}  In this globule, Launhardt \\& Henning (\\cite{launhen}) detected 1.3~mm continuum emission, in the form of an ``extended source with compact components'' at the location of the IRAS source (indicated by the ellipse in Fig.~\\ref{cb188205}). A YSO was detected at the same location by Yun \\& Clemens (\\cite{yc95}), while Yun \\& Clemens (\\cite{yc94b}) mapped a small ($\\leq 2\\arcmin$) outflow, of which the less extended blue lobe overlaps the red one.  \\noindent Our continuum-subtracted H$_2$ and [FeII]-images do not show any line emission in the field down to the sensitivity limit. We also do not see the cometary diffuse emission around the possible NIR counterpart of the IRAS source that was found by Yun \\& Clemens (\\cite{yc94a}; this would be the bright star seen just outside the IRAS error ellipse in Fig.~\\ref{cb188205}), nor do we detect their source F, which they find embedded in this diffuse emission (see their Fig.~10b). It might be either a spurious detection or a variable object.  \\subsection{CB205}  This globule, also known as L810, has an extended NIR nebulosity at the location of the IRAS source, as seen in Fig.~\\ref{cb188205}. A NIR star cluster is associated with it; at least 10 stars can be seen within $10\\arcsec$ of the IRAS position (see the top-inset of Fig.~\\ref{cb188205}). \\hfill\\break\\noindent The major axis of the nebula is oriented in a N-S direction, which is also the orientation of the compact molecular outflow detected by Yun \\& Clemens (\\cite{yc94b}), of which the red and blue lobes show a significant overlap. This coincidence of orientations suggests a connection between the shape of the nebula and the presence of the outflow. However, no line emission has been found in our continuum-subtracted images, although we cannot rule out that some line emission knots may have been hidden in the more intense parts of the diffuse NIR nebulosity, near the IRAS point source location.  \\noindent The nebula has been studied in the NIR (J,H,K) by Yun et al. (\\cite{yetal93}). The object they identified as the illuminator of the nebula, L810IRS, coincides with our star B in Fig.~\\ref{cb188205}. Yun et al. (\\cite{yetal93}) also report the detection of an elongated jet-like feature, located about $35\\arcsec$ from L810IRS, to the SW. This feature is also visible in all our images (both line and continuum) and is seen to be resolved in a coincidental alignment of stars embedded in diffuse emission (see the lower-inset in Fig.~\\ref{cb188205}, at $\\Delta\\alpha \\sim -20\\arcsec$ and $\\Delta\\delta \\sim -45\\arcsec$). No emission is left after continuum subtraction, further excluding a jet origin for it."
    },
    "0403/astro-ph0403415_arXiv.txt": {
        "abstract": "{ We present a quantitative study of massive stars in the High Excitation Blob N81, a compact star forming region in the SMC. The stellar content was resolved by HST and STIS was used to obtain medium resolution spectra. The qualitative analysis of the stellar properties presented in Heydari-Malayeri et al.\\ (\\cite{papI}) is extended using non-LTE spherically extended atmosphere models including line-blanketing computed with the code CMFGEN (Hillier \\& Miller \\cite{hm98}), and the wind properties are investigated. The main results are the following:  \\begin{itemize} \\item The SMC-N81 components are young ($\\sim$ 0--4 Myrs) O stars with effective temperatures compatible with medium to late subtypes and with luminosities lower than average Galactic O dwarfs, rendering them possible ZAMS candidates.  \\item The winds are extremely weak: with values of the order 10$^{-8}$/10$^{-9}$ \\myr\\, the mass loss rates are lower than observed so far for Galactic dwarfs. Only the recent study of SMC stars by Bouret et al.\\ (\\cite{jc03}) show the same trend. The modified wind momenta ($\\dot{M}$ v$_{\\infty}$ $\\sqrt{R}$) are also 1 to 2 orders of magnitude lower than observed for Galactic stars. Both the mass loss rates and the modified wind momenta are lower than the predictions of the most recent hydrodynamical models. \\end{itemize}  The accuracy of the UV based mass loss rate determination, relying in particular on the predicted ionisation fractions, are carefully examined. We find that $\\dot{M}$ could be underestimated by a factor of up to 10. Even in this unlikely case, the above conclusions remain valid  \\textit{qualitatively}.  The reasons for such weak winds are investigated with special emphasis on the modified wind momenta:  \\begin{itemize} \\item There may be a break-down of the wind momentum - luminosity relation (WLR) for dwarf stars at low luminosity (log L/L$_{\\odot}$ $\\la$ 5.5). However, reasons for such a breakdown remain unknown.  \\item The slope of the WLR may be steeper at low metallicity. This is predicted by the radiation driven wind theory, but the current hydrodynamical simulations do not show any change of the slope at SMC metallicity. Moreover, there are indications that some Galactic objects have wind momenta similar to those of the SMC stars.  \\item Decoupling may take place in the atmosphere of the SMC-N81 stars, leading to multicomponent winds. However, various tests indicate that this is not likely to be the case. \\end{itemize}  The origin of the weakness of the wind observed in the SMC-N81 stars remains unknown. We suggest that this weakness may be linked with the youth of these stars and represents possibly the onset of stellar winds in recently formed massive stars. ",
        "introduction": "\\label{s_intro} Massive stars play key roles in various astrophysical contexts all along their evolution: they ionise ultra-compact HII regions while still embedded in their parental molecular cloud; they create ionised cavities and shape the surrounding interstellar medium during the main fraction of their lifetime; they experience strong episodes of mass loss when they become Luminous Blue Variables and Wolf-Rayet stars, revealing their core and enriching the ISM in products of H and He burning; they end their life as supernovae, producing the heavy elements and releasing large amounts of mechanical energy. During all these phases, massive stars lose mass through winds driven by radiation pressure on metallic lines. This affects not only their evolution (e.g. Chiosi \\& Maeder \\cite{cm86}) but also the surrounding interstellar medium in which the release of mechanical energy can trigger instabilities leading to the collapse of molecular clouds and to star formation. Moreover, bubbles and superbubbles observed on galactic scales are powered by such mass ejections. Hence, various astrophysical fields require the knowledge of quantitative wind properties of massive stars.  Several studies have been carried out in the last two decades to determine these properties. At solar metallicity, the observational determinations (e.g. Howarth \\& Prinja \\cite{hp89}, Puls et al.\\ \\cite{puls96}, Herrero et al.\\ \\cite{hpv00}) are on average in good agreement with the most recent hydrodynamical predictions based on the radiation driven wind theory (Vink et al.\\ \\cite{vink00}), both in terms of mass loss rate and of the modified wind momentum - luminosity relation (WLR, e.g. Puls et al.\\ \\cite{puls96}) which quantifies the strength of the wind. At non solar metallicities, we expect the wind properties to vary with $Z$ due to the modified radiative acceleration through metallic lines. In particular, the mass loss rate should be proportional to $Z^{r}$ (Abbott \\cite{abbott82}, Puls et al.\\ \\cite{psl00}) and the WLR should be shifted towards lower values and should have a steeper slope. The most recent theoretical results predict $r \\sim 0.8$ (Vink et al.\\ \\cite{vink01}) but no change in the slope of the WLR, at least for $Z > 10^{-3} Z_{\\odot}$ (Hoffmann et al.\\ \\cite{tadziu02}, Kudritzki \\cite{kud02}). Observational studies indicate a reduction of the mass loss rate and of the terminal velocity in the Magellanic Clouds, but given the small number of objects studied so far, the behaviour of the WLR at low metallicity is still poorly understood. Several groups are currently analysing stars in sub solar (Crowther et al.\\ \\cite{paul02}, Hillier et al.\\ \\cite{hil03}, Bouret et al.\\ \\cite{jc03}) and super solar (Najarro et al.\\ \\cite{paco} , Figer et al.\\ \\cite{figer}) regions for a better understanding of wind properties in different environments. The present work on SMC-N81 stars takes part in this effort.  The SMC-N81 region belongs to the class of the ``High Excitation Blobs'' (HEB) first introduced by Heydari-Malayeri \\& Testor (\\cite{ht82}). These blobs are compact regions of star formation in the Magellanic Clouds (see Heydari-Malayeri \\cite{mhmiau} for a complete review). They have a typical radius of a few pc and display the features of star forming regions: HII cavities, turbulent structures, ionisation fronts and shocks. Recent HST observations (Heydari-Malayeri et al.\\ \\cite{pap0}) have revealed for the first time its stellar content, 4 of the brightest stars being grouped in the central 2 \\arcsec\\ wide region. Subsequently spectra of the main exciting stars have been obtained with STIS onboard HST. The qualitative analysis of these spectra, presented in Heydari-Malayeri et al.\\ (\\cite{papI}, hereafter paper I), have already revealed interesting properties. First, the stars have been identified as mid O dwarfs with surprisingly low luminosities compared to ``classical'' dwarfs. Second, the UV spectra have shown signatures of very weak winds, even weaker than those usually observed in the SMC. These characteristics have lead Heydari-Malayeri et al.\\ (\\cite{papI}) to propose that the SMC-N81 stars could belong to the class of Vz stars which are massive stars thought to lie very close to the ZAMS (Walborn \\& Parker \\cite{wp92}).  As such the properties of these stars, showing unusually weak winds compared to other SMC O stars, seem already quite interesting. Furthermore the association of these objects with a compact star forming region, presumably indicative of a very young age, allows one also to obtain unique constraints on properties of very young massive stars shortly after their birth. In fact such observations appear crucial for a better understanding of the earliest evolutionary phases of massive stars and to constrain their formation process which is still under debate (they may form by accretion on a protostellar core -- Norberg \\& Maeder \\cite{nm}, Behrend \\& Maeder \\cite{raoul}-- or by collisions between low mass components in dense stellar clusters -- Bonnell et al.\\ \\cite{bonnell}).  With such objectives in mind we have carried out a quantitative study of the UV spectra of the SMC-N81 stars. First results have been presented in Martins et al.\\ (\\cite{lanzarote}). In fact we are able to determine upper limits on the mass loss rates of four O stars in this region, which turn out to be surprisingly low (typically $\\mdot \\la$ a few $10^{-9}$ \\myr) compared to predictions of the radiation driven wind theory, even when taking metallicity effects into account. Although no precise physical explanation is found for this behaviour we strongly suggest that this behaviour is related to the very youth of these massive stars.   The remainder of the paper is structured as follows. Section \\ref{observations} briefly summarises the observations and data reduction. Section \\ref{cmfgen} describes the main ingredients of the modeling. In Sect.\\ \\ref{interstellar} we explain how interstellar lines are taken into account. The main results are given in Sect.\\ \\ref{analysis} and discussed in Sect. \\ref{neb_ste_prop} (nebular and stellar properties) and \\ref{wind_properties} (wind properties). Finally, Sect.\\ \\ref{conclusion} summarises the main results. ",
        "conclusions": "\\label{conclusion} Based on UV spectral obtained with STIS/HST we have analysed the stellar and wind properties of the four main exciting stars of the High Excitation Blob SMC-N81 using extensive calculations of spherically expanding non-LTE line blanketed atmosphere models with the code CMFGEN.  The main results are the following:  \\begin{itemize} \\item[$\\diamond$] {The stellar properties (L, \\teff) indicate that the SMC-N81 components are young ($\\sim$ 0--4 Myrs old) O stars which shows, with perhaps the exception of star 1, a lower luminosity than ``normal'' Galactic O dwarfs. This, together with the closeness to the ZAMS for star 3 and 11, confirms the conclusion of paper I that they may belong to the Vz class (Walborn \\& Parker \\cite{wp92}). }  \\item[$\\diamond$] {The UV spectra of the N81 stars show unusually weak stellar winds. The upper limits on mass loss rates are of the order a few $10^{-9}$ \\myr\\ which is low compared to 1) Galactic stars of the same luminosity and 2) the most recent predictions of $\\dot{M}$ as a function of stellar parameters and metallicity. Point 1) could be qualitatively understood due to the reduced metallicity of the SMC but point 2) indicates that this reduction is higher than expected.}  Although the mass loss rates derived from the UV line analysis are potentially affected by uncertainties in the modeled ionisation fractions, various tests indicate that the above conclusions remain qualitatively valid.    \\item[$\\diamond$] {Our objects show modified wind momenta ($M_{\\odot} v_{\\infty} R^{1/2}$) which are, for the same luminosity $L$, lower by typically two orders of magnitude compared to the ``normal'' O star samples. Similarly low wind momenta have also been found by Bouret et al.\\ (2003) for 3 SMC stars in NGC 346.  The modified wind momentum - luminosity relation of all the SMC objects could be interpreted as showing a break-down at low luminosities or a different slope than the Galactic relation. The current sample of SMC stars may indeed indicate a steeper slope at least for giants and dwarfs, but the scatter is still too large to firmly establish this trend. However, the most recent hydrodynamical models (Vink et al.\\ \\cite{vink01}, Kudritzki \\cite{kud02}, Hoffmann et al.\\ \\cite{tadziu02}) do not predict such a change in the slope between solar and SMC metallicities. Furthermore we present the first indications that some Galactic objects have also low wind momenta comparable to the SMC dwarfs. This also tends to exclude explanations based uniquely on metallicity.}  \\item[$\\diamond$] Possible explanations for a breakdown of the modified wind momentum - luminosity relation at low luminosities are discussed. Ionic decoupling appears unlikely according to various estimates. A failure of the CAK parameterisation in high density atmospheres, discussed by Owocki \\& Puls (\\cite{op99}), might be invoked to explain a lower acceleration in the transsonic region where the mass loss rate is set.  Although the physical mechanism leading to such weak winds remains currently unknown, we speculate that the low mass loss rate is probably intrinsically related to the youth of the stars, possibly testifying of a phase of the ``onset'' of radiatively driven winds in young O stars shortly after their formation.  \\end{itemize}  Further studies of very young massive stars, Vz stars, and related objects with indications of weak winds will be of great interest to attempt to understand these puzzling wind properties and to provide interesting constraints on the development of stellar winds in the early phases of massive star evolution or possibly even on the final phases of their birth.   \\appendix"
    },
    "0403/astro-ph0403690_arXiv.txt": {
        "abstract": "We present three-dimensional simulations of viscous dissipation of AGN induced gas motions and waves in clusters of galaxies. These simulations are motivated by recent detections of ripples in the Perseus and Virgo clusters. Although the sound waves generated by buoyant bubbles decay with distance from the cluster center, we show that these waves can contribute substantially to offsetting the radiative cooling at distances significantly exceeding the bubble size. The energy flux of the waves declines more steeply with radius than the inverse-square law predicted by energy conservation, implying that dissipation plays an important role in tapping the wave energy. We show that such dispersing sound waves/weak shocks are detectable as ripples on unsharp-masked X-ray cluster maps, and point out that the interfaces between the intracluster medium and old bubbles are also clearly detectable in unsharp-masked X-ray maps. This opens up the possibility of detecting fossil bubbles that are difficult to detect in radio emission. This mode of heating is consistent with other observational constraints,  such as the presence of cool rims around the bubbles and the absence of strong shocks. Thus, the mechanism offers a way of heating clusters in a spatially distributed and gentle fashion. We also discuss the energy transfer between the central AGN and the surrounding medium. In our numerical experiments, we find that roughly 65 per cent of the energy injected by the AGN is transferred to the intracluster medium and approximately 25 percent of the injected  energy is dissipated by viscous effects and contributes to heating of the gas. The overall transfer of heat from the AGN to the gas is comparable to the radiative cooling losses. The simulations were performed with the FLASH adaptive mesh refinement code. ",
        "introduction": "The long-standing problem of cooling flow clusters of galaxies, in which the central cooling time is much shorter than the Hubble time, is how to prevent the intracluster medium (ICM) from collapsing catastrophically on a short timescale. The original idea for maintaining the overall cluster stability (Fabian 1994) was to postulate that a certain amount of gas decouples from the flow and does not contribute to the cooling of the  remaining gas. This model would require up to 1000 $M_{\\odot}$ yr$^{-1}$ in mass deposition rates to  guarantee cluster stability. This has been found to be inconsistent with recent {\\it Chandra} (e.g., McNamara et al. 2000, Blanton et al. 2001) and XMM-{\\it Newton} observations (e.g., Peterson et al. 2001, 2003; Tamura et al. 2001). {\\it Chandra} observations reveal a number of clusters with X-ray cavities/bubbles created by the central active galactic nuclei (AGN). It has been suggested by many authors that AGN feedback may play a crucial role in self-regulating  cooling flows (e.g. Churazov et al. 2001, Ruszkowski \\& Begelman 2002, Brighenti \\& Mathews 2003). One of the main outstanding issues is how the AGN heating comes about in detail. In principle, strong  shocks generated by AGN outbursts can dissipate in the ICM and heat the gas. However, imaging observations of cooling flow cores do not give evidence for this mode of heating. Recent {\\it Chandra} observations of two well-known clusters,  the Perseus cluster (Fabian et al. 2003a,b) and the Virgo cluster (Forman et al. 2004), suggest that dissipation of sound waves and weak shocks could be an important source of gas heating --- an idea first proposed by Fabian et al. (2003a). Further support for the idea that viscosity may play an important role in the ICM comes from a recent study of density profiles in clusters (Hansen \\& Stadel 2003).  Recently, a number of papers have described simulations of bubble-heated clusters (e.g., Churazov et al. 2001, Br\\\"{u}ggen et al. 2002, Br\\\"{u}ggen \\& Kaiser 2002, Br\\\"{u}ggen 2003, Quilis et al. 2001). Numerical simulations of viscous dissipation of AGN energy in ICM were previously considered by Ruszkowski et al. (2004) (Paper I) and Reynolds et al. (2004).\\\\ \\indent The main purpose of this paper is to extend our previous work on viscous heating of the ICM by waves to three dimensions. The results of a simulation of viscous dissipation in three dimensions could differ from our previous two-dimensional results given that the amplitudes of waves decrease faster with radius in three dimensions. This would directly affect the spatial distribution of the viscous dissipation rate. Apart from performing our simulations in three dimensions,  we extend our previous analysis and that of Reynolds et al. (2004) to include aspects of heating that were previously neglected.  First, we show that the conclusions drawn from our 2D simulations carry over to three dimensions. Second, our new results include (i) X-ray maps and unsharp-masked X-ray images, (ii) extend the discussion of spatial distribution of energy dissipation and overall heating rate, (iii) present details of the flow of heat between the bubbles and the intracluster medium and (iv) discuss the wave decay rates.\\\\ \\indent The outline of this paper is as follows. In the next section we describe the assumptions of the model. Section 3 presents and discusses our results, focusing on the  spatial distribution of heating. In particular, we discuss the detectability of the ripples/sound waves, their decay rate with distance from the center/bubble surface, the energy transfer from the AGN to the kinetic and thermal energy  of the ICM, and the fraction of the dissipated energy that heats the gas. The fourth section discusses the limitations of our model. The final section summarizes our findings. ",
        "conclusions": "To summarize, we have analyzed the energy deposition in the cluster due to rising bubbles, sound waves and weak shocks. This was motivated by the recent discovery of such waves in the Perseus cluster by Fabian et al. (2003a) and in the Virgo cluster by Forman et al. (2004). We found that the dissipated energy may be comparable to the cooling rate, thereby significantly affecting the cooling flow or even quenching it altogether. We showed that about 65 per cent of the energy injected by the central source can be transferred to the ICM. Approximately 25 per cent of the energy injected by the AGN can be converted to heat, assuming Spitzer viscosity. We discussed the wave decay rates and showed that a significant fraction of wave energy is deposited within the cooling radius. The computed decay rates are consistent with linear theory estimates of the damping length. The damped sound waves or weak shocks are still detectable in unsharp-masked X-ray images. Old bubbles become increasingly difficult to detect in the X-ray maps as the contrast between the rising bubbles and the surrounding gas diminishes. However, apart from sound waves and weak shocks in unsharp-masked X-ray maps, the interfaces between the intracluster medium and old bubbles are also clearly visible. This opens up the possibility of detecting fossil bubbles that are difficult to detect in radio emission."
    },
    "0403/astro-ph0403373_arXiv.txt": {
        "abstract": "The progress of optical astronomy in post-apartheid South Africa is discussed. Particular emphasis is given to the socio-political climate which embraced the idea of a 10-m class telescope as a flagship project that would lead to widespread development in science, technology and education - not only in South Africa, but across the subcontinent. ",
        "introduction": "This account of optical astronomy in South Africa starts where Feast (2002, hereafter Paper I) left off, in 1994 with the first democratic elections and the start of a new era. The end of apartheid offered vastly increased opportunities for international collaborations among individuals and institutions, which the astronomy community was quick to take advantage of. Nevertheless, while the historical strength of astronomy laid a firm foundation for growth and success, the reasons why the discipline thrived and grew in the following decade were complex and essentially political.  In the following I attempt a brief description of the policies and socio-political climate that have nurtured astronomy in South Africa, while acknowledging that no two individuals will see this in the same way and that any such account will be incomplete and probably idiosyncratic.  The government-funded facilities available for optical and infrared astronomy in South Africa are described with an emphasis on the 10-m Southern African Large Telescope (SALT), due to be commissioned in early 2005. No attempt is made to describe detailed scientific projects or results, but it is worth noting that productivity remained high throughout the decade, with the South African Astronomical Observatory (SAAO) annual report, for example, recording well over 100 publications per year from its user community. A detailed account of research at the beginning of the period can be found in the compilation edited by Warner (1995). The challenge for the future is to redirect these efforts towards making effective use of SALT.  Two significant transformations were initiated in South African astronomy during this decade: one involves the change to big telescope astronomy and is only just starting; the other is a broad ``Africanization\" of activities which is far from complete but well underway. Prior to 1994 the main interactions had been with the international community (excluding Africa), particularly that in the UK. Most of the local optical astronomers had been born, and many of them trained, outside of Africa. The post-1994 investment in astronomy came with the assumption that this would change - that astronomers would find ways to interact with and influence South African science and that a cohort of indigenous astronomers would be trained and nurtured. The challenge has been to do this in such a way that these young scientists are the peers of their international contemporaries and not merely tokens to fill quotas. Interestingly this challenge is being met through strengthening collaborations and partnerships, both nationally and internationally. ",
        "conclusions": "I want to finish by contextualizing the spending on astronomy in South Africa through comparing expenditure on SALT with that on the Hubble Space Telescope (HST). The cost of HST at launch in 1990 was \\$1.5 billion (or \\$1.59 billion in 1992 assuming 3 percent inflation) while that of SALT in 2002 terms is about \\$25 million. Using data from the World Bank\\footnote{www.worldbank.org/data/countrydata.html} we see that the USA GDP in 1992 was \\$6.262 trillion, while that of South Africa in 2002 was \\$104.2 billion. The cost of SALT compared to the GDP of South Africa, 0.025 percent, is almost identical to that of HST compared to the GDP of the USA. There can be no doubt that the South African astronomy community is extraordinarily fortunate in having access to this level of support.  The politicization of science is often problematic for scientists, and much has been written about the inevitability of this process in connection with big science projects (e.g. Enard 2002). The above comparison with HST demonstrates that SALT is big science for Africa; the people of South Africa will have expectations of it that are comparable to those of Americans for HST. We South Africans will do well to follow the example of the Space Telescope Science Institute in ensuring that: \\begin{itemize} \\item observing time goes to the astronomers with the best projects, \\item those astronomers are fully empowered to do first rate science with SALT \\item and that the outcomes of their research are made accessible to the public and particularly to young Africans.\\end{itemize}  It seems appropriate to give the last word to the President of South Africa, Thabo Mbeki. In opening the South African Pavilion at the 2000 World Expo in Hanover, Germany, President Mbeki said the following: \\begin{quotation}``{\\it Now, in the small town of Sutherland in the semi-desert Karoo region of our country, we are building a gigantic African eye through which we can view the universe. The construction of the single largest telescope in the southern hemisphere, SALT - as it is called - will mean that in this humble home of our earliest humans, we are also building a vast gateway through which we can observe our earliest stars, learn about the formation of our galaxy and the lives of other worlds so as to give us insights into our future.  We are proud that SALT will not only enable South African scientists to undertake important research, but also provide significant opportunities for international collaboration and scientific partnerships with the rest of the world.}\"\\end{quotation}"
    },
    "0403/astro-ph0403145_arXiv.txt": {
        "abstract": "The center of M83, a barred starburst galaxy with a double nucleus, has been observed in  the CO($J$=2--1) and CO($J$=3--2) lines with the Submillimeter Array. The molecular gas shows a distribution and kinematics  typical for barred galaxies at $\\sim$kpc radii, but reveals unusual kinematics around the double nucleus in the central $\\sim$300 pc. Our CO velocity data show that the visible nucleus in M83 is at least 3\\arcsec\\ (65 pc) away from the galaxy's dynamical center, which most likely coincides with the center of symmetry previously determined in $K$ band and is suggested to host another nucleus. We discovered high-velocity molecular gas associated with the visible off-center nucleus, and also found a steep velocity gradient across it. We attribute these features to a gas disk rotating around the off-center nucleus, which may be the remnant  of a small galaxy accreted by M83. The dynamical mass of this component is estimated to be $3\\times 10^8 \\Msol$ within a radius of 40 pc. The dynamical perturbation from the off-center nucleus may have played a key role in shaping the lopsided starburst. ",
        "introduction": "M83 (NGC 5236) is a nearby face-on barred spiral galaxy with a double nucleus and a nuclear starburst ($D=4.5$ Mpc, $1\\arcsec = 22$ pc; Thim et al. 2003). The galaxy is exceptionally complex in the central 300 pc, despite the symmetric appearance of its bar and spiral arms. In $K$ band, the brightest point, called the `visible nucleus', is offset by  3\\farcs4  (75 pc) from the centroid (i.e., center of symmetry) determined from the isophotes at radii of 0.3 -- 0.7 kpc \\citep[hereafter TTG]{Thatte00}. The centroid suffers from large extinction, and does not show a peak in $K$ band. However, the stellar velocity dispersion peaks both at the visible nucleus and near the isophotal centroid, and suggests a mass of $\\sim$10$^{7}$\\Msol\\ for each. This suggests a double nucleus (TTG). The visible nucleus has a hard power-law spectrum in X-ray due either to an accreting supermassive black hole or to X-ray binaries \\citep{Soria02}. \\citet{Elmegreen98} found a double circumnuclear ring in their $J-K$ color index map. The  `inner' and `outer' rings of 190 pc and 60 pc radius are not concentric. The galaxy hosts a circumnuclear starburst mainly in the `starburst arc' of $\\sim$250 pc length that lies between the two rings \\citep{Gallais91, Harris01}. The galactic center has a large amount of molecular gas, estimated to be $10^{7.5\\mbox{--}8.5}\\Msol$ in $r\\leq300$ pc \\citep{Handa90,Israel01}.   The proximity, face-on configuration, abundant molecular gas, starburst, and double nucleus make M83 one of the most interesting targets to study spiral galaxies and their nuclear activity through molecular gas. We observed M83 as a part of the early science program of 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. }. The new telescope in Hawaii can easily observe the galaxy at $\\delta=-30\\degr$. In this {\\it Letter}, we first describe the observations and the bar-driven gas dynamics in the central 2 kpc,  and then show that the visible nucleus is indeed offset from the dynamical center of M83 and has a velocity feature indicative of its own gas disk. We discuss the origin of the double nucleus, and its relation to the starburst activity. ",
        "conclusions": ""
    },
    "0403/astro-ph0403235_arXiv.txt": {
        "abstract": "We present the analysis of  3 hrs of  a rapid time series of precise stellar radial velocity (RV) measurements ($\\sigma$ = 4.5 m\\,s$^{-1}$) of the cool Ap star $\\beta$ CrB. The integrated RV measurements spanning the wavelength interval 5000-6000\\,{\\AA} show significant variations (false alarm probability = 10$^{-5}$) with a period of 16.21 min ($\\nu$ = 1028.17 $\\mu$Hz) and an amplitude of 3.54 $\\pm$ 0.56 m\\,s$^{-1}$. The RV measured over a much narrower wavelength interval reveals one spectral feature at $\\lambda$6272.0 {\\AA} pulsating with the same 16.21 min period and an amplitude of 138 $\\pm$ 23 m\\,s$^{-1}$. These observations establish $\\beta$ CrB to be a low-amplitude rapidly oscillating Ap star. ",
        "introduction": "Kurtz (1989) first suggested that the cool magnetic Ap star $\\beta$ CrB was a prime candidate to be a rapidly oscillating Ap (roAp) star since it had stellar properties similar to several known oscillating Ap stars (HR 1217, 33 Lib, and $\\gamma$ Equ). The issue of whether $\\beta$ CrB is a pulsating star  has important implications for the excitation mechanism in roAp stars as well as the possible existence of an instability strip for roAp stars.  Several investigators have searched for pulsational variations in $\\beta$ CrB using photometric measurements with null results (Heller \\& Kramer 1988; Kreidl 1991). Recently, Kochukhov et al. (2002, hereafter K02) presented radial velocity (RV) measurements for $\\beta$ CrB spanning the wavelength interval 6105 -- 6190 {\\AA}. Measurements of most lines were constant to a level of 20--30 m\\,s$^{-1}$; however, intriguing evidence for RV variations were found for one spectral feature, Fe I $\\lambda$6165.4 {\\AA}. This feature had an amplitude of 71 $\\pm$ 11 m\\,s$^{-1}$ and a period of 11.5 $\\pm$ 0.5 min. Before we can place $\\beta$ CrB in the family of roAp stars the K02 result must be confirmed since the false alarm probability for the detection was rather high (= 0.016).  Here we present our own precise radial velocity measurements for $\\beta$ CrB and show that it is indeed a rapidly oscillating Ap star, but not with the period found by K02. ",
        "conclusions": "Our precise RV measurements for $\\beta$  CrB establish that this star is with high probability a rapidly oscillating Ap star with an amplitude of 3.54 $\\pm$ 0.56 m\\,s$^{-1}$ and a period of 16.2 $\\pm$ 0.7 min. Although additional observations are needed to confirm our result, we believe that the signal we have detected is real for a number of reasons:  \\begin{itemize} \\item {\\it We derive a very low false alarm probability.}  The probability that this is a false signal due to noise is $\\approx$ 10$^{-5}$ as determined through the bootstrap randomization procedure. This is a more rigorous way of determining the FAP than strictly using the Lomb-Scargle power. This result represents the combined mean velocity of all spectral lines spanning the wavelength interval 5000--6000\\,{\\AA}.  Any systematic or instrumental error would have to affect all wavelength chunks used in the analysis in the same way which seems unlikely.   \\item {\\it The variations are not due to a few outliers.}  We think it is unlikely that a few outliers are causing our signal. (We define an outlier as deviating significantly from the mean RV value and not from the best fit sine wave through the data.) If only a few outliers are driving the power in the periodogram then this would be evident in a high false alarm probability.  For a real periodic signal in the presence of noise, the more data one accumulates  the more significant the detection becomes, even with the presence of a few occasional outliers. Thus the power in a Lomb-Scargle periodogram should increase with increasing number of data points.  This is demonstrated in Figure~\\ref{powerinc} which shows the Lomb-Scargle power at $\\nu$ = 1028.17 $\\mu$Hz (circles) as a function of the number of data points used in the periodogram.  For comparison we generated a synthetic signal consisting of a sine wave with the same period and amplitude as found in our data and sampled in the same manner as the real data. Random noise was added with the same rms scatter as our measurements. The crosses in Figure~\\ref{powerinc}  show that the L-S power of the fake data (with a signal present) increases with increasing number of data points in the same way.  An iodine absorption cell was used to perform our RV measurements because it is designed to eliminate, or at least minimize instrumental effects. However, one could argue that there might be residual instrumental shifts not corrected by the reduction process, or that our instrumental profile modelling somehow introduces a false signal into the data. We do not believe this is the case.  We investigated whether periodic instrumental shifts were present in our data using the wavelength solution calculated from the iodine absorption lines. A Fourier analysis of these shifts showed the strongest period at 176 minutes, and smaller variations with periods of 20 and 24 minutes. All periods were significantly different from the one found in our RV data. Furthermore, we also performed a Fourier analysis of all IP parameters used in our RV determination. We found no variations at the period coinciding with the one found in the RV data. This demonstrates that using the iodine cell technique does an excellent job of excluding instrumental variations.  \\item {\\it The same signal is found in the $\\lambda$6272 {\\AA} feature.}  The spectral orders used to determine the integrated RV did not include the order covering the $\\lambda$6272 {\\AA} feature. Thus this spectral line is not contributing to the signal found in our integrated RV measurements. However, a separate analysis of this  spectral chunk did find the same 16.21 min period and with the same statistical significance as with the integrated RV measurements. Although this is not an independent confirmation (we are using the same data set, but different parts of the spectrum), this result and the above arguments argue strongly that the 16.21 min oscillations in $\\beta$ CrB are real.  \\end{itemize}  \\begin{figure} \\epsfxsize=8.5truecm \\epsffile{figure5.eps} \\caption{The  Lomb-Scargle  power  at $\\nu$ = 1028.17 $\\mu$Hz as a function of number of data points used in the periodogram using real (circles) and fake data (crosses). The  fake data was generated using a sine wave with the same amplitude (3.56 m\\,s$^{-1}$) and period (= 16.21 min) as found in our data and sampled in the same manner. Random noise with the same scatter ($\\sigma$ = 4.5 m\\,s$^{-1}$) as our data was also added to the fake signal. } \\label{powerinc} \\end{figure}  Our RV measurements fail to confirm the results of K02 on the Fe I $\\lambda$6165\\,{\\AA} feature. The wavelength chunk containing this spectral line had an rms scatter of 130 m\\,s$^{-1}$ (due to the weak iodine absorption lines in this wavelength region). However, we had a factor of 3 more measurements than K02. Monte Carlo simulations indicate that we would\thave detected an 11.5 min period with an amplitude of 70 m\\,s$^{-1}$ and a FAP of 10$^{-4}$. An amplitude of 50 m\\,s$^{-1}$  would have been detected with a FAP = 0.01. Thus we would have detected any variations in the Fe I $\\lambda$6165\\,{\\AA} line if they were present at the same amplitude as reported by K02.  We believe that the result of K02 on $\\beta$ CrB is spurious. The FAP is much too high (FAP = 0.016) and our measurements have failed to find the presence of an 11.5 min period in either the integrated RV measurements, or in the analysis of the narrow wavelength chunks. More troubling is that the K02 period of 11.5 $\\pm$0.5 min found in $\\beta$\\,CrB is very close to the values of periods ($\\approx$ 11.7 min) found in several spectral lines in the roAp star 10 Aql using the same instrument.  The most likely explanation for the RV signal found in $\\beta$ CrB by K02 is that it is due to short term uncorrected instrumental variations in the Gecko spectrograph  used for the measurements. Circumstantial evidence for this comes from Matthews \\& Scott (1995) who reported an 11.1 min  period in the RV variations in  $\\gamma$\\,Equ. This period is different from known pulsation modes and has never been confirmed by subsequent RV measurements from other investigations. Although the authors used  a superimposed mercury emission line from an arc lamp to eliminate instrumental shifts, the 11.1 min period could still be a residual instrumental effect due to the stellar light a calibration source having a slightly different optical path (not the case when using the iodine absorption cell).  The RV measurements of K02, on the other hand, were made {\\it without} a simultaneous wavelength calibration (either I$_2$ absorption cell or simultaneous Th-Ar calibration), rather the wavelength calibration used a Th-Ar exposure taken before and after the time series used for the RV measurements. This cannot correct for any instrumental variations on time scales significantly shorter than the time between the two calibration measurements. Because K02 have failed to exclude an instrumental origin with a period of $\\approx$ 11.5 min for the RV variations in $\\beta$ CrB their results on this star (and possibly 10 Aql) should be considered suspect. An investigation of the short-term instrumental shifts of the Gecko spectrograph would be useful.  Our RV measurements for $\\beta$ CrB seem to establish that this star is an roAp star with the lowest RV amplitude (3.5 m\\,s$^{-1})$ and one of the longest periods (16.2 min).  $\\beta$ CrB is also unique among roAp stars in that it shows high amplitude pulsational in only one spectral feature, the $\\lambda$6271.9\\,{\\AA}. This showed significant variations (FAP $=$ 1.5$\\times$10$^{-5} $) with the same 16.2 min period and an amplitude of 138 $\\pm$ 23 m\\,s$^{-1}$. We tentatively identify this feature as a  blend of Ce\\,II and Cr\\,II lines. Why $\\beta$ CrB has such a low RV amplitude, almost no high amplitude spectral lines,  and a longer period mode compared to other roAp stars with similar spectral properties should provide clues as to the origin of the roAp phenomenon. We are currently planning observations of this star over a full rotation period not only to confirm our detection, but to search for rotationally modulated amplitude variations of the pulsations."
    },
    "0403/astro-ph0403529_arXiv.txt": {
        "abstract": "Resonant  active-to-active  ($\\nu_a \\rightarrow  \\nu_a$),  as well  as active-to-sterile   ($\\nu_a  \\rightarrow   \\nu_s$)   neutrino  ($\\nu$) oscillations  can take  place during  the core  bounce of  a supernova collapse. Besides, over this phase, weak magnetism increases antineutrino ($\\bar{\\nu}$) mean  free paths, and thus its  luminosity.  Because the oscillation feeds mass-energy into the target $\\nu$ species, the large  mass-squared difference between species ($\\nu_a  \\rightarrow \\nu_s$) implies a huge amount of energy to be given off as gravitational waves ($L_{\\textrm{GWs}} \\sim 10^{49}$~erg s$^{-1}$), due to anisotropic but coherent $\\nu$ flow over the oscillation length. This asymmetric $\\nu$-flux is driven by both the spin-magnetic and the {\\it universal spin-rotation} coupling.  The novel contribution of  this paper stems from 1) the new computation of the anisotropy parameter $\\alpha \\sim 0.1-0.01$, and 2) the use of the tight constraints from neutrino experiments as  SNO and KamLAND,   and    the   cosmic   probe   WMAP,    to   compute   the gravitational-wave emission during neutrino oscillations in supernovae core  collapse  and bounce.  We  show that  the  mass  of the  sterile neutrino  $\\nu_s$ that can  be resonantly  produced during  the flavor conversions makes it a good  candidate for dark matter as suggested by Fuller et {\\it  al.} (2003).  The new spacetime  strain thus estimated is  still several  orders of  magnitude larger  than those  from $\\nu$ difussion (convection  and cooling)  or quadrupole moments  of neutron star matter. This  new feature turns these bursts  the more promissing supernova gravitational-wave signal that may be detected by observatories as LIGO,  VIRGO, etc., for distances  far out to the  VIRGO cluster of galaxies.  \\vskip 0.8 truecm ",
        "introduction": "{\\it  Supernovae  neutrinos  and  gravity waves.---}  That  outflowing neutrinos ($\\nu$s) from a  supernova (SN) generate gravitational waves (GWs)  was firstly  pointed out  by  Epstein (1978). However, over       the        first      $\\sim       10$       milliseconds (ms) (Mayle, Wilson \\& Schramm 1987; Walker \\& Schramm 1987) after the SN core  bounce the central density gets  so high  that no radiation  nor even $\\nu$s  can escape, they are thus frozen-in and strongly coupled to the neutron matter ($N^0$) as  described by the  Lagrangean (see Kusenko \\& Postma 2002 for this dynamics)   \\be L^{int}_{N^0\\leftrightarrow            \\nu}           = \\frac{G_F}{\\sqrt{2}} \\left[ \\bar{N}^0    \\gamma_\\mu(1-\\gamma_5)N^0 \\right] \\left\\{ \\bar{\\psi} \\gamma^\\mu(1-\\gamma_5)\\psi \\right\\}\\; ,\\label{interact} \\ee  with  the $\\nu$  field  $(\\psi)$ satisfying  the time-dependent  Dirac equation  \\be  \\left[i\\gamma^0  \\partial_0  + i\\gamma^\\alpha  \\partial_\\alpha  + \\rho(t)   v_\\beta   \\gamma^\\beta   \\left(\\frac{1-\\gamma_5}{2}\\right) - m_\\nu\\right] \\psi = 0. \\label{dirac-mass} \\ee   At  this  phase  the  whole  proto-neutron  star  (PNS)  dynamics  is dominated by gravity alone, and  can be appropriately described by the general relativistic Oppenheimer-Volkoff equation for both the $N^0$ + $\\nu$ fluid (see Mosquera Cuesta 2002). As discussed by Mayle, Wilson \\& Schramm (1987); and Walker \\& Schramm (1987), it is over this early transient that most $\\nu$ flavor conversions are expected to  resonantly take place  and consequently the  super strong GWs burst from the oscillation  process to be released.  GWs from this decoupling  has  been suggested  to  likely  be  the ultimate  process responsible for the neat kick given  to a nascent pulsar during the SN collapse  (Mosquera Cuesta 2000; 2002).   The contention of this \\textit{paper}  is a) to pave, in the framework of  general   relativity  (GR),   the  pathway  to   this  fundamental astrophysical process of generation  of GWs from $\\nu$ oscillations in a  PNS. b)  to demonstrate,  by taking  into account  experimental and observational  constraints,  that $\\nu$  oscillations  during SN  core bounce do produce  GWs of the sort predicted  by Einstein's GR theory, and more  crucial yet,  c) to  stress that these  bursts are  the more likely SN GWs-signals to  be detected by interferometric observatories as LIGO, VIRGO, GEO-600, etc.  We speculate that such a signal perhaps might have  been detected  during the SN1987a  event, despite  the low sensitivity of the detectors at the time. Some claims in this direction were presented by Aglietta, Amaldi, Pizzella, et al.\\cite{amaldi89}, and related papers. ",
        "conclusions": "One can see that if  $\\nu$ flavor conversions indeed take place during SN  core bounce  inasmuch as  they  take place  in our  Sun and  Earth (Smirnov 2002), then GWs should be released during the transition. The GWs  signal from  the process is  expected to irradiate  much more energy than current mechanisms figured out to drive the NS dynamics at birth do. A luminosity  this large (Eq.(\\ref{GWs-luminosity})) would turn these bursts the strongest GWs signal  to be detected from any SN that may come to occur, futurely, on  distances up to the VIRGO cluster, $R \\sim [10-20]$ Mpc.  It is stressed  that this signal will still be the stronger one  from a  given SN, even  in the  worst case in  which the probability of  $\\nu$ conversion is three orders  of magnitude smaller then the estimated  in the present paper. In proviso,  we argue that a GWs signal  that strong could  have been detected during  SN1987a from the  Tarantula  Nebula  in  the  Milky Way's  satellite  galaxy  Large Magellanic Cloud, despite  of the low sensitivity of  the detectors at the epoch. In such a case,  the GWs burst must have been correlated in time with the earliest arriving neutrino burst constituted of some active species given  off during the  very early oscillation  transient where some  $\\nu_e$s  went   into  $\\nu_{\\mu}$s,  $\\nu_\\tau$s  or  $\\nu_s$s. Thenceforth, it could be of  worth to reanalyze the data collected for from that  event taking careful follow  up of their  arrival times, if appropriate timing was available at that moment."
    },
    "0403/astro-ph0403325_arXiv.txt": {
        "abstract": "The Sloan Digital Sky Survey has validated and made publicly available its Second Data Release.  This data release consists of 3324 square degrees of five-band ($u\\,g\\,r\\,i\\,z$) imaging data with photometry for over 88 million unique objects, 367,360 spectra of galaxies, quasars, stars and calibrating blank sky patches selected over 2627 degrees of this area, and tables of measured parameters from these data. The imaging data reach a depth of $r \\approx 22.2$ (95\\% completeness limit for point sources) and are photometrically and astrometrically calibrated to 2\\% rms and 100 milli-arcsec rms per coordinate, respectively.  The imaging data have all been processed through a new version of the SDSS imaging pipeline, in which the most important improvement since the last data release is fixing an error in the model fits to each object.  The result is that model magnitudes are now a good proxy for point spread function (PSF) magnitudes for point sources, and Petrosian magnitudes for extended sources. The spectroscopy extends from 3800\\AA\\ to 9200\\AA\\ at a resolution of 2000.  The spectroscopic software now repairs a systematic error in the radial velocities of certain types of stars, and has substantially improved spectrophotometry.  All data included in the SDSS Early Data Release and First Data Release are reprocessed with the improved pipelines, and included in the Second Data Release.  Further characteristics of the data are described, as are the data products themselves and the tools for accessing them. ",
        "introduction": "The Sloan Digital Sky Survey (SDSS; York \\etal\\ 2000) is an imaging and spectroscopic survey of the high Galactic latitude sky visible from the Northern hemisphere. The principal survey goals are to measure the large-scale distribution of galaxies and quasars and to produce an imaging and spectroscopic legacy for the astronomical community.  The SDSS data have been used in well over 200 refereed papers to date on subjects ranging from the colors of asteroids (Ivezi\\'c \\etal\\ 2002) to magnetic white dwarfs (Schmidt \\etal\\ 2003) to structures in the Galactic halo (Newberg \\etal\\ 2003) to the star-formation history of galaxies (Kauffmann \\etal\\ 2003) to Type II quasars (Zakamska \\etal\\ 2003) to the large-scale distribution of galaxies (Pope \\etal\\ 2004; Tegmark \\etal\\ 2004).  The survey uses a dedicated 2.5m telescope with a three-degree field of view at Apache Point Observatory, New Mexico.  A 120 mega-pixel camera (Gunn \\etal\\ 1998) images in five broad bands ($u,g,r,i$ and $z$; Fukugita \\etal\\ 1996; Stoughton \\etal\\ 2002) on clear moonless nights of good seeing. These data are photometrically calibrated using an auxiliary 20-inch telescope with a $40^\\prime \\times 40^\\prime$ imager, which determines the photometricity of each night (Hogg \\etal\\ 2001), and measures the extinction and photometric zeropoint using a network of % standard stars (Smith \\etal\\ 2002).  The imaging data are processed through a series of pipelines that locate and measure the properties of all detected objects (Lupton \\etal\\ 2001) and carry out photometric and astrometric calibration (Pier \\etal\\ 2003).  From the resulting catalogs of objects, complete catalogs of galaxies (Eisenstein \\etal\\ 2001; Strauss \\etal\\ 2002) and quasar candidates (Richards \\etal\\ 2002) are selected for spectroscopic followup, and are assigned to spectroscopic tiles of diameter 3 degrees (Blanton \\etal\\ 2003).  Spectroscopy is performed on nights with moonlight, mild cloud cover, and/or poor seeing using a pair of double spectrographs with coverage from 3800--9200\\AA, and resolution $\\lambda / \\Delta \\lambda$ of roughly 2000.  A plug plate for each tile holds 640 optical fibers of $3''$ entrance aperture which feed the spectrographs, together with eleven coherent fiber bundles to image guide stars.  Because of the diameter of the cladding holding the optical fibers, spectroscopy cannot be carried out for objects separated by less than $55''$ on a given plate. ",
        "conclusions": ""
    },
    "0403/astro-ph0403439_arXiv.txt": {
        "abstract": "We present UKIRT UIST spectra of Sakurai's Object (=V4334~Sgr) showing CO fundamental band absorption features around 4.7~$\\mu$m. The line-centres are at heliocentric radial velocity of $-$170$\\pm$30~km~s$^{-1}$. The number and relative strengths of the lines indicate a CO gas temperature of $400\\pm100$~K and CO column density of 7$^{+3}_{-2}\\times10^{17}$~cm$^{-2}$. The gas was moving away from the central star at an average speed of $\\sim$290$\\pm$30~km~s$^{-1}$ in 2003~September. The lines appeared sometime between mid 1999 (well after the opaque dust shell formed) and mid 2000 and may have been somewhat more blue--shifted initially than they are now. The observed CO velocity and temperature indicate the continued presence of a fast wind in the object, previously seen in the He~{\\sc i} 1.083~$\\mu$m line beginning just prior to massive dust formation, and more recently in atomic and ionized lines. The dust continuum is consistent with a temperature of 350$\\pm$30~K, indicating continued cooling of the shell. The similar CO temperature suggests that the bulk of the CO absorption occurs just outside of the dust continuum surface. ",
        "introduction": "\\label{sec-intro}  Stellar evolution theory accounts reasonably well for the effects of thermal pulses on the development of stars on the Asymptotic Giant Branch.  Less well constrained are the effects of a late thermal--pulse (LTP, once the star has shed its envelope) or very--late thermal--pulse (VLTP), once the star has begun to descend the white dwarf cooling track).  However, indications are that perhaps 10 to 20\\%\\ of low-- and intermediate--mass stars undergo LTP or VLTP. It appears that the subsequent ``born--again'' evolution across the HR diagram is short lived (perhaps a few centuries, but perhaps as short as a few decades; see \\citealt{Iben83,Iben96, Lawlor03}), and consequently very few are observable at any one time.  Sakurai's Object (= V4334~Sgr) is probably the first example of a VLTP observable with non--optical instruments in the immediate post--flash epoch.  Sakurai's~Object was first identified as possibly undergoing a helium--shell--flash on 1996~February~23 \\citep{Nakano96, Benetti96}, but it is clear that it started to increase in brightness in mid--1994 \\citep{Takemizawa97}.  In 1995 we began an infrared (IR) spectroscopic monitoring programme using the United Kingdom Infrared Telescope (UKIRT) and the cooled grating spectrometer CGS4 \\citep{Mountain90}; more recently we have used the UKIRT Imaging Spectrometer UIST \\citep{Ramsay-Howat00}. Observations are carried out throughout the star's observable period each year and the latest data, taken in 2003~September, are presented here. ",
        "conclusions": "\\label{sec-conclusion}  A recent spectrum of Sakurai's Object has clearly revealed the presence of highly blue--shifted absorption lines of the fundamental band of CO. The CO was first noted in a spectrum from 2003~September~8.  Re--examination of previous lower resolution spectra shows that the CO features must have arisen between 1999~May~4 and 2000~April~17.  The blue--shift of the lines is consistent with gas moving away from the central star at $\\sim$300~km~s$^{-1}$. The number and relative strengths of the lines suggest a CO gas temperature of $400\\pm100$~K and CO column density of 7$^{+3}_{-2}\\times10^{17}$~cm$^{-2}$ (the lower value corresponding to the higher temperature), and linewidths no greater than 25~km~s$^{-1}$ (FWHM). We rule out interstellar CO as the source of the absorption lines, but the circumstellar CO could be formed shortly before discovery or existed for some time prior to emerging from the dusty shroud. % We note that the CO velocity is additional evidence for the existence of a fast wind in this object, although we cannot make a direct connection with changes in the central star. Taking the 1--5~$\\mu$m spectrum overall we find that the dust continues to cool, and the temperatures of the CO--bearing and dust--bearing materials are similar. Hence the two components presumably lie at a similar distance from the central star. As the CO is moving more rapidly than the dust continuum surface, this suggests an ongoing wind replenishing the CO in the region of the dust."
    },
    "0403/astro-ph0403263_arXiv.txt": {
        "abstract": "Aluminum and other metal abundances were determined in 21 red giants in the globular clusters NGC~6752 and M80 as part of a larger study to determine whether the aluminum distribution on the red giant branch is related to the second parameter effect that causes clusters of similar metallicity to display different horizontal branch morphologies. The observations were obtained of the Al~I lines near 6700 {\\AA} with the CTIO Blanco 4-m telescope and Hydra multi-object spectrograph. The spectra have a resolving power of 18000 or 9400, with typical S/N ratios of 100-200. Mean [Fe/H] values obtained from the spectra are --1.58 for NGC~6752 and --1.73 for M80; this represents the first spectroscopic iron abundance determination for M80. Both NGC~6752 and M80 display a spread in aluminum abundance, with mean [Al/Fe] ratios of +0.51 and +0.37, respectively.  No trend in the variation of the mean Al abundance with position on the giant branch is discernible in either cluster with our small sample. ",
        "introduction": "\\label{sec:intro}  Over 20 years have passed since \\citet{NCFD81} first showed aluminum abundance variations on the red-giant branch (RGB) of the globular cluster NGC~6752, and yet the nature of these inhomogeneities remains a mystery. Observations since then continue to show aluminum (and sodium and magnesium) variations in other clusters \\citep[see, e.g.][]{CD81,WLO87,SHFS87,DSS92, ND95,PSKL96,ZWB96,S96a,Kraft97,SKS97,K98,Ivans1999,CN2000,Ivans2001,RC2001,GratAl2001, GBNF2002}. Although variations in carbon and nitrogen were previously known, these could be described through a simple mixing mechanism proposed by \\citet{SM79}, where rotationally induced meridional circulation currents could carry nuclearly processed materials such as C, N, and O from around the hydrogen-burning shell (H shell) of a red giant across the radiative zone to the outer convective envelope. Heavier elements such as Mg, Na, and Al weren't thought to be processed around the H shell, and any variations in them were taken as evidence that some kind of primordial pollution affected the surface abundances \\citep[see, e.g.,][]{CD81}. Continued work on key nuclear reaction rates \\citep{Champagne98,DD90,Iliadis90, CBS93,Blackmon95,Iliadis96}, however, suggested that these elements could be processed around the H shell under the same conditions that the CN and ON nuclear cycles operated. Using these results, as well as the widely accepted rates of \\citet{CF88}, separate groups showed that it is possible to account qualitatively for the observed variations that showed Na, Al, and N anticorrelated with C, O, and, in some cases, Mg \\citep{LHS93, CSB96, DW96, CSB98, DDNW98}. The challenge has always been describing the results {\\em quantitatively}; in particular, producing [Al/Fe] as high as 1.5 dex without overproducing [Na/Fe], and depleting $^{24}$Mg to the observed levels in M13 \\citep{S96b}, all while remaining within the acceptable proton-capture rates \\citep{LHZ97,NACRE,Powell99,UNC}. Despite any latitude afforded by the reaction rate uncertainties \\citep[see, e.g.,][]{CSB98}, the mixing theory still relies on non-solar abundance ratios in the star prior to mixing \\citep{DDNW98,CN2000}, thus, at least partly relying on primordial influences. Theories other than meridional circulation have been put forth \\citep{LHZ97,FAK99, AFK2001,DW2001}, but a detailed discussion of each is beyond the intent of this paper.  Regardless of the physics behind any mixing mechanism, the question still remains: are the Al (and Na, Mg) variations caused by mixing, primordial sources, or a combination of both? We might be able to answer this question by looking at another long outstanding problem in globular cluster astronomy, namely, the second-parameter effect. First pointed out by \\citet{SW67} and \\citet{vdB67}, the second-parameter effect refers to the phenomenon where the horizontal-branches (HB) of two clusters with similar metallicity (the first parameter) have markedly different color distributions. Possible second parameters that have been investigated include age, initial helium abundance, CNO abundance, and mass loss \\citep[see, e.g.,][]{Faulkner66,Renzini77,Chaboyer98}, among others, but none is applicable to all clusters. One recent suggestion hypothesizes that if deep mixing (i.e., mixing that penetrates the H shell) occurs, then helium will be brought to the surface affecting the HB morphology by making mixed stars both bluer and brighter (Sweigart 1997a,b).  Unfortunately, helium cannot be measured in RGB stars because of their low surface temperatures, and helium settling on the HB precludes an accurate measurement in the hotter stars. However, models by \\citet{CSB98}, which use standard (albeit, uncertain) reaction rates, show that aluminum can be produced only in the H shell of giants within the last magnitude of the RGB, implying that an increase in helium in the envelope must also be accompanied by an increase in aluminum. From this we can postulate that, if [Al/Fe] variations are produced internally and not primordially, Al could be a good surrogate to measure He mixing on the bright RGB, and if He mixing is indeed the second parameter, a relationship should exist between the ratio of Al-strong to Al-normal stars (the Al ratio) and the ratio of blue to red HB stars (the HB ratio). It is the {\\em distribution} of abundances as a function of magnitude that is critical in determining the whether this correlation exists and what its cause might be.  This current work attempts to provide some fresh data on clusters that have been historically under-studied. Both M80 (NGC~6093) and NGC~6752 possess blue HBs relative to other clusters at similar metallicity. For example, \\citet{Ferr98} compared Hubble $U,V$ M80 photometry directly with M13 and M3, with the result that M80's HB is very similar to M13, while M3 lacked the extended blue tail of the other two clusters. Meanwhile, \\citet{GCLSA1999} presented extensive Str\\\"{o}mgren photometry of NGC~6752, M13, and M3, showing extensive blue tails in the former two compared with the latter. Neither M80 nor NGC~6752 had been studied extensively for abundances until \\citet{GratAl2001} and \\citet{GBNF2002} determined [Al/Fe] for 39 stars in NGC~6752 in total, extending the data of \\citet{ND95}. These were significant results because they probed less evolved stars near the main-sequence turnoff, on the subgiant branch, and at the base of the RGB, which are below the point that mixing theories predict that aluminum can be produced. Using a non-LTE analysis, \\citet{GratAl2001} observed dwarfs with [Al/Fe] as low as $-0.76$ dex from the Al~I resonance lines, and subgiants with [Al/Fe] as high as $+0.86$ dex from the doublet at ${\\lambda}{\\lambda}$8773/74~{\\AA}. In fact, the resonance line analysis for the dwarfs yielded [Al/Fe]~$=~-0.18~{\\pm}~0.15$~(s.e.m.)~dex, while the subgiants gave [Al/Fe]~$=~+0.29~{\\pm}~0.11$~dex (s.e.m.). The results for the dwarfs are uncertain due to the difficulty of analyzing resonance lines with non-LTE corrections of as much as $+0.6$~dex; yet, the results are still surprisingly low. What causes the drastic change from the main sequence to the subgiant branch: atmospheric effects, the different choice of lines, or an actual physical phenomenon? The \\citet{GBNF2002} results are more in line with the results of \\citet{ND95} and with other clusters. We discuss implications of the aluminum data in NGC~6752 in more detail in section~\\ref{sec:final_look}.  M80, on the other hand, has no published abundances and is in need of further investigation, especially given its similarities to M13 in metallicity and HB morphology.   The rest of the paper is outlined according to the following: We begin with a description of the observations in Section~\\ref{sec:obs}, followed by our data reduction techniques in Section~\\ref{sec:ccd}. We then discuss membership criteria in Section~\\ref{sec:rv}. After culling the data, we show the results of our abundance analysis in Section~\\ref{sec:analysis} and give our final conclusions in Section~\\ref{sec:conclude}. ",
        "conclusions": "\\label{sec:conclude} \\subsection{Summary} \\label{sec:summary}  We begin the conclusions by summarizing our results: \\begin{itemize} \\item We observed 21 giants in the globular clusters M80 and NGC~6752 with spectra of sufficient quality to determine abundances. \\item Both M80 and NGC 6752 display a spread of aluminum abundances, having mean abundances of $+0.37$~dex and $+0.51$~dex, respectively.  The abundance spreads are 0.43 ({\\stdv}) dex for M80 and 0.36 ({\\stdv}) dex for NGC~6752. \\item The aluminum data cannot resolve the discrepancy in the [Al/Fe] values near the main-sequence turnoff and the subgiants as observed by \\citet{GratAl2001}. \\item No trends between [Al/Fe] ratios and magnitude are discernible in our small sample. \\item The mean [Fe/H] value for M80 is $-1.73$, which is the first spectroscopic determination for this cluster.  For NGC~6752, the mean [Fe/H] is $-1.58$, which is consistent with previous results. \\item The Fe-peak elements chromium and nickel follow iron closely. \\item The [Ca/Fe] enhancements are consistent with the ${\\alpha}$-enhancements observed in other clusters. \\item The [Eu/La] ratio is constant for both clusters at $+0.42$ dex, which does not appear unusual; however, both [Eu/Fe] and [La/Fe] are enhanced in M80 relative to NGC~6752. \\end{itemize}  \\subsection{A Final Look At Aluminum} \\label{sec:final_look}  Now that we've derived aluminum abundances in M80 and NGC~6752, are we any closer to answering whether or not the variations are the result of primordial pollution, mixing processes, or both? Perhaps when combined with the results of \\citet{GBNF2002} we can get some insight into just how complicated this problem really is (we do not include the Gratton data here since they were biased by use of the $c_{1}$ index to select stars.) Given the differences in the quality of the data between the two studies, and the still small number of stars that have been analyzed, it is difficult to draw firm conclusions. If mixing is an ongoing phenomenon, then one would expect that the Al ratio would increase with decreasing magnitude, as it appears to in M13 \\citep{CN2000}; unfortunately, the small numbers make this difficult to discern at this point.  An interesting result also comes from both \\citet{GratAl2001} and \\citet{GBNF2002}, who show that stars on the subgiant/lower red giant branch may have just as much Al in them as stars on the upper RGB, where mixing is theoretically possible. These results might suggest that all Al anomalies are the result of pollution, but this raises yet another question: Can it be demonstrated that a globular cluster main sequence star, subgiant, or lower red giant exists that has strongly enhanced Al and is not depleted in C and/or O? That is, if the CNO variations are the result of deep mixing and the Al (and Na) anomalies are primordial, there should be stars that show uncorrelated abundance patterns. But, if the CNO and Na-Al anomalies are both caused by primordial scenarios, then why does evidence exist, in clusters where it's been studied, that the total C$+$N$+$O remains constant from star to star independent of the Na or Al content \\citep{NBS81, CAP88,DCCB91}, and how does one explain $^{12/13}$C ratios near the CN cycle equilibrium value in many other clusters \\citep{SS91,S96b,ZWB96, BSKL97,BSSBN97}? On the other hand, if the anomalies are all created {\\em in-situ}, what is the physics behind them that can occur without completely contradicting well-established basic theories of stellar evolution? It used to be that one could side with one scenario or the other, then the data became more complicated, and one would say that it was somehow a mixture of both primordial and evolutionary scenarios. Maybe in the end that will be the final conclusion, but we are far from proving it; too many questions remain to be answered and too many observations remain to be made."
    },
    "0403/astro-ph0403055_arXiv.txt": {
        "abstract": "We study  the  off-equilibrium effects of inflaton on the dynamics of  primordial perturbations in the $O(N)$ model. A self-consistent off-equilibrium formalism is employed to investigate the evolution of the inflationary background field and its fluctuations with the back-reaction effects. We find two observable remains left behind the off-equilibrium processes: the running spectral index of primordial density perturbations and the correlations between perturbation modes in phase space, which would serve as  the imprints  to probe the epoch of inflation, even beyond. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403219_arXiv.txt": {
        "abstract": "{ We report the first successful application of the astrometric color-induced displacement technique (CID, the displacement of the photocenter between different bandpasses due to a varying contribution of differently colored components to the total light), originally proposed by \\citet{Wielen-1996} for discovering unresolved binary stars. Using the Sloan Digital Sky Survey (SDSS) Data Release 1 with $\\sim 2.5\\,10^6$ stars brighter than 21$^m$ in the $u$ and $g$ bands, we select 419 candidate binary stars with CID greater than 0.5 arcsec. The SDSS colors of the majority of these candidates are consistent with binary systems including a white dwarf and any main sequence star with spectral type later than $\\sim$K7. The astrometric CID method discussed here is complementary to the photometric selection of binary stars in SDSS discussed by \\citet{Smolcic-2004:a}, but there is considerable overlap (15\\%) between the two samples of selected candidates. This overlap testifies both to the physical soundness of both methods, as well as to the astrometric and photometric quality of SDSS data. } ",
        "introduction": "It is believed that 50\\% of all stars belong to multiple systems \\citep{Heintz-1969:a}.  Nevertheless, being aware that a specific star is a binary is always useful because either one throws it out of the sample or updates the model to describe it and begins some follow-up observations!  Whether one deals with stellar evolution or galactic dynamics, binaries always receive some special considerations.  So, it is important to be able to detect the binary nature of a star at an early stage of an investigation by any possible means.  Besides spectroscopy, photometry, and interferometry, astrometry coupled to photometry has lately emerged as a way of revealing the binary nature of a source \\citep{Wielen-1996}.  That method relies upon either a change in the position of the source as its brightness varies (Variability-Induced Movers, VIM) or a photometric-band dependence of the position (Color-Induced Displacement, CID).  Whereas VIM requires several observations along the brightness variation cycle and only works when at least one component is variable, it only takes one image in each band to identify a CID and can be done for non-variable stars.  Despite the number of multi-band photometric surveys, none so far has carried sufficiently accurate astrometry in at least two distinct bands.  This lack of observations has been recently alleviated by the Sloan Digital Sky Survey.  The Sloan Digital Sky Survey \\citep[SDSS;][ and references therein]{York-2000:a,Abazajian-2003:a} is revolutionizing stellar astronomy by providing homogeneous and deep ($r < 22.5$) photometry in five passbands \\citep[$u$, $g$, $r$, $i$, and $z$;][]{Fukugita-1996:a,Gunn-1998:a,Hogg-2001:a,Smith-2002:b} accurate to 0.02 mag \\citep{Ivezic-2003:a}. Ultimately, up to 10,000 deg$^2$ of sky in the Northern Galactic Cap will be surveyed. The survey sky coverage will result in photometric measurements for over 100 million stars and a similar number of galaxies. Astrometric positions are accurate to better than 0.1 arcsec per coordinate (rms) for point sources with $r<20.5^m$ \\citep{Pier-2003:a}, and the morphological information from the images allows robust star-galaxy separation to $r \\sim$ 21.5$^m$ \\citep{Lupton-2003:a}.  Using the SDSS data, we report on the first successful identification of Color-Induced Displacement (hereafter CID) binaries.  The underlying ideas of that method are given in Sect.~\\ref{sect:cid}.   Sect.~\\ref{sect:simu} describes the simulation that allowed us to optimize the screening of the data described in Sect.~\\ref{sect:data}.  In Sect.~\\ref{sect:results}, we present our results and compare them with those of \\citet{Smolcic-2004:a}, who have recently used color selection to identify a stellar locus made of white dwarf+M dwarf binaries. . ",
        "conclusions": "The color induced displacement method described by \\citet{Wielen-1996} as a way of detecting binaries has been successfully applied to the first public release of the SDSS data.  We identify about 400 systems whose changes in position are essentially consistent with a white dwarf coupled to a lower end (later than $\\sim$K7) main-sequence star.  We therefore expect $\\sim2\\,000$ CID binaries at the completion of the SDSS observation campaign.  This identification of binaries is an independent confirmation of the color based results of \\citet{Smolcic-2004:a}.  However, whereas they had a lower bound on $g-r$ of 0.3, the astrometric criterion allows us to identify candidate binaries down to $g-r=-0.4$.  On the other hand, color selection they utilized is more sensitive to binaries with angular separations smaller than the sensitivity of the CID method.  Though the approach has proven to give results, its efficiency is extremely low.  Whereas \\citet{Marchal-2003:a} quote at least 30\\%\\ of binaries among M stars (the percentage grows with the mass of the star along the main-sequence), only 0.02\\%\\ are detected through their CID effect.  It is noteworthy that with such a low fraction, the CID binaries do not affect the overall SDSS astrometric precision.  Because of their much better astrometric precision (typically a few $\\mu$as), space-based astrometry missions like SIM and Gaia will eventually supersede the SDSS results presented here.  According to a Gaia preparatory study \\citep{Arenou-2001:a}, the latter could, for instance, detect a M0 companion to a G0 dwarf star at a $3\\sigma$ level at a separation as low as 2.3 mas.  In terms of separations, this is $\\sim 200$ times better than the sensitivity of the CID method applied to the SDSS data."
    },
    "0403/astro-ph0403505_arXiv.txt": {
        "abstract": "We here report results of an {\\it INTEGRAL} observation of the X--ray burst and atoll source Ser X-1 performed on May 2003. The object was observed for a total of 400 ks but nearly 8$^\\circ$ off-axis due to the amalgamation with an observation of SS 433, the pointing target source. Ser X-1 was detected up to 30 keV with unprecedented positional accuracy for a high-energy emission; a sharp spectral drop is evident beyond this energy. Significant variability is seen in the 20--30 keV light curve. Comparison with previous observations indicates that the source was in its high (banana) state and displayed a soft spectrum during the {\\it INTEGRAL} pointing. A (non simultaneous) broadband radio-to-$\\gamma$--rays broad-band spectral energy distribution for Ser X-1 is also presented for the first time. ",
        "introduction": "The low-mass X--ray binary Ser X-1, or 4U 1837+04 is known to host a neutron star as the accreting object; it is classified as an Atoll source (e.g., Liu et al. 2001). Archival {\\it EXOSAT} data (Seon \\& Min 2002) showed this object in the banana (i.e. high intensity) state during the observations. More recent {\\it BeppoSAX} and {\\it RXTE} pointing (Oosterbroek et al. 2001) caught the source while it was again in a high activity state, with an unabsorbed flux (1--200 keV) of 8.0$\\times$10$^{-9}$ erg cm$^{-2}$ s$^{-1}$. Its spectrum was well described by a combination of a blackbody-like and Comptonization (Titarchuk 1994) models to account for the observed hard tail. A reflection component could not be excluded, but the data quality could not provide a definitive conclusion. Up to now, however, the X--ray emission observed from Ser X-1 with the most recent high-energy missions was never well representative of the hard (island) state which is generally seen when Atoll sources are undergoing the low intensity phase (Barret 2001).  Several X--ray bursts (lasting tens of seconds at most) were also detected from Ser X-1; moreover, during 2001, {\\it BeppoSAX} pinpointed a very long ($\\sim$4 hours) X--ray burst (Cornelisse et al. 2002), making this source join the group of `superbursters' (see Kuulkers 2003 for a review). By studying X--ray bursts observed with {\\it Einstein}, Christian \\& Swank (1997) deduced a distance to the source of 8.4 kpc. This implies a 1--200 keV luminosity of 6.7$\\times$10$^{37}$ erg s$^{-1}$ during the {\\it BeppoSAX} observation of Oosterbroek et al. (2001), which means roughly one third of the Eddington luminosity for a neutron star.  The optical counterpart to Ser X-1, located in a crowded stellar field (Thorstensen et al. 1980), was correctly identified by Wachter (1997) and, subsequently, spectroscopically confirmed and studied by Hynes et al. (2004). Very recently, the radio counterpart was discovered with the VLA (Migliari et al. 2004).  We here report on a observation of Ser X-1 performed with the INTErnational Gamma--RAy Laboratory ({\\it INTEGRAL}; Winkler et al. 2003) on May 2003, i.e. less than 7 months after the launch of this spacecraft. The high spectral sensitivity of the high-energy instruments onboard this satellite are optimal to study the behaviour of the hard X--ray tail of this source in case of its presence. A more complete analysis of these data can be found in Masetti et al. (2004). ",
        "conclusions": ""
    },
    "0403/astro-ph0403675_arXiv.txt": {
        "abstract": "V407 Vul (RX J1914.4+2456) is a candidate double-degenerate binary with a putative 1.756 mHz (9.5 min) orbital frequency. In a previous timing study using archival ROSAT and ASCA data we reported evidence for an increase of this frequency at a rate consistent with expectations for gravitational radiation from a detached ultracompact binary system.  Here we report the results of new {\\it Chandra} timing observations which confirm the previous indications of spin-up of the X-ray frequency, and provide much tighter constraints on the frequency derivative, $\\dot\\nu$. We obtained with {\\it Chandra} a total of 90 ksec of exposure in two epochs separated in time by 11.5 months. The total time span of the archival ROSAT, ASCA and new {\\it Chandra} data is now $\\approx 10.5$ years. This more than doubles the interval spanned by the ROSAT and ASCA data alone, providing much greater sensitivity to a frequency derivative. With the addition of the Chandra data an increasing frequency is unavoidable, and the mean $\\dot\\nu$ is $7.0 \\pm 0.8 \\times 10^{-18}$ Hz s$^{-1}$. Although a long-term spin-up trend is confirmed, there is excess variance in the phase timing residuals, perhaps indicative of shorter timescale torque fluctuations or phase instability associated with the source of the X-ray flux. Power spectral searches for periods longward of the 9.5 minute period do not find any significant modulations, however, the sensitivity of searches in this frequency range are somewhat compromised by the dithering of the Chandra attitude. The observed spin-up is of a magnitude consistent with that expected from gravitational radiation decay, however, the factor of $\\approx 3$ variations in flux combined with the timing noise could conceivably result from accretion-induced spin-up of a white dwarf.  Continued monitoring to explore correlations of torque with X-ray flux could provide a further test of this hypothesis. ",
        "introduction": "Ultracompact binary systems could provide a promising means to observe directly the influence of gravitational radiation on orbital evolution. Moreover, such systems would be ideal sources for detection with spaced based gravitational radiation observatories (such as the planned NASA/ESA LISA mission), opening up the possibility for detailed studies of compact interacting binaries.  In recent years a pair of candidate ultracompact systems; V407 Vul (also known as RX J1914.4+2456) and RX J0806+1527 (hereafter J0806) have been proposed.  These objects were first discovered by ROSAT (Motch et al.  1996; Israel et al. 1999; Beuermann et al. 1999), and initially were suggested to be members of a ``soft'' class of Intermediate Polars (IPs), with the X-ray periods of 569 and 321 s, respectively, representing the putative spin periods of the accreting white dwarfs.  Since their discovery extensive follow-up observations have identified the optical counterparts (Ramsay et al. 2000; Israel et al. 2002; Ramsay, Hakala \\& Cropper 2002). Their soft X-ray spectra, the shape of the X-ray modulation, the phasing of the X-ray and optical modulations, the lack of additional longer periods, and the lack of strong optical emission lines have all called into question their IP credentials (for a discussion see Cropper et al. 2003). However, Norton, Haswell \\& Wynn (2004) have argued that an IP interpretation is still plausible if the systems are stream-fed, pole-switching accretors (ie. no accretion disk), and are viewed from a nearly face-on geometry.  It was Cropper et al. (1998) who first suggested that V407 Vul might be a double-degenerate compact binary.  They proposed a synchronized, magnetic accretor (polar-like) model with accretion powering the X-ray flux. A non-magnetic variant was subsequently proposed by Marsh \\& Steeghs (2002).  In this Algol-like model, the accretion stream impacts directly onto a non-magnetic primary, and the spins are not necessarily synchronized with the orbit.  An interesting alternative not requiring accretion was proposed by Wu et al. (2002). They suggested a unipolar inductor model, analogous to the Jupiter - Io system (Clarke et al.  1996). If these systems are indeed compact, and thus the observed X-ray period is the orbital period, then an important discriminating factor is the magnitude and sign of the orbital evolution. If the systems are accreting stably from degenerate donors, the expected evolution is for the orbit frequency to decrease. In a previous study we (Strohmayer 2002) used archival ROSAT and ASCA data to explore the evolution of the 1.756 mHz X-ray frequency of V407 Vul over an $\\approx 5$ yr time period, and found evidence for a positive frequency derivative, $\\dot\\nu$, with a magnitude consistent with simple expectations for gravitational radiation induced decay of a circular orbit. Since a measurement of the frequency evolution places severe constraints on possible models, it is crucial to confirm the initial indications of orbital decay and place tighter constraints on $\\dot\\nu$.  In this paper we present the results of new {\\it Chandra} observations which confirm an increase in the X-ray frequency and allow us to place much tighter limits on $\\dot\\nu$. In \\S 2 we describe the {\\it Chandra} observations and the data extraction and analysis.  In \\S 3 we discuss our phase coherent timing study, and we show that the inclusion of the {\\it Chandra} data conclusively indicates a positive $\\dot\\nu = 7.0 \\pm 0.8 \\times 10^{-18}$ Hz s$^{-1}$. We also discuss the flux variability of the source and excess variance (timing noise) in the phase residuals. In \\S 4 we discuss the implications of our findings for the nature of V407 Vul. We conclude in \\S 5 with a brief summary and goals for future observations. ",
        "conclusions": "Our new {\\it Chandra} data confirm that the 1.756 mHz X-ray frequency of RX V407 Vul is increasing at a rate of $\\approx 7 \\times 10^{-18}$ Hz s$^{-1}$.  Although this rate is consistent with that expected from gravitational radiation losses in a detached ultracompact binary, an IP interpretation in the context of accretion onto a spinning white dwarf cannot yet be strongly ruled out.  In some ways the new {\\it Chandra} results have only deepened the mystery surrounding V407 Vul (and by implication, its sister source, RX J0806.3+1527). After a decade of observations we still do not know the nature of the source with any certainty.  However, a number of future observations could help provide the solutions. Deep pointings with XMM could probe more sensitively for an unseen orbital period, and further timing observations with {\\it Chandra} will establish constraints on the dependence of torque fluctuations on X-ray flux. Finally, if these do not suffice to crack the mystery, a detection with a spaced-based gravitational radiation observatory (such as the NASA/ESA LISA mission) would provide definitive evidence for an ultracompact system."
    },
    "0403/astro-ph0403396_arXiv.txt": {
        "abstract": "{Collimated outflows from accreting white dwarfs have a vital role to play in the study of astrophysical jets.}    \\addkeyword{Accretion, Accretion Disks} \\addkeyword{Binaries: Symbiotic} \\addkeyword{Stars: Winds, Outflows} \\addkeyword{White Dwarfs}   \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403443_arXiv.txt": {
        "abstract": "\\xmmn EPIC observations have revealed a bright point-like X-ray source in the nearby Magellanic-type galaxy NGC 55.  At the distance of NGC 55, the maximum observed X-ray luminosity of the source, designated as XMMU J001528.9-391319, is $L_x \\sim 1.6 \\times 10^{39}$ erg s$^{-1}$, placing the object in the ultraluminous X-ray source (ULX) regime.  The X-ray lightcurve exhibits a variety of features including a significant upward drift over the 60 ks observation.  Most notably a series of X-ray dips are apparent with individual dips lasting for typically 100--300 \\,s. Some of these dips reach almost 100 percent diminution of the source flux in the 2.0--4.5 \\,keV band. The EPIC CCD spectra can be modelled with two spectral components, a very soft powerlaw continuum ($\\Gamma \\approx 4$) dominant below 2 keV, plus a multi-colour disc (MCD) component with an inner-disc temperature $kT \\approx 0.8$ \\,keV.  The observed upperward drift in the X-ray flux can be attributed to an increase in the level of the MCD component, whilst the normalisation of the powerlaw continuum remains unchanged. The dipping episodes correspond to a loss of signal from both spectral components, although the blocking factor is at least a factor two higher for the MCD component. XMMU J001528.9-391319 can be considered as a candidate black-hole binary (BHB) system. A plausible explanation of the observed temporal and spectral behaviour is that we view the accretion disc close to edge-on and that, during dips, orbiting clumps of obscuring material enter our line of sight and cause significant blocking or scattering of the hard thermal X-rays emitted from the inner disc.  In contrast, the more extended source of the soft powerlaw flux is only partially covered by the obscuring matter during the dips. ",
        "introduction": "\\chandra and \\xmmn provide powerful facilities for studying the X-ray properties of nearby galaxies.  A focus of recent work in this area has been the ultraluminous X-ray source (ULX) phenomenon, namely point-like X-ray sources located outside the nucleus of the galaxy with X-ray luminosities apparently in excess of $L_x > 10^{39}$ erg s$^{-1}$ (\\citealt{roberts00}; \\citealt{colbert02}; \\citealt{miller03}). It is entirely plausible that sources with X-ray luminosities at, or just above, this threshold are mass-transfer binaries containing a stellar mass black-hole (3--20 M$_{\\odot}$) radiating at close to the Eddington limit. Supporting evidence is provided by the fact that many ULXs display similar characteristics to those of established black-hole binaries (BHBs) (\\eg \\citealt{makishima00}).  However, the nature of the subset of ULXs with X-ray luminosities in excess of a few $\\times 10^{39}$ erg s$^{-1}$ is less certain, since these could be systems harbouring intermediate mass black-holes \\citep{colbert99}, radiating anisotropically \\citep{king01} or possessing truly super-Eddington discs \\citep{begelman02}.  It has been estimated that an accreting black-hole, as opposed to a neutron star, is present in at least 10\\% of all bright X-ray binaries (XRBs)\\citep{mcclintock03}.  At the present time there are only 18 dynamically-confirmed stellar-mass BHBs, most of which were discovered as X-ray novae in our own galaxy \\citep{mcclintock03}.  To this sample we can add a further 20 or more candidate objects which exhibit all the characteristics of black-hole systems \\citep{mcclintock03}.  Using \\chandra and {\\it XMM-Newton}, an individual bright binary X-ray source can be studied out to a distance of about 10 Mpc, hence studies of nearby galaxies have the potential for greatly extending our knowledge of luminous XRBs of all types, including black-hole systems. For example, M31 has been a prime target for recent observations (\\eg \\citealt{kong02}; \\citealt{shirey01}; \\citealt{osborne01}), and at least one good BHB candidate has been identified on the basis of its X-ray properties (RX J0042.3+4115; \\citealt{barnard03}).  The other major Local Group galaxy, M33, also hosts many discrete X-ray sources (\\eg \\citealt{haberl01}), including the most luminous persistent X-ray source in the Local Group (M33 X-8; \\citealt*{trinchieri88}). This source is another good black-hole candidate, with recent \\chandra observations revealing characteristics consistent with accretion onto a $> 5 M_{\\odot}$ object \\citep{laparola03}.  Yet a further example is the discovery of an eclipsing XRB in NGC 253 \\citep{pietsch03}.  In the present paper we discuss the properties of the brightest X-ray source detected in NGC 55. This source sits right on the boundary of the ``normal'' binary/ULX categorisation and, on the basis of its X-ray luminosity and spectral properties, is most probably a black-hole system.  NGC 55 is a member of the nearby Sculptor group of galaxies, located in the region of the South Galactic Pole at a distance of 1.78 Mpc \\citep{kara03}. It is morphologically similar to the Large Magellanic Cloud but viewed edge-on with its bar pointing almost along the line of sight ($i=90 \\dg$, \\citealt{tully88}).  NGC 55 has previously been studied in the X-ray band through \\rosat PSPC (\\citealt*{read97}; \\citealt*{schlegel97}) and HRI \\citep{roberts97} observations.  The PSPC data revealed seven bright point-like X-ray sources coincident with the galaxy (\\citealt{schlegel97}).  A subsequent re-analysis of the PSPC data, in conjunction with the HRI data, revealed 25 X-ray point sources coincident with, or in close proximity to, the disc of the galaxy \\citep{roberts97}.  The \\rosat observations showed one particular object, located $\\sim7'$ to the east of the main bar complex, to be several times brighter than any other X-ray source in the galaxy (Source 7 of \\citealt{schlegel97}; Source 6 of \\citealt{read97}; and Source N55-14 of \\citealt{roberts97}).  The \\rosat PSPC data further revealed this source to be spectrally soft (bremsstrahlung temperature $kT \\sim 0.8 - 1.0$ \\,keV or powerlaw photon index $\\Gamma \\sim 3 - 4$) and mildly absorbed (\\nh $\\sim 2 - 4 \\times 10^{21} \\atpcm$) with a derived X-ray luminosity of $\\sim 7 \\times 10^{38} \\ergsec$ in the 0.1 -- 2.4 \\,keV \\rosat band, adopting a distance of 1.78 Mpc.  Crucially, both long- and short-term variability were seen, suggesting this is a luminous accretion-powered XRB.  Here we revisit this source using new, high quality \\xmmn observations to investigate its spectral and temporal behaviour.  \\begin{figure*} \\begin{center} \\scalebox{3}{{\\includegraphics[width=50mm,angle=0]{figure1psxv.ps}}} \\end{center} \\caption{{\\it Left panel:} The \\xmmn image of the NGC 55 field in a broad (0.3--10.0 \\,keV) bandpass. The field centre is at RA $00^{h} 15^{m} 18.0^{s}$, Dec $-39^{\\circ} 13' 33''$ (J2000) and the image size is $40 \\times 40$ arcminute$^2$. The position of XMMU J001528.9-391319 is highlighted by the arrow. {\\it Right panel:} The equivalent optical DSS-2 (red) image  with the rotation angle aligned precisely to a north-south projection.  The white contour, which is derived from a lightly smoothed version of the X-ray image (using a circular Gaussian mask with $\\sigma = 1$ pixel $ = 4''$), corresponds to a surface brightness of 16 count pixel$^{-1}$. North is up and East is to the left in each case.} \\label{images} \\end{figure*} ",
        "conclusions": "\\xmmn observations have revealed a luminous XRB in the nearby galaxy NGC 55. On the basis of its X-ray luminosity and X-ray spectral properties we believe this object is most likely a black-hole system. The lightcurve reveals very interesting spectral variability including, most notably, pronounced dips. Future observations may reveal whether these dips show a pattern of occurrence consistent with an underlying orbital period as is the case for the well-studied BHBs GRO J1655-40 and Cyg X-1.  However, detailed investigation of the spectral variations which accompany the dips represents a challenge even for the \\xmmn instrumentation."
    },
    "0403/astro-ph0403169_arXiv.txt": {
        "abstract": "A method is presented for computing the 6-D phase-space density $f(\\Bx,\\Bv)$ and its PDF $v(f)$ in an N-body system.  It is based on Delaunay tessellation, yielding $v(f)$ with a fixed smoothing window over a wide $f$ range, independent of the sampling resolution. It is found that in a gravitationally relaxed halo built by hierarchical clustering, $v(f)$ is a robust power law, $v(f) \\propto f^{-2.5 \\pm 0.05}$, over more than 4 decades in $f$, from its virial level to the current resolution limit.  This is valid for halos of different sizes in the $\\Lambda$CDM cosmology, indicating insensitivity to the initial-fluctuation power spectrum as long as the small-scale fluctuations were not completely suppressed. By mapping $f$ in position space, we find that the high-$f$ contributions to $v(f)$ come from the ``cold\" subhalos within the parent halo rather than the halo central region and its global spherical profile.  The $f$ in subhalos near the halo virial radius is more than 100 times higher than at the halo center, and it decreases gradually with decreasing radius.  This indicates phase mixing due to mergers and tidal effects involving puffing up and heating.  The phase-space structure provides a sensitive tool for studying the evolution of subhalos during the buildup of halos. One wishes to understand why the substructure adds up to the universal power law in $v(f)$.  It seems that the $f^{-2.5}$ behavior is related to the hierarchical clustering process and is not a general result of violent relaxation. ",
        "introduction": "\\label{sec:intro}  Dark-matter halos are the basic entities in which luminous galaxies form and live. They dominate the gravitational potential and have a crucial role in determining the galaxy properties.  While many of the systematic features of halo structure and kinematics have been revealed by $N$-body simulations, the origin of these features is still not understood, despite the fact that they are governed by simple Newtonian gravity.  The halo density profile $\\rho(r)$ is a typical example.  It is found in the simulations to have a robust non-power-law shape (originally Navarro, Frenk \\& White 1997, NFW; Power \\etal 2003; Hayashi \\etal 2004 and references therein), with a log slope of $-3$ at large radii, varying gradually toward $-1$ or even flatter at small radii.  The slope shows only a weak sensitivity to the cosmological model and the initial fluctuation power spectrum (e.g. Colin \\etal 2003; Navarro \\etal 2004), indicating that its origin is due to a robust relaxation process rather than specific initial consitions.  In particular, violent relaxation (Lynden-Bell 1967) may be involved in shaping up the density profile, but we have no idea why this profile has the specific NFW shape.  The properties of the velocity dispersion tensor is another puzzle. The velocity dispersion profile is slightly rising at small radii and slightly falling at large radii but is rather flat overall (Huss, Jain \\& Steinmetz 1999a; 1999b).  The profile of the anisotropy parameter $\\beta(r)$ indicates near isotropy at small radii that is developing gradually into more radial orbits at large radii (Colin \\etal 2000). For a sperical system in equilibrium, the $\\sigma(r)$ and $\\beta(r)$ are related to $\\rho(r)$ via the Jeans equation, but it is not at all clear why $\\sigma(r)$ or $\\beta(r)$ have these specific shapes.  An interesting attempt to address the origin of the halo profile has been made by Taylor \\& navarro (2001), who measured a poor-man phase-space density profile by $f_{\\rm TN}(r) =\\rho(r)/\\sigma(r)^3$, and found that it displays an approximate power-law behavior, $f_{\\rm TN} \\propto r^{-1.87}$, over more than two decades in $r$.  Using the Jeans equation, they showed that this power law permits a whole family of density profiles, and that a limiting case of this family is a profile similar to NFW, but with an asymptotic slope of $-0.75$ as $r \\rightarrow 0$.  This scale-free behavior of $f_{\\rm TN}(r)$ is intriguing, and it motivates further studies of halo structure by means of phase-space density.  The simulations of the $\\Lambda$CDM cosmology also reveal that the halos are bulit by a rougly self-similar hierarchical clustering process, where smaller building blocks accrete and merge into bigger halos. At every snapshot, every halo contains a substructure of subhalos on top of a smooth halo component that has been tidally stripped from an earlier generation of substructure.  Some of the important dynamical processes invloved in this hierarchical halo buildup are understood qualitatively, including dynamical friction, tidal stripping and mergeres.  However, a complete understanding of how these processes work in concert to produce the halo structure and kinematics is lacking.  Attempts have been made to explain an inner density cusp using toy models of dynamical stripping and tidal effects during the halo buildup by mergers (e.g. Syer \\& White 1998; Dekel, Devor \\& Hetzroni 2003; Dekel \\etal 2003). However, a similar halo density profile seems to be produced also in simulations where substructure has been artificially suppressed (Moore \\etal 1999b; Alvarez, Shapiro \\& Martel 2002), indicating that the process responsible for the origin of this density profile might be a more robust feature of gravity.  The issue of halo substructure has become timely both observationally and theoretically. Tidal streams associated with dwarf satellite galaxies are observed in the halos of the Milky Way and M31 and reveal their histories (Ibata \\etal 2001; this proceedings). Gravitational-lenses provide preliminary indicatons for the presence of substructure in halos at the level predicted by the $\\Lambda$CDM scenario (Dalal \\& Kochaneck 2002). In contrast, the observed number density of dwarf galaxies seems to be significantly lower, thus posing a ``missing dwarf problem\" (Klypin \\etal 1999; Moore \\etal 1999a). Also, the ``angular-momentum problem\" of disk galaxies (e.g. Navarro \\& Steinmetz 2000; Bullock \\etal 2001a) is probably associated with the evolution of substructure in halos (Maller \\& Dekel 2002; Maller, Dekel \\& Somerville 2002). While these problems necessarily invlove baryonic proceses, understanding the gravitational evolution of substructure is clearly a prerequisite for solving them.  Aiming at the origin of halo structure, we report here on a first attempt by Arad, Dekel \\& Klypin (2004, ADK) to address directly the halo phase-space structure.  The fundamental quantity in the dynamical evolution of gravitating systems is the full, 6D, phase-space density $f(\\Bx,\\Bv)$, which intimately relates to the underlying Vlasov equation, and lies behind any (violent) relaxation process that gives rise to the virialized halo structure.  Ideally, one would have liked to compute $f$ free of assumptions regarding spherical symmetry, isotropy, or any kind of equilibrium, but computing densities in a 6D space is a non-trivial challenge.  The state-of-the-art N-body simulations, with more than million particles per halo, allow for the first time an attempt of this sort. ADK developed a succesful algorithm for measuring $f(\\Bx,\\Bv)$, and studied its relevant properties and the associated systematic and random uncertainties. They then applyied this algorithm to simulated virialized halos in the $\\Lambda$CDM cosmology, and obtained two surprising new results. First, the phase-space volume distribution of $f$ is a universal power-law, valid in all virialized halos that form by hierarchical clustering. Second, this power law is not directly related to the overall density profile, but is rather driven by the halo substructure.  We thus learn that $f(\\Bx,\\Bv)$ provides a useful tool for studying the hierarchical buildup of dark-matter halos and the evolution of substructure in them. ",
        "conclusions": "\\label{sec:conc}  It is important to verify that these results are not numerical artifacts. Based on the error analysis and tests with mock datasets, we believe that the $v(f)$ measured by the DTFE algorithm genuinely reflects the true phase-space properties of the given $N$-body system over a broad range of $f$. The question is whether the phase mixing suffered by the subclumps is an artifact of numerical effects such as few-body relaxation, leading to underestimated inner densities and/or overestimated internal velocities (Binney 2003).  The apparent agreement between simulations run with different codes and different resolutions is encouraging. In order to specifically address the effect of two-body relaxation, we intend to run twice a simulation of the same halo with the same number of particles but with a different force resolution (ongoing work with F. Stoehr).  Assuming that the simulations genuinely reflect the true physical behaviour, the origin of the robust power-law shape of $v(f)$ from the merging substructure becomes a very interesting theoretical issue.  As demonstrated in \\S5, a simple model using the mass function and the scaled profiles of the general halo population in the $\\Lambda$CDM scenario does not reproduce the correct power law.  This, and the apparent trend of the $f$ spikes with radius, indicate that the structural and kinematical evolution of the subhaloes in the parent halo are important. Studies of tidal heating and stripping may be found useful in this modelling.  It would be interesting to follow the phase-space evolution and the contribution to the overall $v(f)$ by a single, highly resolved subhalo, or many of those, as they orbit within the parent halo and approach its center. This may help us understand the nature of the interaction between the parent halo and its subhaloes, and the origin of the $v(f)$ power law (ongoing works with E. Hayashi and with B. Moore).  We saw that the power-law behavior of $v(f)$ is limited to the virial regime.  It would be interesting to learn how this shape evolves in time as the halo virializes.  A preliminary study (to be concluded and reported in another paper) indicates that in the intermediate-$f$ regime the $v(f)$ of a pre-virialized system is significantly flatter than $f^{-2.5}$, while in the high-$f$ regime it drops in a much steeper way.  The $f^{-2.5}$ behavior seems to be a feature unique to virialized systems.  We learned that in the haloes that are built by hierarchical clustering, the power-law behavior $v(f)\\propto f^{-2.5}$ reflects the halo substructure. It would be interesting to find out whether this power-law behavior actually requires substructure, or it is a more general phenomenon of virialized gravitating systems, valid independently of substructure.  One way to answer this question would be to analyse simulated haloes in which all fluctuations of wavelengths smaller than the halo scale were removed, resulting in a smooth halo formed by monolithic collapse, with no apparent substructure in the final configuration. As described in \\S1, such haloes are known to still have NFW-like density profiles in real space, and one wonders whether they also have the magic power-law $v(f)$. There are preliminary indications for a steeper $v(f)$ in this case (Arad, Dekel \\& Moore, in preparation). If confirmed, it would indicate that the $f^{-2.5}$ behavior, while insensitive to the exact slope of the initial power spectrum, is unique to the hierarchical clustering process, and is not a general result of violent relaxation.  Our current results are just first hints from what seems to be a promising rich new tool for analysing the dynamics and structure of virialized gravitating systems. The analysis could become even more interesting when applied to haloes including the associated gaseous and stellar components."
    },
    "0403/hep-th0403132_arXiv.txt": {
        "abstract": "{ The exact computation of asymptotic quasinormal frequencies is a technical problem which involves the analytic continuation of a Schr\\\"odinger--like equation to the complex plane and then performing a method of monodromy matching at the several poles in the plane. While this method was successfully used in asymptotically flat spacetime, as applied to both the Schwarzschild and Reissner--Nordstr\\o m solutions, its extension to non--asymptotically flat spacetimes has not been achieved yet. In this work it is shown how to extend the method to this case, with the explicit analysis of Schwarzschild de Sitter and large Schwarzschild Anti--de Sitter black holes, both in four dimensions. We obtain, for the first time, analytic expressions for the asymptotic quasinormal frequencies of these black hole spacetimes, and our results match previous numerical calculations with great accuracy. We also list some results concerning the general classification of asymptotic quasinormal frequencies in $d$--dimensional spacetimes. } ",
        "introduction": "A long time has passed since research first focused on analyzing the linear stability of four dimensional black hole solutions in general relativity \\cite{regge-wheeler, zerilli-1}. However, it was not until very recent times that this stability problem was addressed within a $d$--dimensional setting \\cite{kodama-ishibashi-1, kodama-ishibashi-2, kodama-ishibashi-3}. These papers tried to be as exhaustive as possible, studying in detail the perturbation theory of spherically symmetric black holes in $d$--dimensions and allowing for the possibilities of both charge and a background cosmological constant. Having thus acquired a list of stable black hole solutions, the next question to address within this problem are quasinormal modes---the damped oscillations which describe the return to the initial configuration, after the onset of a linear perturbation (see \\cite{nollert, kokkotas-schmidt} for reviews).  Besides their natural role in the perturbation theory of general relativity, quasinormal modes have recently been focus of much attention following suggestions that they could have a role to play in the quest for a theory of quantum gravity \\cite{hod, dreyer}. The idea is to look at those special modes which are infinitely damped, and thus do not radiate. It was suggested in \\cite{hod} that an application of Bohr's correspondence principle to these asymptotic quasinormal frequencies could yield new information about quantum gravity, in particular on the quantization of area at a black hole event horizon. It was further suggested in \\cite{dreyer} that asymptotic quasinormal frequencies could help fix certain parameters in loop quantum gravity. Both these suggestions lie deeply on the fact that the real part of the asymptotic quasinormal frequencies is given by the logarithm of an integer number, a fact that was analytically shown to be true, for Schwarzschild black holes in $d$--dimensional spacetime, in \\cite{motl, motl-neitzke}. A question of particular relevance that immediately follows is whether the suggestions in \\cite{hod, dreyer} are universal or are only applicable to the Schwarzschild solution. Given the mentioned analysis of \\cite{kodama-ishibashi-1, kodama-ishibashi-3}, one has at hand all the required information to address this problem and compute asymptotic quasinormal frequencies of $d$--dimensional black holes. A preliminary clue is already present in \\cite{motl-neitzke}, where the analysis of the four dimensional Reissner--Nordstr\\o m solution yielded a negative answer: the asymptotic quasinormal frequencies obeyed a complicated relation which did not seem to have the required form. While extending this result to both the $d$--dimensional and the extremal  Reissner--Nordstr\\o m solutions did not pose great obstacles \\cite{natario-schiappa}, an extension of the analytical techniques in \\cite{motl-neitzke} to non--asymptotically flat spacetimes proves to be a greater challenge. It is the goal of this paper to carry out an extension of the techniques in \\cite{motl-neitzke} to non--asymptotically flat spacetimes, with the explicit analysis of Schwarzschild de Sitter and large Schwarzschild Anti--de Sitter black holes, both in four dimensions. The detailed study of these solutions in $d$--dimensions will appear elsewhere \\cite{natario-schiappa}, including charged solutions in asymptotically de Sitter and asymptotically Anti--de Sitter spacetimes, as well as an analysis of the implications of our results on what concerns the proposals of \\cite{hod, dreyer}, dealing with the application of quasinormal modes to quantum gravity.  It is important to stress that even if the ideas in \\cite{hod, dreyer} turn out not to be universal, it is still the case that quasinormal frequencies will most likely have a role to play in the quest for a theory of quantum gravity. Indeed, quasinormal frequencies can also be regarded as the poles in the black hole greybody factors which play a pivotal role in the study of Hawking radiation. Furthermore, the monodromy technique introduced in \\cite{motl-neitzke} to analytically compute asymptotic quasinormal frequencies was later extended, in \\cite{neitzke}, so that it can also be used in the computation of asymptotic greybody factors. It was first suggested in \\cite{neitzke} that the results obtained for these asymptotic greybody factors could be of help in identifying the dual conformal field theory which microscopically describes the black hole, and these ideas have been taken one step forward with the recent work of \\cite{krasnov-solodukhin}. It remains to be seen how much asymptotic quasinormal modes and greybody factors can help in understanding quantum gravity.  Let us conclude this introduction with some generics concerning quasinormal frequencies (we refer the reader to the upcoming \\cite{natario-schiappa} for a full list of conventions and details). Later, in section 2, we shall compute asymptotic quasinormal frequencies for a Schwarzschild de Sitter black hole in four dimensional spacetime. Our results will also be shown to match earlier numerical computations with great accuracy. In section 3, we shall study large Schwarzschild Anti--de Sitter black holes in four dimensions, and analytically compute their asymptotic quasinormal frequencies. Again, our results match earlier numerical computations to great accuracy. We end with some comments concerning the general classification of asymptotic quasinormal frequencies in $d$--dimensional spacetimes \\cite{natario-schiappa}.  For a four dimensional Schwarzschild black hole, one has the asymptotic quasinormal frequencies  $$ \\lim_{n \\to + \\infty} \\omega_{n} \\sim [ {\\mathrm{offset}} ] + i n [ {\\mathrm{gap}} ] + {\\mathcal{O}} \\left( \\frac{1}{\\sqrt{n}} \\right), $$  \\noindent where the real part of the offset is the frequency of the emitted radiation, and the gap are the quantized increments in the inverse relaxation time. Here, the gap is given by the surface gravity. One can try to extend this analysis to more general situations and also include spacetimes with two horizons, but then generic results become much harder to obtain \\cite{mmv-1, padmanabhan, mmv-2, choudhury-padmanabhan}. We shall take the time dependence for the perturbation to be $e^{i\\omega t}$, so that ${\\mathbb{I}}{\\mathrm{m}} (\\omega) > 0$ for stable solutions. There is also a reflection symmetry $\\omega \\leftrightarrow - \\bar{\\omega}$ which changes the sign of ${\\mathbb{R}}{\\mathrm{e}} (\\omega)$. In this case, our quasinormal mode conventions are the following (see \\cite{natario-schiappa} for a full list of conventions in $d$--dimensions). The perturbation master equations of \\cite{kodama-ishibashi-1, kodama-ishibashi-3} can be cast in a Schr\\\"odinger--like form as  \\begin{equation} \\label{schrodinger} - \\frac{ d^{2} \\Phi_{\\omega}}{dx^{2}} (x) + V (x) \\Phi_{\\omega} (x) = \\omega^{2} \\Phi_{\\omega} (x), \\end{equation}  \\noindent where the potential will vary according to the specific case at hand. The boundary conditions are the usual: incoming waves at the black hole horizon and outgoing waves at infinity (or at the cosmological horizon, for the asymptotically de Sitter case)\\footnote{For the asymptotically Anti--de Sitter situation things will be different.}. These can be written as  \\begin{eqnarray*} \\Phi_{\\omega} (x) &\\sim& e^{i\\omega x}\\;\\, {\\mathrm{as}}\\;\\, x \\to - \\infty, \\\\ \\Phi_{\\omega} (x) &\\sim& e^{-i\\omega x}\\;\\, {\\mathrm{as}}\\;\\, x \\to + \\infty, \\end{eqnarray*}  \\noindent where $x$ is the tortoise coordinate. Indeed, if the metric is chosen as $g = - f(r)\\ dt \\otimes dt + {f(r)}^{-1}\\ dr \\otimes dr + r^{2} d\\Omega_{2}^{2}$, with parameters $M = \\MM$ for the black hole mass and $\\Lambda = 3 \\LL$ for the background cosmological constant, then at any (event or cosmological) horizon, $f(R_{H})=0$. One can expand near the horizon $f(r) \\simeq (r-R_{H}) f'(R_{H}) + \\cdots$, and it follows for the tortoise  $$ x \\equiv \\int \\frac{dr}{f(r)} \\simeq \\int \\frac{dr}{(r-R_{H}) f'(R_{H})} = \\frac{1}{f'(R_{H})} \\log (r-R_{H}) \\equiv \\frac{1}{2k_{H}} \\log (r-R_{H}) \\equiv \\frac{1}{4 \\pi T_{H}} \\log (r-R_{H}), $$  \\noindent locally near the chosen horizon. Here $k_{H}$ is the surface gravity and $T_{H}$ is the Hawking temperature. ",
        "conclusions": ""
    },
    "0403/hep-th0403060_arXiv.txt": {
        "abstract": "We present a Mathematica package for performing algebraic and numerical computations in cosmological models based on supersymmetric theories. The programs allow for (I) evaluation and study of the properties of a scalar potential in a large class of supergravity models with any number of moduli, an arbitrary superpotential, \\Ka\\ potential, and  D-term; (II) numerical solution of a system of scalar and  Friedmann equations for the flat FRW universe with any number of scalar moduli and arbitrary moduli space metric. We are using here a simple set of first order differential equations which we derived in a Hamiltonian framework. Using our programs we present some new results: (I) a shift-symmetric  potential of the inflationary model with a mobile D3 brane in an internal space with stabilized volume; (II) a KKLT-based dark energy model with the acceleration of the universe due to the evolution of the axion partner of the volume modulus. The gzipped package can be downloaded from \\url{http: //www.stanford.edu/~prok/SuperCosmology/} or from  \\url{http: //www.stanford.edu/~rkallosh/SuperCosmology/} ",
        "introduction": "The studies of the cosmological aspects of supergravity and string theory have a long history, going back to the beginning of the 80's.  At present, there is a new wave of interest in the cosmological aspects of string theory. The subject is rather complicated, partly due to the complexity of the analytical study of these theories. For example, to find the F-term part of the effective potential in supergravity, one should specify the expressions for the \\Ka\\  potential $K$ and superpotential $W$. Even with the simplest  $K$ and $W$, the computation of  potential is tedious, especially if there is more than one superfield, and the resulting expression is hard to analyse.    Deriving cosmological consequences from the models based on string theory and supergravity is intricate too. For non-canonical \\Ka\\ potentials (which are the rule rather than the exception) the equations of motion acquire additional velocity-dependent terms whose effects are not easily understood using intuition based on  the simplest  scalar field models. For the case of one field, one can always reduce the theory to the canonical form, but  this method does not work for the description of the simultaneous motion of several different fields, so one should really solve the system of equations keeping the non-canonical kinetic terms throughout.   Our Mathematica-based package ``SuperCosmology'' is intended to simplify the study of supergravity potentials and of the cosmological models based on string theory and supergravity. The programs we describe here were used in papers \\cite{Kallosh:2002gf}-\\cite{Hsu:2003cy}, where only the final results of computations were presented. The purpose of this paper is to present the explanation of the programs and methods we used in computations. Also we present some new cosmological models and use them to demonstrate how our package works.  The SuperCosmology package consists of two parts. Part~I has the following Mathematica nb-files:\\\\   \\noindent {\\tt SuperPotential.nb}, \\\\ {\\tt SuperPotential\\_KKLT.nb}, \\\\ {\\tt SuperPotential\\_fine\\_tune.nb}, \\\\ {\\tt SuperPotential\\_D3.nb}. \\\\  In {\\tt SuperPotential.nb}, {\\tt SuperPotential\\_KKLT.nb}, {\\tt SuperPotential\\_fine\\_tune.nb} one finds examples from \\cite{Hsu:2003cy}, \\cite{Kachru:2003aw} and  \\cite{Kachru:2003sx}, respectively,  of using our  program ``SuperPotential'' in computation of a scalar potential. In  {\\tt SuperPotential\\_D3.nb}, a new example of an inflationary potential with a mobile D3 brane is studied, which is based on the D3/D7 inflationary model investigated before in \\cite{Herdeiro:2001zb}-\\cite{Hsu:2004hi}.  \\noindent In part~II,  we present the program ``FRW'' used  in \\cite{Kallosh:2002gf} for the numerical solution of  the Friedmann equations for a system with any number of scalar fields with geometric kinetic terms of the form ${1\\over 2} G_{ij} (\\phi, \\phi^*) \\partial\\phi^i \\partial \\phi^{j}$ specified by a metric $G_{ij} (\\phi)$ on the scalar manifold. Part~II has the following Mathematica nb-files:\\\\  \\noindent {\\tt FRW\\_N2.nb}, \\\\ {\\tt FRW\\_DarkE.nb}, \\\\ {\\tt FRW\\_LateDarkE.nb}. \\\\  In  {\\tt FRW\\_N2.nb}, we show a dark energy model based on the N=2 supergravity model  \\cite{Fre:2002pd}, \\cite{Kallosh:2002wj} which was also discussed in \\cite{Kallosh:2002gf}.  New results on the dark energy model based on the KKLT model \\cite{Kachru:2003aw}, are presented in {\\tt FRW\\_DarkE.nb} and {\\tt FRW\\_LateDarkE.nb}. The examples in {\\tt FRW\\_DarkE.nb}  describe the situation when the system has not yet reached the dS minimum  (so that the scalars are still moving and $\\Omega_D$ is still increasing). In {\\tt FRW\\_LateDarkE.nb} we study the long term evolution including the time when the dS minimum is reached by the scalars. ",
        "conclusions": "The SuperCosmology Mathematica package presented in this paper proved to be useful in numerous applications, both in our previous work as well as in the new models described in this paper.  The interesting features of the new models are due to the special choice of potentials allowed in supersymmetric theories,  and  the non-canonical geometric kinetic terms.   One of the new  models studied here is the KKLT-based model of dark energy. The kinetic term of the model has an $SL(2,R)$-symmetry typical for string theory and supergravity (see \\cite{Horne:1994mi} for earlier studies of cosmology with $SL(2,R)$-symmetry). One can find a change of variables which will bring one of the scalars to the canonical form, $\\rho= \\sigma +i \\alpha$, $\\sigma=e^{\\sqrt{2/3} \\phi}$, however, the other one  cannot be canonical: $$L_{ kin}= 3{ \\partial \\rho \\partial \\bar \\rho\\over (\\rho+\\bar \\rho)^2}={1\\over 2} [(\\partial\\phi)^2+ {3\\over 2} e^{-2\\sqrt{2/3} \\phi} (\\partial\\alpha)^2]$$ The KKLT non-perturbative potential depends on the volume modulus $\\sigma$ and on the axion $\\alpha$ and has a complicated profile. We have shown the contour plot of this potential as well as some trajectories of scalar fields in Fig. 2.   Consider, for example, the evolution of the model from the point $\\sigma=130$, $\\alpha=15$. This initial point corresponds to the top left (orange) dot on the Fig. \\ref{fig:Fig2}. The solution obtained numerically shows that the volume modulus $\\sigma$ after some initial  increase  stops and waits, while the axion evolves towards $\\alpha = 0$ and oscillates around it. Then $\\sigma$  moves back and eventually both fields get trapped at the minimum of the potential. The unusual behavior of the fields is explained by an interplay between  the potential and  non-canonical kinetic terms. From a more general perspective: in the second order equation for the scalars (\\ref{second}) there is an extra term $\\Gamma^i_{jk}\\dot \\phi^j \\dot \\phi^k$ in addition to the standard friction due to the Hubble parameter and also the contribution of the potential depends on the metric,  $G^{ij} {\\partial V\\over \\partial \\phi^j}$. That makes it hard to guess,  before a numerical solution of equations is found, why some trajectories end up at the minimum of the potential whereas some other trajectories lead to the volume de-compactification.  Our package ``SuperCosmology'' is intended to aid in the study of models related  to string theory as shown in our examples. We hope it will prove to be useful for further investigations of the interface between string theory, supergravity and cosmology.    \\subsection*"
    },
    "0403/astro-ph0403166.txt": {
        "abstract": "{Granato et al. (2004) have elaborated a physically grounded model exploiting the mutual feedback between star-forming spheroidal galaxies and the active nuclei growing in their cores to overcome, in the framework of the hierarchical clustering scenario for galaxy formation, one of the main challenges facing such scenario, i.e. the fact that massive spheroidal galaxies appear to have formed much earlier and faster than predicted by previous hierarchical models, while the formation process was slower for less massive objects. Adopting the choice by Granato et al. (2004) for the parameters governing the history of the star formation, of chemical abundances and of the gas and dust content of galaxies, we are left with only two, rather constrained, but still adjustable, parameters, affecting the time- and mass-dependent SEDs of spheroidal galaxies. After having complemented the model with a simple phenomenological description of evolutionary properties of starburst, normal late--type galaxies, and of AGNs, we have successfully compared the model with a broad variety of observational data, including deep $K$-band, ISOCAM, ISOPHOT, IRAS, SCUBA, radio counts, and the corresponding redshift distributions, as well as the 1--$1000\\,\\mu$m background spectrum. Special predictions have been made for the especially challenging counts and redshift distributions of EROs. We also present detailed predictions for the GOODS and SWIRE surveys with the Spitzer Space Telescope. We find that the GOODS deep survey at $24\\,\\mu$m and the SWIRE surveys at 70 and $160\\,\\mu$m are likely to be severely confusion limited. The GOODS surveys in the IRAC channels (3.6 to $8\\,\\mu$m), reaching flux limits of a few mJy, are expected to resolve most of the background at these wavelengths, to explore the full passive evolution phase of spheroidal galaxies and most of their active star-forming phase, detecting galaxies up to $z\\simeq 4$ and beyond. A substantial number of high $z$ star-forming spheroidal galaxies should also be detected by the $24\\,\\mu$m SWIRE and GOODS surveys, while the 70 and $160\\,\\mu$m will be particularly useful to study the evolution of such galaxies in the range $1 \\lsim z \\lsim 2$. However, starburst galaxies at $z \\lsim 1$--1.5 are expected to be the dominant population in MIPS channels, except, perhaps, at $160\\,\\mu$m. ",
        "introduction": "The standard Lambda Cold Dark Matter ($\\Lambda$CDM) cosmology is a well established framework to understand the hierarchical assembly of dark matter (DM) halos. Indeed, it has been remarkably successful in matching the observed large-scale structure. However the complex evolution of the baryonic matter within the potential wells determined by DM halos is still an open issue, both on theoretical and on observational grounds.  Full simulations of galaxy formation in a cosmological setting are far beyond present day computational possibilities. Thus, it is necessary to introduce at some level rough parametric prescriptions to deal with the physics of baryons, based on sometimes debatable assumptions (e.g.\\ Binney 2004). A class of such models, known as semi-analytic models, has been extensively compared with the available information on galaxy populations at various redshifts (e.g.\\ Lacey et al.\\ 1993; Kauffmann, White \\& Guiderdoni, 1993; Cole et al.\\ 1994; Kauffmann et al.\\ 1999; Somerville \\& Primack 1999; Cole et al.\\ 2000; Granato et al.\\ 2000; Benson et al.\\ 2003).  The general strategy consists in using a subset of observations to calibrate the many model parameters providing a heuristic description of baryonic processes we don't properly understand. Besides encouraging successes, current semi-analytic models have met critical inconsistencies which seems to be deeply linked to the standard recipes and assumptions. These problems are in general related to the properties of elliptical galaxies, such as the color-magnitude and the [$\\alpha$/Fe]-M relations (Cole et al. 2000; Thomas 1999; Thomas et al. 2002), and the statistics of sub-mm and deep IR selected (I- and K-band) samples (Silva 1999; Chapman et al. 2003; Kaviani et al. 2003; Daddi et al. 2004; Kashikawa et al. 2003; Poli et al. 2003; Pozzetti et al. 2003; Somerville et al. 2004; see Cimatti 2003 for a review).  However, the general agreement of a broad variety of observational data with the hierarchical scenario and the fact that the observed number of luminous high-redshift galaxies, while substantially higher than predicted by semi-analytic models, is nevertheless consistent with the number of sufficiently massive dark matter halos, indicates that we may not need alternative scenarios, but just some new ingredients.  Previous work by our group (Granato et al.\\ 2001; Romano et al.\\ 2002; Granato et al.\\ 2004) suggests that a crucial ingredient is the mutual feedback between spheroidal galaxies and active nuclei at their centers.  Granato et al.\\ (2004, henceforth GDS04) presented a detailed physically motivated model for the early co-evolution of the two components, in the framework of the $\\Lambda$CDM cosmology.   \\begin{figure}[tbp] \\centering \\includegraphics[width=9truecm]{c15dstd6090.ps} \\includegraphics[width=9truecm]{c15istd6090.ps} \\caption{Differential (upper panel) and integral (lower panel) $15 \\mu$m counts. The solid line is the sum of contributions from spheroids (dot-dashed line; long dashes single out passively evolving spheroids), spirals (short dashes), starburst galaxies (dotted line) and (type 1 + 2) AGN (filled triangles). Data are from Elbaz et al.\\ (1999), Gruppioni et al.\\ (2002).} \\label{c15std} \\end{figure}   \\begin{figure}[tbp] \\centering \\includegraphics[width=9truecm]{c175std6090.ps} \\includegraphics[width=9truecm]{nz175std6090_223mjy.ps} \\caption{$170\\,\\mu$m counts (upper panel) and redshift distribution of sources with S$_{170}>223\\,$mJy over an area of $3\\,\\hbox{deg}^{-2}$. The dotted, dashed, dot-dashed lines and the filled triangles show the contributions of starburst, spiral, (star-forming) spheroidal galaxies and AGN, respectively. Data in the upper panel are by Dole et al.\\ (2001). In the lower panel, the thin continuous line with asterisks is the sum of the various contributions and the thick continuous histogram shows the data by Rowan-Robinson et al.\\ (2003).} \\label{c175std} \\end{figure}   In this paper, we present a comprehensive comparison of the model with the available data (number counts and redshift distributions) in near-IR (NIR) to sub-mm bands and extensive predictions relevant for surveys such as GOODS and SWIRE, which are being carried out with NASA Spitzer (formerly SIRTF) Observatory. In Sect.~2 we give a short overview of the GDS04 model for spheroidal galaxies, a description of the phenomenological approach adopted to model the evolution of starburst and normal late-type galaxies, and of active galactic nuclei (AGNs). In Sect.~3 we discuss the determination of the two main parameters controlling the time-dependent spectral energy distributions (SEDs) of spheroidal galaxies. In Sect.~4, the model counts and redshift distributions are compared with data from the ISO surveys and follow-up, from the IRAS $60\\,\\mu$m survey, from the SCUBA $850\\,\\mu$m surveys and follow-up, from the radio 1.4 GHz surveys down to sub-mJy flux densities, and from deep K-band surveys and follow-up. In Sect.~5, we present our predictions for Spitzer  GOODS and SWIRE surveys. The main conclusions are summarized in Sect.~6.  We adopt the following cosmological parameters: $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, $H_0=70$ km s$^{-1}$.   \\begin{figure}[tbp] \\centering \\includegraphics[width=9truecm]{nz15std6090_0.1mjydat.ps} \\includegraphics[width=9truecm]{nz15std6090_0.1mjymod.ps} \\caption{Redshift distribution of sources brighter than $0.1\\,$mJy at $15 \\mu$m, within an area of $6\\cdot 10^{-3}\\,\\hbox{deg}^{2}$. In the upper panel the global (spheroids plus spiral and starburst galaxies) redshift distribution predicted by the model (thin continuous line with asterisks) is compared with data by Elbaz et al. (2002, thick solid histogram with error bars) and by Franceschini et al. (2003, three dot-dashed line). In the lower panel we show the various contributions to the global model distribution (again shown by the thin solid histogram with asterisks): starbursts (dots), spirals (dashes), star-forming spheroids (dots-dashes), passively evolving spheroids (long dashes).} \\label{nz15std01} \\end{figure}   \\begin{figure}[tbp] \\centering \\includegraphics[width=9truecm]{nz15std6090_1mjydat.ps} \\includegraphics[width=9truecm]{nz15std6090_1mjymod.ps} \\caption{Redshift distribution of sources brighter than $1\\,$mJy at $15 \\mu$m, within an area of $5.46\\,\\hbox{deg}^{-2}$. In the upper panel the global (spheroids plus spirals, starbursts and AGN) model $z$-distribution (thin continuous line with asterisks) is compared with data by Rowan-Robinson et al. (2003, thick continuous line) and Pozzi et al (2003; three dots-dash, scaled to the same area). In the lower panel, the dotted, short-dashed, dot-dashed lines and the filled triangles show the contributions from starburst, spiral, spheroidal galaxies, and AGN respectively, to the global redshift distribution, represented again by the thin continuous line with asterisks. The long-dashed line singles out the contribution of passively evolving spheroids.} \\label{nz15std1} \\end{figure} ",
        "conclusions": "Granato et al. (2001, 2004) have shown that the mutual feedback between star-forming spheroidal galaxies and the active nuclei growing in their cores can be a key ingredient towards overcoming one of the main challenges facing the hierarchical clustering scenario for galaxy formation, i.e. the fact that the densities of massive high redshift galaxies detected by SCUBA and by deep near-IR surveys are well above the predictions.  However, to take full advantage of the wealth of data on extragalactic sources that are rapidly accumulating in the IR to mm region to test evolutionary models and to assess their parameters, we need to deal with complex and poorly understood processes that determine the time-dependent SEDs of the various populations of galaxies. Indeed, semi-analytic models must rely on a large number of adjustable parameters. We have carried out a detailed comparison of the physically grounded GDS04 model, keeping their choice for the parameters controlling the star-formation history, the chemical enrichment and the evolution of dust and gas content of massive spheroidal galaxies. We are therefore left with only two adjustable parameters, affecting their near-IR to mm SED (see Sect.~\\ref{sect:SED}) computed using the code GRASIL that includes a full treatment of star-light reprocessing by dust; as described in Sect.~\\ref{sect:param}, their values are constrained mostly by $15\\,\\mu$m and $K$-band counts. A simplified phenomenological approach (Sect.~\\ref{sec:other}) has been adopted to deal with the other relevant galaxy populations (normal late-type and starburst galaxies), and the contribution by AGN has been estimated by coupling the cosmological evolution of AGN in the X-ray bands with detailed SEDs (Sect.~\\ref{sec:agn}).  The model predictions have then been tested against a broad variety of observational data, including, in addition to the $15\\,\\mu$m and $K$-band counts, the redshift distributions of sources brighter than 0.1 and 1 mJy at $15\\,\\mu$m, the SCUBA counts at $850\\,\\mu$m, the available (although still scanty) data on the redshift distribution of sources brighter than 5 mJy at $850\\,\\mu$m, the ISOPHOT 90 and $170\\,\\mu$m counts and the corresponding redshift distributions, the IRAS $60\\,\\mu$m counts, the radio 1.4 GHz counts, the ISOCAM $6.7\\,\\mu$m counts and redshift distribution, the redshift distributions of galaxies to the magnitude limits $K= 20$, 23, and 24, and the 1--$1000\\,\\mu$m background spectrum. Specific predictions for the $K$-band counts and redshift distributions of EROs have been worked out and compared with data.  Encouraged by the good agreement of model predictions with all these data sets, we have worked out detailed predictions for the GOODS and SWIRE surveys with the Spitzer Space Telescope. In agreement with previous estimates, we find that the GOODS deep survey at $24\\,\\mu$m and the SWIRE surveys at 70 and $160\\,\\mu$m are likely to be severely confusion limited. The GOODS surveys in the IRAC channels (3.6 to $8\\,\\mu$m), reaching flux limits of a few $\\mu$Jy, are expected to resolve most of the background at these wavelengths, to explore the full passive evolution phase of spheroidal galaxies and most of their active star-forming phase, detecting galaxies up to $z\\simeq 4$ and beyond. A substantial number of high $z$ star-forming spheroidal galaxies should also be detected by the $24\\,\\mu$m SWIRE and GOODS surveys, while the 70 and $160\\,\\mu$m surveys will be particularly useful to study the evolution of such galaxies in the range $1 \\lsim z \\lsim 2$. However, starburst galaxies at $z \\lsim 1$--1.5 are expected to be the dominant population in MIPS channels, except, perhaps, at $160\\,\\mu$m.  We plan to apply our model to make predictions for the forthcoming surveys with Herschel, Planck, LMT and ALMA.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0403/astro-ph0403661_arXiv.txt": {
        "abstract": "The diffuse galactic gamma-ray spectrum measured by the EGRET experiment \\citep{Hunter:1997} are interpreted within a scenario in which cosmic rays (CRs) are injected by three different kind of sources, (i) supernovae (SN) which explode into the interstellar medium (ISM), (ii) Red Supergiants (RSG), and (iii) Wolf-Rayet stars (WR), where the two latter explode into their pre-SN winds \\citep{Biermann:2001iu,Sina:2001}. ",
        "introduction": "\\begin{table}[b] \\label{tab_sources} \\centering \\caption{Interaction spectra for the different type of supernovae. $E_{\\rm knee}$ and $E_{\\rm cut-off}$ are in GeV. $\\gamma_1$ and $\\gamma_2$ are the spectral indices ($\\Phi=\\Phi_0 E^\\gamma$) below and above $E_{\\rm knee}$. $Z$ is the charge of the nucleus.} \\vspace{0.3cm} \\begin{tabular}{ccccc} \\hline SN type&  $E_{\\rm knee}$ & $E_{\\rm cut-off}$ &$\\gamma_1$& $\\gamma_2$    \\\\ \\hline ISM       &                 & $3\\,Z \\,10^{5}$   & -2.75 &                   \\\\ RSG       &                 & $3\\,Z \\, 10^{5}$  & -2.33 &                    \\\\ WR        &  $2\\,Z \\,10^{6}$& $Z \\, 10^{8}$     & -2.88 &      -3.21 \\\\ \\hline \\end{tabular} \\end{table} Conventional models of diffuse galactic gamma-ray production are based on three main processes. Gamma-rays are produced through the decay of $\\pi^0$ as secondary particles of hadronic collisions of CRs with the ISM, such as proton-hydrogen or proton-helium collisions, in bremsstrahlung processes of CR electrons with the ISM, and through inverse Compton scattering of CR electrons with interstellar radiation fields. These models can explain a wide range of observations like the energy spectra below $\\approx 1$~GeV or the integrated flux from the outer parts of our Galaxy. Nonetheless the diffuse galactic gamma-rays observed by the EGRET experiment from the inner Galaxy above $\\approx$~1~GeV exceeds by about 60$\\,$\\% the intensity predicted by these calculations (the measured spectrum is too hard).  In this contribution we shall present first results on the diffuse gamma-ray production expected in the model by \\citet{Biermann:2001iu}. In this model, in addition to the aforementioned processes, interactions of hadrons of the relatively hard CR injection spectrum (see Tab.~\\ref{tab_sources}) with the pre-SN winds are considered as further possible sources of diffuse galactic gamma-rays \\citep{Biermann:2001iu}.  \\begin{figure}[t] \\centering \\includegraphics[width=7.5cm]{0.05new.ps} \\caption{The log of the mean energy of the CR particles versus the number of shock crossings for $V_{sh}=0.01c$. Starting from the bottom for $5^{o}, 25^{o}, 65^{o}, 80^{o}$ and $85^{o}$ respectively. We see the difference in the energy gain of CR particles for the almost perpendicular case compared to smaller shock inclinations. An effect that \\citep{Jokipii:1987} pointed out as well.} \\label{fig_athina1} \\end{figure}  In section~2 a Monte Carlo code is developed to model particle acceleration in non-relativistic near parallel and highly oblique shock configurations, including cross-field diffusion, with application to RSG and WR winds. In section~3 we review the adopted model of diffuse galactic gamma ray production and introduce all necessary input parameters. In section~4 and 5, we present our results, compared to EGRET, CASA-MIA and KASCADE data. In section~6, we summarize our findings and discuss some further improvements in our models. ",
        "conclusions": "The presented model is very promising in explaining the measured diffuse gamma-ray flux. Also it seems to be capable to reproduce the C/B ratio as well as the antiproton flux \\citep{Sina:2001}. It is quite a plausible scenario in which the CR interact mostly in the environment close to their sources. Some stars, like RSG or WR, provide sufficient material, ejected as powerful winds at the end of their lives, before they explode as supernovae. This wind material provides most of the grammage crossed by CR particles seen at Earth. We expect the data in the TeV range from MAGIC and HESS experiments to allow further insight or put constraints on the present model. More detailed studies are planned to calculate the various predictions of the model by \\citet{Biermann:2001iu} and improve the 3D of the interstellar matter, cosmic ray, and radiation field distributions. \\\\  \\vspace{1cm}  \\textit{The present work is being supported by AUGER theory and membership grant 05CU1ERA/3 through DESY/BMBF.}  \\vspace{0.5cm}"
    },
    "0403/astro-ph0403511_arXiv.txt": {
        "abstract": "{ We report on  a calculation of the growth of the mass of supermassive black holes  at galactic centers from dark matter and Eddington - limited baryonic accretion. Assuming that dark matter halos are made of fermions and harbor compact degenerate Fermi balls of masses from $10^{3}M_{\\odot}$ to $10^{6}M_{\\odot}$, we find that dark matter accretion can boost the mass of seed black holes from about $\\sim 5M_{\\odot}$ to $10^{3-4}M_{\\odot}$ black holes, which then grow by Eddington - limited baryonic accretion to supermassive black holes of $10^{6 \\, - \\, 9}M_{\\odot}$. We then  show that the formation of the recently detected supermassive black hole of $3\\times 10^{9}M_{\\odot}$  at a redshift of $z = 6.41$ in the quasar SDSS J114816.64+525150.3 could be understood if the black hole   completely consumes  the degenerate Fermi ball and then grows by Eddington - limited baryonic accretion. In the context of this model we constrain the dark matter particle masses to be within the range from 12 ${\\rm keV/c}^{2}$ to about 450 ${\\rm keV/c}^{2}$. Finally we investigate the black hole  growth dependence on the formation time of the seed BH and on the mass of the seed BH. We find that in order to fit the observed data point of $M_{BH} \\sim 3 \\times 10^{9}M_{\\odot}$ and $z \\sim 6.41$, dark matter accretion cannot start later than about $2 \\times 10^{8}$ years and the seed BH cannot be greater than about $10^{4}M_{\\odot}$. Our results are in full agreement with the WMAP observations that indicate that the first onset of star formation might have occurred at a redshift of $z \\sim 15 \\; - \\; 20$. For other models of dark matter particle masses, corresponding constraints may be derived from the growth of black holes in the center of galaxies. ",
        "introduction": "Over the past few years, the idea of dark matter (DM) and the possible  existence of  supermassive black holes (BH) of masses from $10^{6.5}$ to $10^{9.5}M_{\\odot}$ at the center of galaxies (Macchetto et al. \\cite{macchetto97}; Sch\\\"odel et al. \\cite{schodel02}) have  become deeply rooted within the astrophysical community. The exploration of the relationship between these two intriguing problems of modern astrophysics has been the subject of an increasing number of papers.  It has been established  that the mass of the central BH is tightly correlated with the velocity dispersion $\\sigma$ of its host bulge, where it is found that $M_{BH} \\sim \\sigma^{4-5}$ (Faber et al. \\cite{faber97}; Magorrian et al. \\cite{magorian98}; Ferrarese \\& Merritt \\cite{ferrase00}; Gebhardt et al. \\cite{gebhardt00};  Ferrarese \\cite{ferrase02}; Haering \\& Rix \\cite{rix04}). This tight relation between  the masses of the BHs and the gravitational potential well that hosts them suggests that the formation and evolution of supermassive BHs and the bulge of the parent galaxy may be closely related, e.g.  Wang, Biermann \\& Wandel (\\cite{wang00}). In addition,  the recent discovery of high redshift quasars with $z > 6 $ (Fan et al. \\cite{fan01}) implies that the formation  of supermassive BHs took place over fewer than $10^{9}$years. In spite of the vast and tantalizing work undertaken on BHs, their  genesis and evolution  are not well understood (see Rees \\cite{rees84} for a review). Since the discovery of quasars in the early 1960s, it has been suggested that these objects are powered by accretion of gas onto the supermassive BHs of masses $10^{6}-10^{9}M_{\\odot}$ (Lynden-Bell \\cite{bell69}). Two scenarios have been discussed in modeling the growth of BHs. One is that BHs grow out of a low mass `seed' BH through accretion (Rees \\cite{rees84}), and another one is that BHs grow by merging (Barkana et al. \\cite{barkana01}; Wang, Biermann \\& Wandel \\cite{wang00}; Gopal-Krishna, Biermann \\& Wiita \\cite{bk03,bk04}). In a recent paper, Duschl \\& Strittmatter (\\cite{ds04}) have investigated a model for the formation of supermassive black holes using a combination of merging and accretion mechanisms.  The purpose  of this paper is to study the growth of BHs from dark matter and Eddington-limited baryonic accretion. Also, we would like to obtain some constraints on the DM particle masses. In the past, self gravitating neutrino matter has been suggested  as a model for quasars, with neutrino masses in the range $0.2 {\\rm keV} \\stackrel {\\textstyle <} {\\sim} m_{f}   \\stackrel {\\textstyle <} {\\sim} 0.5 {\\rm MeV}$ (Markov \\cite{markov64}). Later, neutrino matter was suggested to describe DM in clusters of galaxies and galactic halos with masses in the range of $1 {\\rm eV} \\stackrel {\\textstyle <}{\\sim} m_{f}   \\stackrel{\\textstyle <}{\\sim} 25 {\\rm eV}$ (Cowsick \\& McClelland \\cite{cowsik73}; Ruffini \\cite{ruffini80}). More recently, fermion balls (FBs) made of degenerate fermionic matter of  $10 \\ {\\rm keV} \\stackrel {\\textstyle <}{\\sim} m_{f}   \\stackrel{\\textstyle <} {\\sim} 25 \\ {\\rm keV}$ were suggested as an alternative to supermassive BHs in galaxies (Viollier \\cite{viollier94}; Bili\\'c, Munyaneza \\& Viollier \\cite{bmv99}; Tsiklauri \\& Viollier \\cite{tv98}; Munyaneza \\& Viollier \\cite{mv02}). It has been also suggested that if the Galaxy harbors a supermassive BH, then there should be a density spike in which dark matter (DM) falling towards the center could annihilate and the detection of these annihilation signals could be used as a probe for the nature of DM (Bertone, Silk \\& Sigl \\cite{bertone02}; Gondolo \\& Silk \\cite{gondolo99} and  Merritt et al. \\cite{merritt02}). The current belief is that DM particles are bosonic and very massive, i.e. $m_{DM} \\stackrel {\\textstyle >}{\\sim} 1 GeV/c^{2}$. However, the absence of experimental constraints on the weakly interacting massive particles (WIMPs) that probably constitute DM leaves the door open for further investigation of the hidden mass of the Universe.  In this paper, we will assume DM to be of fermionic matter and described by a Fermi - Dirac distribution with an energy  cutoff in phase space (King \\cite{king66}). We will then explore the limits for the DM particle masses in order to reproduce the mass distribution in the Galaxy (Wilkinson \\& Evans \\cite{evans99}) and then study the growth of a seed BH immersed at the center of the DM distribution in galaxies. We use degenerate FBs at the center of DM halos not as replacements  for the BHs but as necessary ingredients to grow the BHs in galactic centers.  The resulting distribution of stars around a massive BH was studied in detail in the  1970s and early 1980s in the context of globular clusters (Hills \\cite{hills75}; Frank \\& Rees \\cite{frank76}; Bahcall \\& Wolf \\cite{bahcal76};  Duncan \\& Shapiro \\cite{duncan82}; Shapiro  \\cite{shapiro85}). Peebles (\\cite{peeble72})  studied the adiabatic growth of a BH in an isothermal sphere and showed that  the BH would alter the matter density to  an adiabatic cusp with $\\rho \\sim r^{-3/2}$. A few years later, Young (\\cite{young80}) constructed numerical models that  confirmed Peebles' results and showed that the BH induces a tangential anisotropy in the velocity dispersion. In this paper, we use the results of previous calculations on  the growth of BHs that are accreting stars (Lightman \\& Shapiro \\cite{lightman78}) to  investigate the growth of a BH that  accretes DM. Here we assume that the physics driving the formation of the power law cusp in the star - star case is the same as in the case of DM particle orbits being perturbed by molecular clouds.  Julian (1967) investigated a similar scenario in which the stellar orbits in our Galaxy were perturbed by molecular clouds to explain the stellar velocity dispersion  dependence on the star's age. Moreover, Duncan \\& Wheeler (\\cite{dw80}) investigated the anisotropy of the velocity dispersions of the stars around the BH in M87.    Given the density distribution in DM halos, we are interested in establishing how a seed BH would grow by accreting DM. Moreover, the comparison of the growth of the BH from accretion of DM and Eddington - limited baryonic matter would give us another piece of information in the debate surrounding the nature of DM. We therefore investigate how a BH seed of $5 \\  M_{\\odot}$ could grow to a $3 \\times 10^{9}M_{\\odot}$ BH as recently detected in  quasar  SDSS J1148+5251 at z=6.41 (Willot, McLure \\& Jarvis \\cite{willot03}). Such a seed BH of a typical mass between 5 and 9 $M_{\\odot}$ could have evolved in BH binaries (Podsiadlowski, Rappaport \\& Hau \\cite{pod03}). Here, we note that the seed BH could  be  an intermediate mass black  hole (IMBH) of $10^{3-4}M_{\\odot}$ that might have formed from collisions in  dense star-forming regions (Portegies Zwart \\& McMillan \\cite{porte02}, Coleman Miller \\cite{miller03} and van der Marel \\cite{van03} for a review). In fact,  Wang \\& Biermann (\\cite{wang98}) have established that a BH would grow exponentially with time accreting baryonic matter as long as the supply lasts. In addition, they were able to reproduce the observed correlation $M_{BH}/M_{sph}$ using standard disk galaxy parameters. Assuming a cusp - like distribution of self-interacting DM (Spergel \\& Steinhardt \\cite{spergel00}), Ostriker (\\cite{ostriker00}) has estimated that BHs could grow to $10^{6}-10^{9}M_{\\odot}$ from DM accretion. For completeness, we note that accretion of DM particles by BHs has recently been studied in Zhao, Haehnelt \\& Rees (\\cite{zhao02}) and Read \\& Gilmore (\\cite{read03}).  Cosmological parameters of $H_{0}=70 \\ {\\rm km \\ s^{-1} Mpc^{-1}}$, $\\Omega_{m}=0.3$ and $\\Omega_{\\Lambda}=0.7$ are assumed throughout this paper. In section~2, we establish the main equations to describe DM in galaxies. We then discuss the  growth of seed BHs from  DM and Eddington - limited baryonic matter accretion in section~3 and  conclude with a discussion in section~4. ",
        "conclusions": "In this paper, we have investigated the growth of  a stellar seed BH immersed at the center of DM halos with degenerate FBs of mass from $\\sim 10^{3}M_{\\odot}$ to $\\sim 10^{6}M_{\\odot}$. Using the Pauli exclusion principle, we have established that the BH accretion rate strongly depends on the mass $m_{f}$ of the fermions as  $\\dot{M}_{BH}\\sim m_{f}^{4}M_{BH}^{2}$ and thus establish for the first time the relationship between  BH growth and fermionic DM. We have shown that in order to fit the DM distribution in the Galaxy with such degenerate cores, the  DM particles should be in the  $ 12 {\\rm keV/c}^{2} \\stackrel {\\textstyle <}{\\sim} m_{f}\\stackrel {\\textstyle <}{\\sim} 450{\\rm keV/c}^{2}$ mass range. We have shown that such DM masses could be  used to fit the distribution of DM in dwarf galaxies. FBs of masses $10^{3 \\; - \\; 6} M_{\\odot}$ could only exist in such galaxies where the density drops off as  $1/r^{2}$  at large distances. Dwarf galaxies as well as cluster of galaxies do not host FBs as their data can only be fitted by a non - degenerate Fermi - Dirac distribution of the King type. We have argued that the merging of dwarf galaxies would lead to the formation of galaxies with degenerate FBs. The wideness of  the fermion mass in the keV range is due to the FB mass range from $10^{3}$ to $10^{6} M_{\\odot}$ that we have adopted in this paper. Our main assumption is the  use of fermions as DM candidates in galaxies. However, if one uses bosons instead of fermions, the range of DM particle masses would of course differ  from the one  obtained in this paper. The range of fermions used in our paper is in conflict with the Lee-Weinberg lower  limit of $\\sim 2 {\\rm GeV/c^{2}}$ on the fermion mass (Lee \\& Weinberg \\cite{weinberg77}). This is due to the fact that the derivation of the Lee-Weinberg limit assumes a  freezout from equilibrium distributions. Our simple  model uses the chemical potential which  allows for non equilibrium distributions and this might modify the Lee-Weinberg argument; this remains to be demonstrated.  If we use heavy fermions with  masses of about  $1 {\\rm GeV/c^{2}}$, then  according to equation (\\ref{eq:ov}) the maximum mass allowed for the degenerate Fermi core would only be of about $1M_{\\odot}$, which is not enough to grow  a stellar mass BH to $10^{3}M_{\\odot}$ in about $10^{8}$ years. On the other hand,  very light fermions i.e. $m_{f}   \\stackrel {\\textstyle <}{\\sim} 12 {\\rm keV/c}^{2}$ would generate very massive degenerate FBs of masses greater than $10^{6}M_{\\odot}$ and the BH would grow to $10^{9}M_{\\odot}$ in a very short time, i.e $z << 6.41$.  The growth of seed BHs from DM accretion is investigated using the quantum cascade mechanism upon which low angular momentum DM particles at the inner Fermi surface are first consumed by the BH and then due to a high degeneracy pressure, higher angular momentum particles are pushed inwards and the process continues until the entire degenerate FB is consumed by the BH. Moreover, molecular clouds have been used  as perturbers of DM particle orbits outside the FB and we have shown that the BH grows faster than the FB. After the BH has consumed  the entire FB, it then grows by Eddington - limited baryonic accretion to higher masses of $\\sim 10^{9}M_{\\odot}$ at redshifts $z \\sim 6.41$. We also point out that  molecular clouds of mass $10^{10}M_{\\odot}$ have also been detected in the host galaxy of the same quasar at a redshift of $z \\sim 6.41$ (Walter et al. \\cite{walter03}).  We have also constrained  the possible starting time of accretion, i.e. the time of BH seed formation. From our analysis, the  mass of a $3\\times 10^{9}M_{\\odot}$ BH  in the quasar SDSS J114816.64+525150.3  at a redshift of $z=6.41$ can be fitted exactly if the accretion process starts at a time of about  $2 \\times 10^{8}$ years, which corresponds to the reionization time.  The seed BH mass is found to be in the range from a few solar masses up to an upper limit of $\\sim 10^{4}M_{\\odot}$. For a seed BH mass  of $ 10^{3\\, - \\, 4}M_{\\odot}$, Eddington baryonic matter accretion would be enough to cause the seed BH to grow into  a supermassive BH of $3\\times 10^{9}M_{\\odot}$ mass. The data point at a redshift of $z=6.41$  can be fitted  by only Eddington baryonic matter accretion with an  efficiency of $\\epsilon \\sim 0.01$.  Our model provides a method to find the DM particles mass. If it is found that there is a clear lower mass cut of $10^{3}$ to $10^{6} M_{\\odot}$ in the distribution of BH masses, then this  mass can be used for the mass of the FB to find the corresponding mass $m_{f}$ of the fermions which will be in the range of 12 keV to 450 keV. If on the other hand the BH mass distribution is a continuous function, then our model of BH growth  with DM will probably be ruled out.  The postulated DM particles in this paper were non - relativistic at the decoupling time and are usually called cold dark matter particles (CDM). The latter have to be neutral, stable or quasi-stable and have to weakly interact with ordinary matter. As mentioned in section 2.1 , these particles could be axions which have been investigated by DAMA/NAI (Bernabei et al. \\cite{bernabei01}). The heavier particles  of mass above $1 {\\rm GeV/c^{2}}$ could also be of the class of DM candidates named WIMPS (Weakly Interactive Massive particles). However, in the standard model of particle physics, CDM cannot be suitable candidates for particles. Thus, a new window beyond the standard model of particle physics has to accommodate these particles for our model of BH growth to work. The DAMA/NAI experiment (see Bernabei et al. \\cite{bernabei05} for a review) which aims at the verification of the presence of DM particles in the Galactic halo, will be able to confirm whether GeV WIMPS or  axions of masses of 12 to 450 keV could exist in nature.  Observations have shown that the masses of supermassive BHs at galactic centers correlate with the masses of the host bulges, i.e. $M_{BH} \\approx  0.002 \\  M_{bulge}$ (Haering \\& Rix \\cite{rix04}). This result is obtained in our model as long as the Eddington limited accretion dominates the final growth of the BH. This happens for BH masses of  more than about $ 10^{5} M_{\\odot}$ (Wang, Biermann \\& Wandel \\cite{wang00}).  Mergers  of dwarf galaxies   as well as the spinning of BHs would play an important role in the growth of the BHs. The consideration of these two effects will be the subject of further investigations. In addition, it would be of great interest to study the growth of BHs from boson DM particles. Although it has been shown that bosons could provide a good fit to the rotation curves in dwarf galaxies, it is not yet clear whether an analogous mechanism could work in galaxies with a $1/r^{2}$ density fall off.  While  it takes only $8.4\\times 10^{8}$ years to grow supermassive BHs in most distant quasars, the Galactic center might have grown to its current mass of $\\sim 10^{6}M_{\\odot}$ with only DM accretion in  a Hubble time. In the following  paper, we will address the growth of the Galactic center BH."
    },
    "0403/astro-ph0403402.txt": {
        "abstract": "Using a high resolution spectrum of the secondary star in the black hole binary \\mbox{A0620$-$00}, we have derived the stellar parameters and  veiling caused by the accretion disk in a consistent way. We have used a $\\chi^{2}$ minimization procedure to explore a grid of 800\\,000 LTE synthetic spectra computed for a plausible range of both stellar and veiling parameters. Adopting the best model parameters found, we have determined atmospheric abundances of Fe, Ca, Ti, Ni and Al. The Fe abundance of the star is $\\mathrm{[Fe/H]}=0.14 \\pm 0.20$. Except for Ca, we found the other elements moderately over-abundant as compared with stars in the solar neighborhood of similar iron content. Taking into account the small orbital separation, the mass transfer rate and the mass of the convection zone of the secondary star, a comparison with element yields in supernova explosion models suggests a possible explosive event with a mass cut comparable to the current mass of the compact object. We have also analyzed the Li abundance, which is unusually high for a star of this spectral type and relatively low mass. ",
        "introduction": "The system \\mbox{A0620$-$00} (V616 Mon) is a low mass X-ray binary (LMXB) discovered as an eruptive X-ray source by {\\it Ariel V} in August 1975 (Elvis et al., 1975). During the outburst, it brighte\\-ned by 6 magnitudes in the optical and after 15 months it had returned to its quiescent magnitude of $m_{V}=18.3$ mag. Spectroscopic observations during quiescence revealed a K5\\,V--K7\\,V stellar spectrum plus an emission line component from an accretion disk surrounding the compact object (Oke 1977; Murdin et al., 1980). Further optical photometric and spectroscopic studies established the orbital period at $P = 0.323$ d and a secondary radial velocity semiamplitude of $K_2 = 457$ {${\\rm km}\\:{\\rm s}^{-1}$} (McClintock \\& Remillard 1986), which implied a mass function of $f(M) = 3.18 \\pm 0.16$ {$M_\\odot$} and thus firm dynamical evidence for a massive compact object---a black hole---in this system. From measurements of the orbital inclination, the compact object mass was estimated at $\\sim 11$ {$M_\\odot\\; $} and the companion star mass at $\\sim 0.7$ {$M_\\odot\\; $} (Shahbaz et al., 1994; Gelino et al., 2001).  Many aspects of the origin and evolution of low mass X-ray binaries (LMXBs) still remain unclear. It is believed that these systems begin as wide binaries with extreme mass ratios and orbital separations of $a \\sim 1000$ {$R_\\odot\\; $} (Portegies Zwart et al., 1997; Kalogera \\& Webbink 1998). After filling its Roche lobe, the massive star engulfs its low mass companion and the latter starts to spiraling in to the massive star's envelope (van den Heuvel \\& Habets 1984; de Kool et al., 1987). A close binary forms if the spiral-in ceases before the low mass companion coalesces with the compact helium core of the primary. The helium core continues its evolution and after SN explosion may turn into a neutron star or a black hole. The system becomes an X-ray binary once the secondary star fills its Roche lobe and begins to transfer matter to the compact object.  The spiral-in process could give rise to a naked He core, identified with Wolf--Rayet stars that have lost their envelopes (Woosley et al., 1995). The high mass-loss rate (Chiosi \\& Maeder 1986; Nugis \\& Lamers 2000) of these stars makes difficult to understand the formation of compact objects as massive as the black hole in \\mbox{A0620$-$00} (Meynet \\& Maeder 2003; Woosley et al., 1993). However, if the hydrogen envelope of the massive star is removed at the end of the He core burning phase (the so-called {\\it Case C} mass transfer, Brown et al., 1999), the mass lost by wind in the short-lived ($\\sim 10^4$ yr) supergiant stage will not be large.  Convection (Langer 1991) and rotation (Maeder \\& Meynet 2000; Heger et al., 2000) influence the structure and evolution of massive stars and subsequently the uncertainties in the treatment of these parameters limit our understanding of the evolution of the progenitors of compact objects. In addition, uncertainties in various aspects of the supernova explosion models affect the predictions of the final remnant mass and the chemical composition of any ejecta captured by the companion. Among the  least known ingredients of these models, we may list:  \\begin{itemize}  \\item The {\\it mass cut}, i.e. the mass above which the matter is expelled at the time of the supernova explosion and below which it remains locked into the compact remnant.  \\item The {\\it amount of fallback} or of the mass which is eventually accreted by the compact core (Woosley \\& Weaver 1995; MacFadyen et al., 2001).  \\item Possible {\\it mixing} during the collapse phase (Herant \\& Woosley 1994; Herant et al., 1994; Kifonidis et al., 2000; Fryer \\& Warren 2002).  \\item The energy of the supernova explosion (Nakamura et al., 2001).  \\item The symmetry of the supernova explosion (MacFadyen \\& Woosley 1999; Maeda et al., 2002).  \\end{itemize}  With the aim of obtaining information on the link between compact objects and their progenitor stars, Israelian et al. (1999) measured element abundances in the secondary star of the black hole binary Nova Scorpii 1994 (GRO J1655$-$40) and found several $\\alpha$-elements (O, Mg, Si, S, and Ti) enriched by a factor of 6--10. Since these elements cannot be produced in a low mass se\\-con\\-da\\-ry star, this was interpreted as evidence of a supernova event that originated the compact object. Taking into account the supernova yields from explosion models of massive stars, the re\\-la\\-ti\\-ve abundances of these elements suggested that the supernova progenitor was in the mass range 25--40 {$M_\\odot$}. Afterwards, these over-abundances were compared with a variety of supernova models, including standard as well as hypernova models (for various helium star masses, explosion energies, and explosion geometries) and a simple model of the evolution of the binary and the pollution of the secondary (Brown et al., 2000; Podsiadlowski et al., 2002). Additional independent evidence for the existence of a supernova event in this system has also been found by Mirabel et al. (2002).  In this paper we analyze the chemical abundances of the secondary star in the LMXB \\mbox{A0620$-$00} with the aim of searching for any evidence of nucleosynthetic products from the progenitor of the compact object. ",
        "conclusions": "We have obtained a high quality spectrum  of the secondary star in \\mbox{A0620$-$00} and derived atmospheric chemical abundances. We have set up a technique that provides a determination of the stellar parameters taking into \\mbox{account} any possible veiling from the accretion disk. We find $T_{\\mathrm{eff}} = 4900 \\pm 150$ K, $\\log g = 4.2 \\pm 0.3$, and a veiling (defined as $F_{\\rm disk}/F_{\\rm cont,star}$) of less than 15 per cent at 5000 {\\AA}  and decreasing towards longer wavelengths. Assuming a mass for the secondary of $M_2 = 0.68 \\pm 0.18$ {$M_\\odot$}, the  estimated  surface gravity leads to a stellar radius of $R_2 = 1.1 \\pm 0.4$ {$R_\\odot\\; $},  consistent with the size of the Roche lobe for the secondary given by Gelino et al. (2001).  The abundances of Fe, Ca, Ti, Al, and Ni are slightly higher than solar. The abundance ratios of each element with respect to Fe were compared with these ratios in late-type main sequence metal-rich stars. Moderate anomalies for Ti, Ni, and especially Al have been found. A comparison with element yields from spherically symmetric supernova explosion models suggests that the secondary star captured part of the ejecta from a supernova that also originated the compact object in \\mbox{A0620$-$00}. The abundances can be explained if a progenitor with a $\\sim 14$ {$M_\\odot\\; $} helium core exploded with a mass cut in the range 11--12.5 {$M_\\odot$}, such that no significant amount of iron could escape from the collapse of the inner layers. Elements such as O, Mg, Si, S, and C, with unavailable transitions in our spectrum, will be studied to confirm this scenario.  The Li abundance in the secondary star in \\mbox{A0620$-$00} is dramatically enhanced in comparison with field  late-type main sequence stars, possibly indicating either that this is a young system ($\\sim0.5$--$2\\times10^8$ yr), or the existence of a Li production-preservation mechanism, such as the $\\alpha$--$\\alpha$ reactions, which have to be tested analyzing the $^7{\\rm Li} / ^6{\\rm Li}$ isotopic ratio using future higher S/N optical spectroscopic observations."
    },
    "0403/astro-ph0403041_arXiv.txt": {
        "abstract": "We present a measurement of the evolution of the stellar mass function in four redshift bins at $0.4 < z < 1.2$, using a sample of more than 5000 $K$-selected galaxies drawn from the MUNICS (Munich Near-Infrared Cluster Survey) dataset.  Our data cover the stellar mass range $10^{10} \\leq M/(\\hMsun) \\leq 10^{12}$. We derive K--band mass--to--light ratios by fitting a grid of composite stellar population models of varying star formation history, age, and dust extinction to BVRIJK photometry. We discuss the evolution of the average mass--to--light ratio as a function of galaxy stellar mass in the K and B bands.  We compare our stellar mass function at $z > 0$ to estimates obtained similarly at $z=0$. We find that the mass--to--light ratios in the K--band decline with redshift. This decline is similar for all stellar masses above $10^{10}\\,\\hMsun$. Lower mass galaxies have lower mass--to--light ratios at all redshifts. The stellar mass function evolves significantly to $z = 1.2$. The total normalization decreases by a factor of $\\sim 2$, the characteristic mass (the knee) shifts toward lower masses, and the bright end therefore steepens with redshift. The amount of number density evolution is a strong function of stellar mass, with more massive systems showing faster evolution than less massive systems. We discuss the total stellar mass density of the universe and compare our results to the values from the literature at both lower and higher redshifts.  We find that the stellar mass density at $z \\sim 1$ is roughly 50\\% of the local value. Our results imply that the mass assembly of galaxies continues well after $z \\sim 1$. Our data favor a scenario in which the growth of the most massive galaxies is dominated by accretion and merging rather than star formation which plays a larger role in the growth of less massive systems. ",
        "introduction": "\\label{sec:introduction}  The stellar mass content in galaxies as a function of redshift is one of the most fundamental observables in the quest to understand galaxy formation and evolution. It provides information on the coupling between the growth of structure through the collapse and subsequent merging of dark matter halos and the physical processes governing the evolution of the baryonic matter. Stellar mass in galaxies grows by star formation within galactic disks, as well as by accretion and the merging of galaxies.  In fact, these two processes are related, because star formation in disks can be triggered or enhanced by tidal interaction in close encounters and by merging events.  This interplay between cosmological structure formation and star formation is believed to govern the mass assembly history of galaxies.  The stellar mass of a galaxy at a given time is difficult to measure, however. While dynamical mass measures are considered to be most reliable, they measure the total mass of an object. The dark matter and gas contributions (which are a function of galaxy type) have to be removed to obtain the stellar mass. These kinds of measurements depend on model assumptions for the dark matter contribution, are observationally very costly, and have therefore only been possible in the local universe so far.  The alternative is to convert the luminosity of a galaxy into a stellar mass by means of a model of its stellar population (derived from photometry or spectroscopy) predicting a mass--to--light ratio (\\ML) in a certain wavelength band.  Near--infrared (NIR) luminosities of galaxies are believed to be well suited for this approach, as the \\ML values vary only by a factor of roughly 2 across a wide range of star formation histories (SFHs; see, e.g., \\citealp{RR93,KC98a,BD01}). This compares to a variation of a factor of $\\sim 10$ in the B--band. In addition, the optical regime is strongly affected by dust extinction which becomes negligible in the K band for the vast majority of galaxies \\citep{TPHSVW98}. By correlating photometric properties of disk galaxies with inclination, \\citet{MGH03} found the edge--on to face--on extinction correction to be 0.1~mag in the K band.  Measuring the stellar masses of galaxies in the local universe by means of modeling their stellar populations has been re--attempted recently using newly available wide--area galaxy surveys. \\citet{Kauffmannetal03a} used spectroscopic data from the Sloan Digital Sky Survey (SDSS), while \\citet{2dF01} and \\citet{BMKW03} combined NIR photometry from the Two Micron All Sky Survey (2MASS) with optical photometry from the 2dF Galaxy Redshift Survey (2dFGRS) and SDSS, respectively, to derive \\ML values and study the stellar mass function (MF) of galaxies.  At $z > 0$, suitable multi--wavelength and redshift data are still sparse. Therefore, the integrated stellar mass density, \\rhosz, has been studied, instead of the stellar MF using the available deep field observations. \\citet{DPFB03} and \\citet{Fontanaetal03} studied \\rhosz\\ in the Hubble Deep Fields (HDFs) over the redshift range $0<z<3$ and compared the global integrated SFH of the universe to their measurement of the total stellar mass density.  These surveys, however, cover only very small fields (several square arcminutes) and therefore cannot measure the evolution of \\rhosz\\ below $z\\sim 1$.  In addition, at higher $z$, cosmic variance, selection biases, and dust extinction are a concern, as the ultraviolet spectral region has redshifted into the observable bands.  At $0<z<1$, the total stellar mass density has been estimated from optical and NIR surveys (\\citealp{BE00}; \\citealp{MUNICS3}, hereafter MUNICS~III; \\citealp{Cohen02}), yet the results concerning the evolution of \\rhosz\\ in this redshift range are still inconclusive.  New wide--angle, NIR selected surveys are promising rapid progress (\\citealp{K20-03,LCIRS02};\\citealp{MUNICS1}, hereafter MUNICS~I) in studying the stellar MF itself instead.  In this work we use data from the wide--field K-band selected field galaxy sample of the MUNICS project (MUNICS~I; \\citealp{MUNICS5}, hereafter MUNICS~V) to study the stellar MF evolution to $z \\sim 1.2$ directly. We derive stellar masses by converting K-band luminosities to mass, by comparing BVRIJK photometry of galaxies to a grid of composite stellar population (CSP) synthesis models based on the simple stellar populations published by \\citet{Maraston98}.  We introduce the sample used in this work in Sect.~\\ref{sec:galaxy-sample} and the method used to derive stellar masses and its uncertainties in Sect.~\\ref{sec:deriving-masses}. We present the results for the stellar MF in Sect.~\\ref{sec:mass-function} and discuss the number density of galaxies with $10^{10} \\leq M/(\\hMsun) \\leq 10^{12}$ in Sect.~\\ref{sec:num-density}. We present our estimate of the total stellar mass density at $0.4 < z < 1.2$ and compare it to those in the literature in Sect.~\\ref{sec:mass-density}. Finally, we summarize and discuss our results in Sect.\\ref{sec:discussion}.  Throughout this work we assume $\\Omega_M = 0.3$, $\\Omega_{\\Lambda} = 0.7$. We write Hubble's constant as $H_0 = 100\\ h\\ \\mathrm{km\\ s^{-1}\\ Mpc^{-1}}$, unless noted otherwise. We denote absolute magnitudes in band $X$ by the symbol $M_X$ and masses by the symbol $M$. ",
        "conclusions": "\\label{sec:discussion}  The results of this study of the evolution of stellar masses of galaxies to $z = 1.2$ can be summarized as follows.  1. The mass--to--light ratios (\\ML s) in the K--band decline with redshift by similar amounts for all stellar masses above $10^{10}\\,\\hMsun$.  Lower mass galaxies have lower \\ML\\ values at all redshifts.  2. The stellar mass function evolves significantly to $z = 1.2$. The total normalization decreases by a factor of $\\sim 2$, the characteristic mass (the knee) shifts towards lower masses and the bright end therefore steepens with redshift.  3. The amount of number density evolution is a strong function of stellar mass. More massive systems show stronger evolution in their number density.  4. In total, roughly half the stellar mass in the present day universe forms at $0 < z < 1$.  At first sight these results seem to disagree with recent measurements of the number density of morphologically selected or color--selected (extremely red objects, or EROs) early--type systems (\\citealp{Imetal96,HSTCFRS399,MCD00,DCetal00a}; but see also \\citealp{TS99,MEABC99}). However, our sample is K--band flux limited and therefore largely morphologically blind. The statement that the number densities of galaxies as a function of their stellar mass content evolves does not contradict a constant number density of elliptical galaxies, as long as their morphologies do not change (on average) as their stellar mass grows. Early type galaxies in clusters, however, show clear evidence for passive evolution, at least up to $z \\sim 1$ (e.g., \\citealp{BZB96,ZB97,SED98,KIDF00c,KIFD01,DF01,DS03}) which implies a constant stellar mass. In the field \\citep{DFKI01,Treuetal01c}, early--type galaxies seem to show larger age scatter, and in a recent study of a sample of strong gravitational lens galaxies, \\citet{VDF03} found that about half of the field elliptical galaxies have younger stellar populations and are best fitted by formation redshifts extending down to $z_f \\sim 1$, a scenario that is fully compatible with our results.  Furthermore, we still find that more massive galaxies have {\\em older} luminosity weighted mean ages (higher \\ML) and that while galaxies with young populations are present in larger numbers at $z \\sim 1$, almost maximally old systems are still present even at the highest redshift we probe. This has also been demonstrated by spectroscopic studies of samples of EROs, most recently by the K20 survey \\citep{K20-02}, using deep optical spectroscopy, and by \\citet{Saraccoetal03}, using spectroscopic follow-up of MUNICS--selected EROs in the NIR.  While half the present day stellar mass seems to form after $z \\sim 1$, it seems that the oldest stellar populations at each epoch are harbored by the most massive objects. The growth of stellar mass at the high mass end of the stellar MF is driven by accretion and merging of objects having similarly old populations, rather than by star formation.  This must be the case because if star formation were to contribute equally to the growth of stellar mass at all masses, the ratio of young to old stars, and hence \\ML, would be the same in all galaxies, which is not the case. The higher \\ML\\ of the more massive galaxies indicates that star formation since $z \\sim 1$ is not as important in these objects as it is at lower masses. A similar result in the local universe was shown by \\citet{KTC94}, and \\citet{BE00} explicitly arrived at the same conclusion by looking at the star formation rate per unit stellar mass in their high-$z$ sample.  To account for their rapid change in number density, however, merging and accretion have to dominate. Indeed, \\citet{CBDP03} find that in the HDF, the merger fraction of massive galaxies increases more rapidly than that of less massive galaxies. Another argument in support of this picture has been put forward by \\citet{CSHC96}. By combining K--band imaging and star formation measurements from OII equivalent widths, they showed that the star formation rate per unit K band luminosity (as a surrogate for stellar mass) is lower in higher mass systems and that this quantity increases with redshift, the increase being slower for higher mass systems (i.e., they formed earlier). In addition, many authors noted a steep increase in the star formation rate at $0 < z < 1$, (e.g., \\citealp{CFRS96,CFRS1297,RRetal97,Floresetal99}).  We wish to point out that the observed trend in density evolution as a function of mass is qualitatively consistent with the expectation from hierarchical galaxy formation models. The most rapid evolution is predicted for the number density of the most massive galaxies, while the number density of less massive galaxies is predicted to evolve less. While older models tended to over--predict this evolution, more recent models seem to yield results more consistent with the redshift distribution of K--band selected galaxies \\citep{Fontanaetal99,LCIRS02}. A detailed comparison of the predicted stellar MFs and colors of galaxies with these observations has yet to be done.  Upcoming large area surveys will advance this field significantly by providing more colors, morphological information, and spectroscopic information. This will allow us to use better CSP model grids, divide the samples up by morphology, and include direct star formation measurements from spectroscopy. This will allow us to determine the relative contributions of star formation and accretion/merging to the mass build--up as a function of cosmic epoch and galaxy properties. Hence these datasets will provide a much more complete picture of the assembly process of galaxies."
    },
    "0403/astro-ph0403277_arXiv.txt": {
        "abstract": "{We present a determination of the distance modulus of the globular cluster NGC 5904 (M 5), obtained by means of the subdwarf main--sequence fitting on the (V,V$-$I) color--magnitude diagram. The subdwarf sample has been selected from the HIPPARCOS catalog in a metallicity range homogeneous with the cluster ([Fe/H] $\\simeq$ $-$1.1). Both the cluster and the subdwarfs have been observed with the same telescope+instrument+filters setup (namely, ESO--NTT equipped with the SUSI2 camera), in order to preserve homogeneity and reduce systematic uncertainties. A set of archival HST data has then been used to obtain a deep and precise ridge line. These have been accurately calibrated in the ground photometric system by using the NTT data and used to fit the cluster distance modulus. By adopting the most commonly accepted values for the reddening, E(B$-$V) = 0.035 and 0.03, we obtain respectively $\\mu_0 = 14.44 \\pm 0.09 \\pm 0.07$ and $\\mu_0 = 14.41 \\pm 0.09 \\pm 0.07$, in agreement with recent determinations. ",
        "introduction": "The ages of the Galactic Globular Clusters (hereafter GGCs) represent one of the milestones on the road of our comprehension of the Universe. In order to derive reliable absolute ages estimates, the most important parameter to determine is the distance \\citep[see][]{renz91}. Errors on the distance propagates through the age determination procedure in such a way that a 0.08 error in the distance modulus (DM) propagates as a 12\\% error in the age. Different ages (younger or older) mean different epoch of formation of the Galaxy. Therefore, it can be easily seen how the uncertainty in GGCs ages estimates significantly affects models of galaxy formation and cosmological evolution.  The recent availability of the \\textit{HIPPARCOS} catalog of trigonometric parallaxes \\citep{perryman97} has greatly improved the knowledge of GGCs distances. A large number of papers on the subject have been published since then \\citep[e.g.][]{reid97,gratton97,pont98,vdb02,gratton03}, most of which present larger distances, and hence younger ages, with respect to previous works \\citep[with the exception, among the cited works, of][who recover the pre-Hipparcos results of 1992]{pont98}.  The method used consists of a fit of the cluster main sequence to the subdwarfs main-sequence of corresponding metallicity, once the reddening has been subtracted. This procedure is, of course, prone to uncertainties, both random and systematic: on the subdwarf side errors on parallax measurements, Lutz-Kelker correction, metallicity; on the cluster side reddening, metallicity, narrowness and accuracy of the mean ridge line. A further effect is represented by the presence of binary stars in the cluster CMD and in the subdwarfs sample. In the first case the determination of the mean ridge line is altered, if a proper fitting procedure is not taken into account. In the second, the magnitude of the field star is different from stars of analogous metallicity. This issue has been widely discussed in, e.g., \\cite{gratton97,pont98,carretta00}, and will be analyzed in detail later. If \\textit{HIPPARCOS} gave the chance to significantly reduce the random errors on parallax measurements, nonetheless other sources of errors still remain. Metallicity, for example, represents a common problem that significantly affects the determination of the distance modulus, because of either the spread in metallicity of the sequence of local subdwarfs or the error on the absolute values. There are two strategies to go around this problem: to obtain a sufficiently large set of subdwarfs in the right metallicity range, or to apply a relation that transforms the colors of the stars of different metallicity. The second introduces a model-dependent relation and should be used with caution. Of course, the first solution is better, especially if the sample of subdwarfs draws a main sequence down to very low masses. In fact, recent HST color-magnitude diagrams of GGCs reach very faint magnitudes, hence small masses, but a similar sequence for subdwarfs is currently not available. \\cite{reid98} fit a distance modulus to NGC 6397 by using a sequence of lower MS subdwarfs, but their work stands on the fact that the sequence of very metal-poor subdwarfs define a reasonably narrow sequence. This is not true, instead, for subdwarfs of higher metallicities (see their Fig. 2). On the other hand, the most widely used set of subdwarfs cover magnitudes ranging from $\\mathrm{M}_{V}\\sim7$ and brighter, leaving uncovered the range between this value and the extreme subdwarfs. This is explained with the relative lack of halo stars near the Sun \\citep{reid98}. One of the possibilities we discuss in order to reduce the uncertainties is to perform a uniform study of a subdwarfs sample and a GGC by using the same telescope, instrumentation and, possibly, observing nights, thus eliminating the systematics due to the relative calibrations among different instrumental setups. For this reason, we observed the target cluster and a sample of subdwarfs in the proper metallicity range matching the cluster's with the same telescope+instrument (NTT+SUSI2) and then used archival data from HST to go faint and calibrated the HST magnitudes into the ground system, obtaining a deep and fully homogenous data set. The selected object is the well known cluster NGC 5904 (M 5), for which a wide set of photometric and spectroscopic studies is available in the literature. Moreover, \\citet[S96]{sandquist96}, \\cite{reid97}, \\cite{gratton97}, \\cite{chaboyer98} and, more recently, \\cite{vdb02} determined the DM of the cluster by using the trigonometric parallaxes of subdwarfs and various sets of models. These works will be often referred to in the following.  NGC 5904 is a typical old halo cluster of intermediate metallicity -{[}Fe/H{]} $\\sim -1.1$ in the scale of \\cite{cg97} (CG), $-$1.4 in the \\cite{zw84} scale (ZW). A recent study of \\cite{ramirez03} gives {[}Fe/H{]} $\\simeq -1.3$. The morphology of the horizontal branch (HB) does not show signs of the second parameter effect and several age determinations are available. The structure of the paper is outlined as follows: Sect. 2 describes the observations and the archival material, Sect. 3 presents the reduction and calibration of the data, in Sect. 4 we present and analyze the subdwarfs sample. In Sect. 5 the cluster color-magnitude diagram (CMD) is presented and discussed, and in Sect. 6 the DM-fit is analyzed. In Sect. 7 a final discussion is summarized. ",
        "conclusions": "This study was aimed at estimating the DM of M 5 through the main-sequence fitting method, using as calibrating sequence a set of local subdwarfs for which the trigonometric parallax from HIPPARCOS was known. In order to achieve our goal we used the same setup telescope+instrument+filters both for the cluster and subdwarfs data, namely the ESO-NTT equipped with the camera SUSI2 and standards ESO V and I filters.  Later, a set of subdwarfs from literature has been added to the original sample in order to improve the sample size, and hence increase the accuracy of the fitting. All the comparisons with recent data from literature show that our data are fully consistent with previous determinations, whose accuracy is known, both for the cluster and the local subdwarfs. A final sample of $\\sim$10 subdwarfs has been used to fit a DM to the cluster. By taking into account the two most commonly accepted values for the reddening (i.e. E(B$-$V)=0.03 and 0.035), we obtained two very close values for the (dereddened) DM: $\\mu_{0}=14.41\\pm0.09\\pm0.07$ and $\\mu_{0}=14.44\\pm0.09\\pm0.07$, respectively, where the first number comes from random error sources, the second from systematic error sources. These values are in good agreement both with \\cite{reid97} ($\\mu_{0}=14.45$ obtained with the subdwarfs having the smallest error on the parallax) and \\cite{carretta00} ($\\mu_{V}=14.59$ corresponding to $\\mu_{0}=14.48$), with a slight preference for the shorter value. \\cite{vdb02}, instead, found $\\mu_{0}=14.38$ with the use of non-diffusive isochrones and $\\mu_{0}=14.30$ with diffusive isochrones. In their work, they use a reddening value $\\mathrm{E(B-V)} = 0.038$ and a metallicity [Fe/H] $= -1.4$ for the cluster, discussing their preference for the ZW scale over the CG  on the argument that the metallicity for M~5 is inconsistent with the metal determination for the field stars, if the value is as high as $-$1.1. Moreover, \\cite{vdb02} choose to not report the field stars to a common metallicity, but this effect acts on the spread of the relation rather than on the absolute value.  In order to try to obtain a further constraint on the DM, we tried to use a set of extreme subdwarfs, using an analogous approach to \\cite{reid98}, that has been found to be unapplicable because both the main sequence locus at the faint magnitude level of the extreme subdwarfs ($\\mathrm{M}_{V} \\sim 10$) is dispersed and the subdwarfs of intermediate metallicity define a sequence dispersed as well. In order to reduce the uncertainty, it would have been desirable to have a set of subdwarfs of fainter magnitudes (i.e. below $\\mathrm{M}_{V} \\sim 7$) to take full advantage of the narrowness of the HST main sequence, but they lack in the HIPPARCOS sample. Moreover, at such faint magnitude level, the lower main sequence stars are very cool and the determination of their metallicity much more uncertain because it is very difficult to take into account all the molecular species in the athmosphere at these low temperatures."
    },
    "0403/astro-ph0403107_arXiv.txt": {
        "abstract": "{  A new application of a fast CCD drift scanning technique that allows us to perform speckle imaging of binary stars is presented. For each observation, an arbitrary number of speckle frames is periodically stored on a computer disk, each with an appropriate exposure time given both atmospheric and instrumental considerations. The CCD charge is shifted towards the serial register and read out sufficiently rapidly to avoid an excessive amount of interframe dead time. Four well-known binary systems (\\object{ADS 755}, \\object{ADS 2616}, \\object{ADS 3711} and \\object{ADS 16836}) are observed in to show the feasibility of the proposed technique.  Bispectral data analysis and power spectrum fitting is carried out for each observation, yielding relative astrometry and photometry. A new approach for self-calibrating this analysis is also presented and validated.  The proposed scheme does not require any additional electronic or optical hardware, so it should allow most small professional observatories and advanced amateurs to enjoy the benefits of diffraction-limited imaging. ",
        "introduction": "\\label{introduction}  Over the last few years, CCDs have been used with increasing frequency for speckle imaging. This work actually started more than a decade ago with the work of Beletic \\& Zadnik (\\cite{zadnik93}), and has more recently been extended by Horch et al. (\\cite{horch97a}, \\cite{horch99}) and Kluckers et al (\\cite{kluckers97}).  There are three main reasons for this change. First, CCDs have dramatically improved in terms of their readout noise and speed characteristics. Second, it has been realized that changing the readout pattern allows one to use large-format CCDs effectively in speckle imaging. Finally, there has been the hope that CCDs would allow diffraction-limited photometry in a way that intensified cameras, such as ICCDs and other microchannel--plate-based devices have not been able to do up to the present.  In this paper, we show that CCD-based speckle imaging can be extended to a relatively inexpensive detection system, namely the Santa Barbara Instruments Group (SBIG) ST-8 camera, which in particular does not have the speed and readout flexibility that other CCDs being used for the same purpose have.  Fast drift scanning has already been shown to be a useful technique for recording lunar occultations at millisecond sampling rate (Fors et al. \\cite{fors01}, \\cite{fors03}).  In the case of speckle observations, a fast drift scanning acquisition scheme is adapted to meet the specific needs of speckle imaging. The data are read out off the chip as fast as possible, but with pauses in readout corresponding to the collection of speckle patterns. In addition, we introduce a method of self-calibration of speckle data in the Fourier plane which makes possible to make reconstructed images without taking data on an unresolved point source. Finally, we make some recommendations which may be valuable for small observatories or advanced amateurs should they wish to carry this work forward. ",
        "conclusions": "\\subsection{New CCDs improvements} \\label{improv}  The performance of the fast drift scanning technique depends on camera readout rate, which directly fixes the minimum dead time attainable. If this is excessive, streaking could handicap the scientific usefulness of the data. In general, the readout rate of a camera will depend on its digitization level and data transfer rate. The digitization rate is fixed by CCD microcontroller design and digitization depth. The data transfer rate is specified by the port architecture being used between the camera and computer.  In the case of the camera used in current paper, its readout rate of 30~kpix~s$^{-1}$ can be considered as moderately low compared with CCDs currently in the market. The fact that it is controlled through parallel port interface fairly limits the final readout rate. This is not surprising, since SBIG-ST8 was designed for general CCD stare imaging purposes, where a long download time was not the main concern.  However, in the last few years, technologies directly related to CCD performance have undergone significant developments. Taking into account only those which apply to full frame CCDs, i.e. the type used most in astronomy, we can consider the following advances:  $\\bullet$ readout noise, which has been continuously dropping in all kind of cameras. Still on the edge of the professional market, the recently available L3Vision technology is able to offer cameras with sub-electron readout noise (e2v Technologies, \\cite{e2v02}).  $\\bullet$ multi output CCD, which increases the frame rate by dividing the data stream to be readout into several channels.  $\\bullet$ new data transfer interfaces, which have noticeably increased the throughput in comparison to parallel port (see Table~\\ref{transfer}). Apogee (\\cite{apogee03}), SBIG (\\cite{sbig02}) and Ethernaude (\\cite{ethernaude01}) constitute recent examples of this improvement, with the application of USB 2.0 and Ethernet interfaces to CCDs in the high-end amateur market. These kinds of initiative can supply readout rates typically 10 to 30 times faster than that offered by our port-parallel camera.  Therefore, the benefits to the fast drift scanning technique from all CCD improvements above are straightforward. On one hand, lower readout noise will increase the SNR of the specklegrams. On the other hand, a faster readout rate will certainly decrease dead time and, as a result, low-level streaking between speckle frames would be effectively reduced.  \\begin{table} \\begin{center} \\caption[]{Data transfer rate for different port architectures.} \\label{transfer} \\begin{tabular}{lc} \\hline \\hline Type & Data transfer rate \\\\ & (Mbit s$^{-1}$) \\\\ \\hline Serial & 0.115 \\\\ Parallel Port EPP/ECP & 0.5-1 \\\\ Firewire & 200 \\\\ USB 2.0 & 480 \\\\ Ethernet & 10/100/1000 \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  \\subsection{Limitations of the self-calibration technique} \\label{limit}  The self-calibration method used here cannot be used in all situations. Indeed, the principal limitation is due to the zenith angle. As the zenith angle increases, the dispersion of the atmosphere elongates the speckles so that the speckle transfer function is no longer radially symmetric, and therefore the point source estimate generated is not an appropriate representation of the speckle transfer function at high zenith angles. This in turn can affect the relative astrometry and photometry derived from such data.  In considering differential photometry, one would expect that this is more sensitive to calibration effects than the astrometric results, since the process of deriving the magnitude difference amounts to estimating the fringe depth in the Fourier plane.  If one uses a symmetric point spread function estimate to deconvolve an asymmetric binary power spectrum, the fringe depth can be severely affected while the fringe spacing and orientation would remain essentially the same.  It is also quite likely that in the case of a faint binary star, it is probably better to use a brighter binary to obtain the one-dimensional trace simply due to signal-to-noise considerations."
    },
    "0403/astro-ph0403331_arXiv.txt": {
        "abstract": "Core radii of globular clusters in the Large and Small Magellanic Clouds show an increasing trend with age.  We propose that this trend is a dynamical effect resulting from the accumulation of massive stars and stellar-mass black holes at the cluster centers.  The black holes are remnants of stars with initial masses exceeding $\\sim 20-25\\msun$; as their orbits decay by dynamical friction, they heat the stellar background and create a core. Using analytical estimates and $N$-body experiments, we show that the sizes of the cores so produced and their growth rates are consistent with what is observed. We propose that this mechanism is responsible for the formation of cores in all globular clusters and possibly in other systems as well. ",
        "introduction": "An enduring problem is the origin of {\\it cores}, regions near the center of a stellar or dark matter system where the density is nearly constant. Resolved cores clearly exist in some stellar systems, e. g. globular clusters \\citep{Harris:96}. In other systems, such as early-type galaxies, cores were long believed to be generic but were later shown to be artifacts of the seeing \\citep{Schweizer:79}. Nevertheless a few elliptical galaxies do exhibit bona-fide cores \\citep{Lauer:02} while many others show a central density that rises only very slowly toward the center \\citep{MF:95}. Density profiles of structures that form from gravitational clustering of density perturbations in an expanding universe are believed to lack cores \\citep{Power:03}, although there is evidence for dark matter cores in the rotation curves of some late-type galaxies (e.g. \\cite{Jimenez:03}).  The existence of a core is usually deemed to require a special explanation.  For instance, galaxy cores may form when binary black holes eject stars via the gravitational slingshot \\citep{Ebisuzaki:91}.  A useful sample for testing theories of core formation is the ensemble of globular clusters (GCs) around the Large and Small Magellanic Clouds (LMC/SMC). These clusters have masses similar to those of Galactic GCs, but many are much younger, with ages that range from $10^6$ -- $10^{10}$ yr. Furthermore ground-based \\citep{Elson:89,Elson:91,Elson:92} and HST \\citep{Mackey:03a,Mackey:03b} observations reveal a clear trend of core radius with age: while young clusters ($\\tau\\ll 10^8$ yr) have core radii consistent with zero, clusters older than $\\sim 10^9$ yr exhibit the full range of core sizes seen in Galactic GCs, $0$ pc $\\lap r_c\\lap 10$ pc (Figure~\\ref{fig:mackey}).  The maximum core radius observed in the LMC/SMC GCs is an increasing function of age and is given roughly by $r_{c}\\approx 2.25{\\rm pc}\\log_{10}\\tau_{\\rm yr} - 14.5$. Attempts to explain the core radius evolution in terms of stellar mass loss \\citep{Elson:91}, a primordial population of binary stars, or time-varying tidal fields \\citep{Wilkinson:03} have met with limited success.  The difficulty is to find a mechanism that can produce substantial changes in the central structure of a GC on time scales as short as a few hundred Myr, while leaving the large-scale structure of the cluster intact.  In this paper, we describe a new mechanism for the formation of GC cores and their evolution with time. Massive stars and their black hole remnants sink to the center of a GC due to dynamical friction against the less massive stars. The energy transferred to the stars during this process, and during the three- and higher-$N$ encounters between the black holes that follow, has the effect of displacing the stars and creating a core. The rate of core growth implied by this model is consistent with the observed dependence of core size on age in the LMC/SMC clusters.  \\begin{inlinefigure} \\begin{center} \\resizebox{0.85\\textwidth}{!}{\\includegraphics{fig_mackey.ps}} \\end{center} \\figcaption{\\label{fig:mackey} Core radius  versus age for LMC and SMC GCs from the samples of Mackey \\& Gilmore (2003a,b). Lines show core radius evolution from the $N$-body simulations with initial cusp slope $\\gamma=1$ and three different scalings to physical units; see text for details.  } \\end{inlinefigure} ",
        "conclusions": "The core formation mechanism proposed here could begin to act even before the most massive stars had evolved into black holes. Evolution times for $20\\msun$ stars are $\\sim 8$ Myr \\citep{Schaller:92}. The earliest phases of core formation, $\\tau\\lap 10^7$ yr, would therefore be driven by the accumulation of massive stars rather than by their remnants. Figure 1 shows possible evidence of core growth on time scales $\\lap 20$ Myr in a few clusters. In this context it is interesting to mention the so-called ``young dense star clusters.'' These clusters have ages $\\aplt 10$\\,Myr, sizes $\\sim 1$\\,pc, and contain $\\aplt 10^5$ stars. Well-known examples are NGC2070 \\citep{Brandl:96}, NGC 3603 \\citep{Vrba:00}, and Westerlund 1 \\citep{Brandl:99}. All of these young clusters have small but distinct cores. The young cluster R136 in 30 Doradus ($\\tau\\approx 5$ Myr) shows clear evidence of mass segregation among the brightest stars \\citep{Brandl:96}.  The mechanism described here is similar to core formation by a binary supermassive black hole in a galactic nucleus via the gravitational slingshot \\citep{Quinlan:96}. The latter process produces cores with masses $\\sim$ a few times the binary mass, assuming that the binary separation decays all the way to the point that coalescence by gravitational wave emission can ensue \\citep{Merritt:03}. If the decay stalls at a larger separation, the displaced mass will be smaller. It is currently uncertain how often the decay would stall \\citep{Milos:03}. Core formation by a {\\it population} of massive remnants also displaces a mass that is a few times the total mass in black holes (Figure 2), and because the smaller black holes are freer to move about, there is less prospect of stalling due to a local depletion of stars. In galactic nuclei, the imprints left on the stellar distribution by the clustering of stellar-mass black holes were probably long ago erased by the growth of the supermassive black hole, by the formation and decay of binary supermassive black holes during galaxy mergers, and by star formation.  A speculative application of these results is to cores formed at the centers of dark matter halos by the clustering of Population III remnants in the early universe \\citep{Volonteri:03}. The latter are believed to contain at least one-half the mass of their stellar progenitors when $m\\gap 250\\msun$ \\citep{Fryer:01}, and the cosmological density of remnants may be similar to that of the supermassive black holes presently observed at the centers of galaxies \\citep{Madau:01}. It follows that the Population III remnants could create cores of appreciable size, {\\it if} a number of them can accumulate in a single halo at a given time, and if the time for their orbits to decay is shorter than the time between halo mergers. Both propositions will require further investigation.  \\bigskip  DM and MH were supported by NSF grant AST02-0631, NASA grant NAG5-9046, and grant HST-AR-09519.01-A from STScI. SPZ was supported by the Royal Netherlands Academy of Sciences (KNAW), the Dutch Organization of Science (NWO), and by the Netherlands Research School for Astronomy (NOVA)."
    },
    "0403/hep-ph0403143_arXiv.txt": {
        "abstract": "\\PRE{\\vspace*{.1in}} Many extensions of the Standard Model have to face the problem of new unsuppressed baryon-number violating interactions. In supersymmetry, the simplest way to solve this problem is to assume R-parity conservation. As a result, the lightest  supersymmetric particle becomes stable and a well-motivated dark matter candidate. In this paper, we show that solving the problem of baryon number violation in non supersymmetric grand unified theories (GUT's) in warped higher-dimensional spacetime can lead to a stable Kaluza--Klein particle. This exotic particle has gauge quantum numbers of a right-handed neutrino, but carries fractional baryon-number and is related to the top quark within the higher-dimensional GUT. A combination of baryon-number and $SU(3)$ color ensures its stability. Its relic density can easily be of the right value for masses in the 10 GeV--few TeV range. An exciting aspect of these models is that the entire parameter space will be tested at near future dark matter direct detection experiments. Other exotic GUT partners of the top quark are also light and can be produced at high energy colliders with distinctive signatures. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403689_arXiv.txt": {
        "abstract": "In an effort to understand unusual power-law spectral slopes observed in some hypercompact HII regions, we consider the radio continuum energy distribution from an ensemble of spherical clumps.  An analytic expression for the free-free emission from a single spherical clump is derived. The radio continuum slope (with $F_\\nu \\propto \\nu^\\alpha$) is governed by the population of clump optical depths $N(\\tau)$, such that (a) at frequencies where all clumps are thick, a continuum slope of $+2$ is found, (b) at frequencies where all clumps are optically thin, a flattened slope of $-0.11$ is found, and (c) at intermediate frequencies, a power-law segment of significant bandwidth with slopes between these two limiting values can result.  For the ensemble distribution, we adopt a power-law distribution $N(\\tau) \\propto \\tau^{-\\gamma}$, and find that significant power-law segments in the SED with slopes from $+2$ to $-0.11$ result only for a relatively restricted range of $\\gamma$ values of 1 to 2.  Further, a greater range of clump optical depths for this distribution leads to a wider bandwidth over which the intermediate power-law segment exists.  The model is applied to the source W49N-B2 with an observed slope of $\\alpha\\approx +0.9$, but that may be turning over to become optically thin around 2 mm. An adequate fit is found in which most clumps are optically thin and there is little ``shadowing'' of rearward clumps by foreground clumps (i.e., the geometrical covering factor $C \\ll 1$).  The primary insight gained from our study is that in the Rayleigh-Jeans limit for the Planck function that applies for the radio band, it is the distribution in optical depth of the clump population that is solely responsible for setting the continuum shape, with variations in the size and temperature of clumps serving to modulate the level of free-free emission. ",
        "introduction": "During the past few years a new class of super-compact HII regions has been recognized via high resolution radio observations.  This class of objects is now referred to as hypercompact (HC) HII regions (Gaume 1994; Kurtz 2000).   HC HII regions are typically about ten times smaller and about a hundred times denser than ultracompact (UC) HII regions (Kurtz and Franco 2000; Kurtz 2002) and have emission measures typically $\\geq 10^8$ pc cm$^{-6}$.  HC HII regions also have rising spectral energy distributions (SEDs) from short cm to mm wavelengths (Hofner \\etal\\ 1996; Kurtz 2002; Sewilo \\etal\\ 2003 and references therein) and often have unusually broad radio recombination lines ($\\geq 40$ km s$^{-1}$) (Gaume 1994; Gaume \\etal\\ 1995; Shepherd \\etal\\ 1995; De Pree, Mehringer, \\& Goss 1997; Sewilo \\etal\\ 2003).  Some HC HII regions (see Sewilo \\etal\\ 2003), like some UC HII regions (see Koo \\etal\\ 1996; Garay \\etal\\ 1993; and Kurtz \\etal\\ 1999) are surrounded by extended low-density ionized halos.  The most likely explanation for the halos is that  HC and UC HII regions are highly clumped, thereby producing porous nebulae to UV photons.  An example of a HC HII region with a power-law SED is G75.78+0.34-H$_2$O, whose spectrum from 6~cm to 7~mm is shown in Figure~\\ref{fig1}.  HC HII regions are found in the vicinity of massive star formation and are often coincident with strong H$_2$O masers.  Some HC HII regions appear to be driving bipolar molecular outflows (Hofner \\etal\\ 1996; Shepherd \\etal\\ 1998).  Sewilo \\etal\\ (2003) have suggested that HC HII regions may represent an evolutionary stage between hot molecular cores and UC HII regions.  During this period, rapid accretion onto the central protostar shuts down, and a circumstellar HII region first becomes large enough to be detected.  A property of at least some HC HII regions that concerns this paper is their power-law SEDs at radio wavelengths.  The SEDs of HC HII regions have spectral indices $\\alpha$ between $+0.3$ to $+1.6$ ($S_{\\nu}\\propto\\nu^{\\alpha}$) with typical values of $\\alpha\\approx +1$ from short cm to mm wavelengths. All HC HII regions seem to have rising power spectra in the range $\\approx 3.6$~cm to $\\leq 3$~mm, but the slopes may differ from source to source.  These spectral indices are especially interesting because they cannot easily be explained by thermal or nonthermal radio continuum emission for constant density nebulae.  They are too shallow for optically thick thermal emission, too steep for optically thin thermal emission, and the wavelength interval is too broad to represent the transition from optically thick to thin thermal emission.  The spectral indices are inconsistent with optically thin synchrotron emission; also, radio recombination lines require that HC HII regions are thermal. In this paper, we explore the possibility that the power-law SEDs of HC HII regions might possibly be the result of hierarchial clumping.  It has been known for some time that power-law density gradients with distance from the ionizing star of HII regions produce radio power-law SEDs (Olnon 1975; Panagia \\& Felli 1975; Wright \\& Barlow 1975). Olnon (1975) showed that the relationship between the slope of the SED ($\\alpha$) and the slope of the power-law density gradient ($\\omega$ where $n_{e}\\propto r^{-\\omega}$) is $\\alpha=(2\\omega-3.1)/(\\omega-0.5)$ for $\\omega>1.5$. Thus, for $\\alpha=0.6$, $\\omega=2.0$; for $\\alpha=+1$, $\\omega=2.6$; and for $\\alpha=+1.5$, $\\omega=4.7$.  The $\\alpha=0.6$, $\\omega=2.0$ parameters correspond to the classical values for a constant velocity wind (Wright \\& Barlow 1975; Panagia \\& Felli 1975). Olnon (1975) points out that when the density gradient has a power-law dependence on radius, the slope of the SED will be determined by the value of $\\omega$ at the radius where the optical depth is about unity. That is, the size of the effective radiating surface depends on both the density gradient and frequency.   Hartmann \\& Cassinelli (1977) showed that a radial outflow whose velocity is a power-law with radius of index $\\beta$ has a density power-law dependence on radius of $\\beta-2$, resulting in a SED power-law ($S_{\\nu}\\propto\\nu^{2/3}$ for a constant velocity wind).  The density power-law index $\\omega$ increases so rapidly with $\\alpha$ that $\\omega$ is improbably large for $\\alpha\\geq +1$ ($\\omega\\geq 2.6$).  In light of the observed density structure of HII regions discussed above, it seems unlikely that real HII regions have such steep and well-behaved density structures with radius, especially in the very early stages of evolution expected for HC HII regions.  We therefore investigate an alternate possible explanation for the observed radio power-law SEDs of HC HII regions, namely hierarchal clumping of nebular gas.  As used here, ``hierarchial clumping'' refers to a region filled with clumps of ionized gas having a range of sizes, temperatures, and optical depths defined by power-law distributions.  There need not be a medium in which all the clumps are embedded although such a structure could be accommodated in our analysis.  A hierarchial clump distribution is not the same as a fractal distribution which posits clumps within clumps within clumps.\tIn a fractal structure the emergent SED is complicated by the fact that every clump is embedded in clumps of larger size, whereas in a hierarchically clumped structure the main complication arises when clumps begin to shadow other clumps, otherwise one does not have to be concerned with radiation transfer through a myriad of larger clumps.  The fact that the interstellar medium (ISM) seems to be clumped on all observed size scales in HII regions, planetary nebulae, and neutral atomic and molecular clouds is a strong motivation for the study presented here.  High resolution Hubble Space Telescope (HST) images of the Orion nebula (O'Dell 2001 and references therein) have revealed an array of small-scale, ionized structures down to the resolution limit of the HST.  The small scale structures (clumps, filaments, knots, etc) are easiest to recognize in Orion because of its proximity to us, but high resolution observations of other HII regions such as M16 also indicate that they are composed of a complex of many clumps of varying sizes (Hester \\etal\\ 1996). It is unlikely that the clumpy structures in Orion and M16 are unique; rather, they probably indicate that such structure is inherent in all HII regions.  Small scale ionized clumps are also seen in planetary nebulae (O'Dell \\etal\\ 2002, 2003). High resolution VLBI observations of Galactic HI absorption toward quasars (Faison \\& Goss 2001; Faison \\etal\\ 1998) show that very small clumps (on the order of a few AU) exist in neutral Galactic HI clouds. Extensive CO observations (e.g., Falgarone \\& Phillips 1996; Elmegreen \\& Falgarone 1996; Falgarone \\etal\\ 1998) have clearly demonstrated the existence of small-scale structures in molecular clouds.  The origin of the clumpy structure in the various phases of the ISM is controversial and may have different explanations in different environments.  For example, turbulence has been suggested by several authors as the origin of structure in molecular and HI clouds (Elmegreen \\& Falgarone 1996; Falgarone \\etal\\ 1998; Lazarian \\& Pogosyan 2000).  However, at least some of the structure in HII regions may be due to hydrodynamical instabilities and/or to pre-existing structure in the natal cloud of an emerging HC HII region.  The reason for clumpy structures in the ISM around massive stars is beyond the scope of this paper.  Here, the observationally established clumpiness of the ISM plus the presence of extended halos around UC and HC HII regions motivates our analysis of the radio free-free spectra of hierarchically clumpy HC HII regions.  In the following section, we derive an analytic expression for the free-free emission from a single spherical clump, and employ the result to consider the radio SEDs from an ensemble of clumps.  In \\S 3, the model is applied to the source W49N-B2. A brief discussion of the model and its results appear in \\S 4, and concluding are remarks given in \\S 5. ",
        "conclusions": "We have used expressions for the radio continuum flux density and optical depth of a single, unresolved, uniform (i.e., temperature and density are constant), spherical clump to calculate the SEDs for an ensemble of many such clumps with a power-law distribution of optical depths. The motivations for this morphology are: (1) the empirical evidence for clumping over a wide range of scale sizes in ionized and neutral atomic and molecular phases of interstellar and circumstellar media; and, (2) the intermediate sloped power-law SEDs observed toward a growing number of HC HII regions.  The primary thrust of this investigation was to determine if a power-law SED with slopes intermediate between the optically thick and thin limits of $\\alpha=+2$ and $-0.1$ (where $F_{\\nu}\\propto\\nu^{\\alpha}$) can be understood as a consequence of emission from a hierarchically clumped medium; and, if so, to investigate under what conditions intermediate sloped SEDs are formed and over how large a frequency interval they may occur.  We have found that it is possible for an ensemble of clumps with a power-law distribution of optical depths to produce power-law SEDs of intermediate slope over a limited bandwidth.  The frequency interval over which an intermediate slope holds is controlled by the range of clump optical depths for an appropriate distribution of clumps N($\\tau$). The greater the range in optical depths, $\\tau_{\\rm max}$ to $\\tau_{\\rm min}$, the broader the bandwidth over which an intermediate slope persists.  The slope of the intermediate SED power-law is determined primarily by the parameter $\\gamma$ which specifies the fraction of the nebula that is filled with optically thick clumps at a given frequency. In our models, intermediate power-law slopes only appear for values of $\\gamma$ between 1 and 2, however, we have not explored all possible values of parameter space.  We find a good fit to the SED of W49N-B2 using the measured radio parameters for this source from De Pree \\etal\\ (2000) for an ensemble of clumps with $\\gamma = 1.5$.  The clump optical depths vary from a maximum of 300 to 0.3 at 7 mm and the mean optical depth at this wavelength is 9.5.  The geometric covering factor is $C \\approx 0.15$. The covering factor is small enough to easily produce a low density halo around the dense ionized core of this HC~HII region, for which some observational evidence exists.  Power-law SEDs of intermediate slopes result from the additive effect of many individual clumps whose turn-over frequencies occur in sequential order over a limited range in frequency.  The primary insight gained from our study of hierarchically clumped nebulae is that the distribution of optical depths of the clump population is solely responsible for determining the continuum shape, with variations in size and temperature of the clumps serving only to modulate the level of the free-free emission.  Logical extensions of this work would be to investigate the effects introduced by non-spherical clump morphologies, non-uniform clumps (i.e., variations of temperature and density within clumps), and clump shadowing."
    },
    "0403/astro-ph0403640_arXiv.txt": {
        "abstract": "We present the model of cosmic rays acceleration at ultrarelativistic subshocks and confront it with the observations of gamma-ray bursts (GRBs) and blazars. We investigate cosmic rays acceleration in shocks with Lorentz factors ($\\gamma$) in the range 3 - 40. We show that fluctuations of the magnetic field downstream of the shock do not play an important role in the acceleration process. Results of numerical simulations for shocks with considered Lorentz factors and perpendicular magnetic field inclinations are presented. We fit the derived particle energy spectral index ($\\sigma$) dependence on fluctuations of the magnetic field upstream and $\\gamma$ with a function. ",
        "introduction": "The acceleration mechanism which operate at ultrarelativistic shock fronts was discovered by Bednarz \\& Ostrowski (1998). The mechanism is different from the diffusive particle acceleration which is assumed to be suppressed at superluminal shock fronts (Bell 1978; Drury 1983). The effect of the magnetic field direction is important for ultrarelativistic shocks because all of them are superluminal. Medvedev \\& Loeb (1999) have shown that the relativistic two-stream instability naturally generate strong magnetic fields which are parallel to the shock front. Therefore ultrarelativistic shocks have to be superluminal even if one could imagine an external magnetic field with the angle between the upstream field and the shock normal smaller than $\\sim 1/\\gamma$. Ultrarelativistic shocks without any mean magnetic fields cannot be considered as real physical phenomena by the same reason.  It have appeared a few papers about a particle acceleration at ultrarelativistic shocks without mean magnetic fields (Kirk et al. 2000; Achterberg et al. 2001; Vietri 2003; Lemoine \\& Pelletier 2003) or with subluminal shock geometry (Ellison \\& Double 2002 - parallel shocks). Their acceleration is similar to the diffusive shock acceleration but it includes anisotropies in the angular distribution upstream of the shock. However, particles in this acceleration are able to return to the shock from downstream to upstream due to large magnetic field fluctuations downstream of the shock or due to subluminal shock geometry as in non-relativistic and mildly relativistic regime. Thus, they have failed to understand the actual ultrarelativistic shock acceleration mechanism because the returning is due to small fluctuations of the magnetic field upstream of the shock (the needed fluctuations decrease when the Lorentz factor of the shock increases) and relativistic effects providing a small change of the particle trajectory in the mean field upstream of the shock to be large as measured downstream. To date, the only numerical calculations performed by Bednarz \\& Ostrowski (1998), Bednarz (2000) and in this paper consider the problem of particle acceleration in ultrarelativistic shocks.  It is known for a long time that relativistic shocks occur in regions of efficient acceleration of leptons. The acceleration to non-thermal distributions is observed at hot spots of extragalactic radio sources, in blazars, GRBs and pulsar wind nebulae. In order to account for the presence of these high energy leptons, some authors have tried to find the acceleration mechanism. Begelman \\& Kirk (1990) proposed shock-drift acceleration at relativistic shocks to operate at hot spots of extragalactic radio sources. In the mechanism, particles are accelerated in a single shock crossing by drifting parallel (or anti-parallel) to the electric field. Afterwards, Hoshino et al. (1992) described a process of shock acceleration of positrons to non-thermal distributions devoted to account for the synchrotron radiation observed in the Crab Nebula and hot spots. In the model, the gyrating reflected protons downstream of the shock dissipate their energy in the form of collectively emitted, left-handed magnetosonic waves which are resonantly absorbed by the positrons immediately behind the ion reflection region. The dynamics of the Weibel instability has recently been simulated by several research groups using 3D plasma kinetic code. These simulations confirm both the generation of the magnetic field and the particle acceleration downstream of the shock. The particle energising in Silva et al. (2003) simulations (electron-positron plasma) is due to pitch angle scattering in the produced magnetic field after saturation. The energy stored in the magnetic field is transfered back to the plasma particles. Simulations of Frederiksen et al. (2003) have yielded the energy transfer from protons to leptons similar to Hoshino et al. (1992). In their description, the scattered protons create a fluctuating electric field which tends to equilibrate the energy between protons and electrons. Nishikawa et al. (2003) results suggest that electrons and ions are accelerated in the direction transverse to the shock normal only.  All the described mechanisms suffer from small energies the particles are able to acquire. The shock-drift acceleration allow for the energy increase of about ten times. Leptons can receive the energy from protons which is about ten times ($\\sim\\gamma$ times in Frederiksen et al. 2003 does not necessarily depend on $\\gamma$) above the thermal energy downstream of the shock. Silva et al. (2003) simulations have led to the generation of high-energy tails in the distribution function, with energies few times above the thermal energy downstream of the shock. In our model presented below, we apply some of these mechanisms to production of seed particles. ",
        "conclusions": ""
    },
    "0403/astro-ph0403195_arXiv.txt": {
        "abstract": "We study the evolution of supernova remnants in the circumstellar medium formed by mass loss from the progenitor star. The properties of this interaction are investigated, and the specific case of a 35 $\\msun$ star is studied in detail. The evolution of the SN shock wave in this case may have a bearing on other SNRs evolving in wind-blown bubbles, especially SN 1987A. ",
        "introduction": "Type II Supernovae are the remnants of massive stars (M $>$ 8 M$_{\\odot}$). As these stars evolve along the main sequence, they lose a considerable amount of mass, mainly in the form of stellar winds. The properties of this mass loss may vary considerably among different evolutionary stages. The net result of the expelled mass is the formation of  circumstellar wind-blown cavities, or bubbles, around the star, bordered by a dense shell. When the star ends its life as a supernova, the resulting shock wave will interact with this circumstellar bubble rather than with the interstellar medium. The evolution of the shock wave, and that of the resulting supernova remnant (SNR), will be different from that in a constant density ambient medium.  In this work we study the evolution of supernova remnants in circumstellar wind-blown bubbles. The evolution depends primarily on a single parameter, the ratio of the mass of the shell to that of the ejected material. Various values of this parameter are explored. We then focus on a specific simulation of the medium around a 35 $\\msun$ star, and show how pressure variations within the bubble can cause the shock wave to be corrugated. Different parts of the shock wave collide with the dense shell at different times. Such a situation is reminiscent of the evolution of the shock wave around SN 1987A. ",
        "conclusions": ""
    },
    "0403/astro-ph0403476_arXiv.txt": {
        "abstract": "{ Using  HST/WFPC2 imaging in  F606W (or F450W)  and F814W filters,   we  obtained   the  color   maps in observed frame  for  36   distant (0.4$\\,<\\,z\\,<\\,$1.2) luminous infrared galaxies (LIRGs, L$_{\\rm IR}(8-1000\\,\\mu$m) $\\geq$\\,$10^{11}$\\,L$_\\odot$), with average star  formation  rates  of $\\sim$100\\,M$_\\odot$\\,yr$^{-1}$. Stars  and compact  sources are  taken as  references to  align images after correction of geometric distortion. { This} leads  to an alignment accuracy  of  0.15\\,pixel,  which  is  a prerequisite for studying the detailed color properties of galaxies with complex morphologies. A new method is  developed to  quantify the reliability  of each pixel  in the color map without any bias against very red or blue color regions.  Based on analyses of  two-dimensional structure and spatially resolved color  distribution, we carried out morphological classification for LIRGs. About 36\\% of  the LIRGs were classified as disk galaxies and  22\\% as  irregulars.  Only 6  (17\\%) systems  are obvious ongoing  major mergers. An upper  limit of  58\\%   was found for the  fraction of  mergers in  LIRGs with all  the possible merging/interacting systems included.  Strikingly,  the  fraction  of  compact sources  is as high as 25\\%, similar to  that found  in optically selected samples.  From their K band luminosities, LIRGs are relatively  massive  systems, with  an average stellar  mass of about 1.1$\\times$10$^{11}$\\,M$_\\odot$. They  are related to the  formation of massive and large disks,  from their morphologies and also  from the fact that they  represent a significant  fraction of  distant disks  selected by their sizes.  If sustained at  such large rates, their  star formation can double their stellar masses in less than 1 Gyr.  The compact LIRGs show blue cores,  which could be associated with  the formation of the central  region  of  these  galaxies.   We find  that  all  LIRGs  are distributed along a sequence which relate their central color to their concentration  index. This  sequence links  compact objects  with blue central  color to  extended ones  with relatively  red  central color, which are  closer to  the local disks.  We suggest that there are many massive disks which are  still forming a large fraction of their stellar mass since $z$\\,=\\,1. For most of them, their central parts (bulge?) were formed prior to the  formation of  their disks. ",
        "introduction": "The evolution of the cosmic star formation density (CSFD) exhibits the history of  the stellar mass  assembly averaged over all  galaxies.  A sharp decline of the CSFD  since $z\\,\\sim$\\,1 has been found, whereas large uncertainties still remain at higher redshifts, particularly due to the uncertainties and biases  regarding dust extinction (e.g. Madau et al.~\\cite{Madau};  Hammer et al.~\\cite{Hammer97}).   Investigations of global stellar mass  density as a function of  redshift indicate that more than one quarter, probably  up to half of the present day stars were formed  since  $z\\,\\sim$\\,1  (Dickinson  et al.  \\cite{Dickinson}  and references therein).  This is in  agreement with an integration of the CSFD if the latter accounts for all the light re-radiated at IR wavelengths (Flores  et  al.~\\cite{Flores}).   Hence the  star-forming  activities since $z\\,\\sim$\\,1  still play an  important role in the  formation of the galaxy Hubble sequence seen in the local universe.  {\\it Hubble Space Telescope} (HST) observations show that  the merger rate increases significantly at $z\\,\\sim$\\,1,  compared  with   that  in   the  local   universe  (Le F$\\grave{\\rm e}$vre     et      al.     \\cite{Fevre};     Conselice     et al. \\cite{Conselice}).  Such events were claimed to be related to dwarf galaxies while massive systems have  almost formed before  redshift 1 (Brinchmann     \\&      Ellis~\\cite{BrinchmannEllis};     Lilly     et al. \\cite{Lilly98}; Schade et  al. \\cite{Schade}).  However, with {\\it Infrared  Space  Observatory} (ISO)  mid-infrared  imaging, Flores  et al.  (\\cite{Flores})  inferred that  a  substantial  fraction of  star formation since $z\\,\\sim$\\,1 is associated with the LIRGs. These objects are luminous star-forming galaxies  at intermediate  redshifts ($z\\,\\sim\\,$0.5  to  1), different from the faint  blue galaxy population (Genzel  \\& Cesarsky \\cite{Genzel}; Franceschini  et al.~\\cite{Franceschini03}).   It  is widely  accepted that merger/interaction is very  efficient in pushing gas into nuclear region and triggering violent star formation. Therefore LIRGs are suspected to  be merging systems  and the evolution  of these  galaxies is linked to the decline of  the merger rate (Elbaz et al. \\cite{Elbaz}). Although HST imaging  showed that most of  the LIRGs are  luminous disk/interacting galaxies  (Flores et al.~\\cite{Flores}),  systematic investigation  of their  properties is still required  to understand their  formation and evolution, as well as link them to the counterparts in the local universe.  Morphological classification  is essential to revealing  the nature of the distant LIRGs. However, at  high redshifts, it becomes difficult to classify galaxy  morphology securely because, the images  of the high-z  galaxies suffer from  reduced resolution, band-shifting  and cosmological  surface  brightness dimming  effects, compared with the local objects.  With  HST {\\it Wide Field Planet Camera 2} (WFPC2)  high resolution imaging in  two or more  bands, spatially resolved  color distribution can be used to investigate the distribution of the stellar population, which is  complementary  to  addressing   the  appearance  in  single  band. Furthermore, the star-forming regions  and dusty regions can stand out in the color map. This is  very important to the study of LIRGs, in which these regions  are  expected to be numerous.  Canada-France Redshift Survey (CFRS) fields are among the most studied fields at various wavelengths.   Two CFRS fields 0300+00 and 1415+52 had  been   observed  deeply  by  ISOCAM  at   15\\,$\\mu$m  (Flores  et al.  ~\\cite{Flores}; 2004 in  preparation) and  by HST  (Brinchmann et al.~\\cite{Brinchmann}).   Aimed   at  performing  detailed   analyses  of morphology,  photometry  and  color  distribution for  distant  LIRGs, additional HST images through blue  and red filters have been taken to complement the color information for  the two CFRS fields (PI: Hammer, Prop. 9149).  In this work,  we present the preliminary results of the color distribution of  the distant LIRGs.  We correct additional effects in HST images and  recenter them accurately, which allow us to access the color maps of complex galaxies. We also implement a method to quantify the signal-to-noise  (S/N) ratio of  the color image in order  to give a reasonable cut for the  target area in color maps.  This paper is organized as  follows.  Sect. 2 describes the HST imaging observations and  the archive data we  adopt. In Sect.  3, we describe the various effects which have  to be corrected for aligning images in different WFPC2 filters.  In Sect. 4, we describe a method we use to generate the color  maps.   In Sect.  5, we  summarize  the morphological properties of the distant LIRGs. The results we obtained of the LIRGs are discussed in Sect. 6.  Brief conclusions are given in Sect.      7.       Throughout      this      paper      we      adopt H$_0$\\,=\\,70\\,km\\,s$^{-1}$\\,Mpc$^{-1}$,  $\\Omega_{\\rm  M}$\\,=\\,0.3 and $\\Omega_\\Lambda$\\,=\\,0.7. Unless specified, we exclude PC chip and the unit  of pixel  refers  to that  in  WF chips.   The bands  B$_{450}$, V$_{606}$ and  I$_{814}$ refer to  HST filters F450W, F606W  and F814W, respectively.  Vega system is adopted for our photometry.  \\begin{table*} \\centering \\caption[]{HST imaging with two bands observations in CFRS fields 0300+00 and 1415+52} \\label{hstlist} \\begin{tabular}{cccccccccc} \\hline \\noalign{\\smallskip} Field &  BlueFilter & Total Exp. & N$^{\\mathrm{a}}$ & Dither$^{\\mathrm{b}}$ & RedFilter & Total Exp. & N$^{\\mathrm{a}}$ & Dither$^{\\mathrm{b}}$ & Prop.ID$^{\\mathrm{c}}$\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 030226+001348 & F606W & 6400 & 5 & 20& F814W & 6000 & 5 & 20& 9149 \\\\ 030227+000704 & F450W & 7000 & 5 & 20& F814W & 6700 & 5 & 20& 6556,5996 \\\\ 030233+001255 & F450W & 6600 & 6 & 0 & F814W & 6400 & 6 & 0 & 5449 \\\\ 030237+001414 & F606W & 6400 & 5 & 20& F814W & 6400 & 5 & 20& 9149 \\\\ 030240+000940 & F606W & 6400 & 5 & 20& F814W & 7000 & 5 & 12.5& 9149,8162 \\\\ 030243+001324 & F450W & 6600 & 6 & 0 & F814W & 6400 & 6 & 0 & 5449 \\\\ 030250+001000 & F606W & 6400 & 5 & 20& F814W & 7000 & 5 & 12.5& 9149,8162 \\\\ 141743+523025 & F450W & 7800 & 6 & 0 & F814W & 7400 & 6 & 0 & 5449 \\\\ 141803+522755 & F606W & 6400 & 5 & 20& F814W & 6800 & 5 & 20& 9149 \\\\ 141809+523015 & F450W & 7800 & 6 & 0 & F814W & 7400 & 6 & 0 & 5449 \\\\ 141724+522512 & F606W & 2800 & 4 & 0 & F814W & 4400  & 4 & 0 & 5090 \\\\ 141731+522622 & F606W & 2800 & 4 & 0 & F814W & 4400  & 4 & 0 & 5090 \\\\ 141737+522731 & F606W & 2800 & 4 & 0 & F814W & 4400  & 4 & 0 & 5090 \\\\ 141750+522951 & F606W & 2800 & 4 & 0 & F814W & 4400  & 4 & 0 & 5090 \\\\ 141743+522841 & F606W &24400 &12 & 0 & F814W &25200  &12 & 0 & 5109 \\\\ 141757+523101 & F606W & 2800 & 4 & 0 & F814W & 4400  & 4 & 0 & 5090 \\\\ 141803+523211 & F606W & 2800 & 4 & 0 & F814W & 4400  & 4 & 0 & 5090 \\\\ \\noalign{\\smallskip} \\hline \\end{tabular}  \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] Number of exposures. \\item[$^{\\mathrm{b}}$] Dither offset among the consecutive exposures, aimed to remove cosmic-rays and hot/bad pixels.  Here the largest offset in unit of pixel along x axis is present. \\item[$^{\\mathrm{c}}$] HST Proposal ID. One ID means that both the blue and the red observations were carried out in the same proposal. Two IDs refer to the proposals that the blue and the red observations were observed  respectively. \\end{list} \\end{table*} ",
        "conclusions": "Specific efforts were made to  obtain the color maps for galaxies with complex morphology,  including the accurate alignment of  the blue and red  band images  and a  new  method to  quantitatively determine  the reliability  of each pixel  in the color  map. These  efforts allow  us to access the spatially resolved color distribution of the distant LIRGs, which    often     have    complex    morphologies,     relating    to interactions/mergers.   In   two  10$\\arcmin$$\\times$10$\\arcmin$  CFRS fields 0300+00 and 1415+52, HST  WFPC2 imaging in F606W (or F450W) and F814W filters is available  for 87 square arcminute area. From these fields , we select a representative  sample of 36 distant ($0.4\\,<\\,z\\,<\\,1.2$) LIRGs detected  in  deep  ISOCAM  observations.   Two-dimensional  structure analysis  is  carried  out   using  GIM2D  software.   With  structure parameters   and   color   distribution,   a   careful   morphological classification was performed for the distant LIRGs.  We find that about 36\\% LIRGs are spiral galaxies and about 25\\% LIRGs show compact morphology.  About 22\\%  LIRGs are  classified as  irregular galaxies, showing complex and  clumpy structures. Among 36 LIRGs,  only 6 (17\\%) of  them   were  undergoing  a  major  merger   episode,  revealed  by distinctive close  galaxy pair with distorted  morphology and apparent tidal tails. The  fraction of mergers could reach  58\\% if all of the possible post-mergers/pre-mergers are included.  Inspection  of  their stellar  masses  derived  from  K band  absolute magnitude  evidences  that  LIRGs  are  massive  systems.   The  LIRGs classified as  disk galaxies  indeed belong to  the large  disk galaxy population, and  become a significant fraction of  large distant disks selected by their sizes.  We find  that LIRGs  are distributed along  a sequence in  the central color versus compactness diagram. The sequence links the compact LIRGs with  relatively blue  central color  to that  of extended  LIRGs with central  color and  compactness close  to  those of  the local  normal galaxies.  The compact LIRGs  showing blue central color are suggested to  be  the  systems  forming  their bulges,  in  agreement  with  the suggestion of Hammer et al. (\\cite{Hammer01}). We argue that the sequence suggests that distant compact LIRGs  would eventually evolve  into the spiral galaxies in the local universe."
    },
    "0403/astro-ph0403706_arXiv.txt": {
        "abstract": "The \\al radionuclide can be detected through its decay emission line at 1.809 MeV, as was first observed by Mahoney et al. (1982). Since then, COMPTEL on board of the CGRO satellite, performed a sky survey in this energy range, and provided maps of the \\al distribution in the Galaxy. These results revealed that the main contributors to the synthesis of \\al are most likely the massive stars, which contribute through their winds (Wolf-Rayet stars) and through their supernova explosion.\\\\ Comparison between these observations (in particular observations in localized regions such as the Vela region and the Cygnus region) and the models available at that moment, showed however the need for improvements from both theoretical and observational points of view, in order to improve our understanding of the \\al galactic distribution as well as that of its synthesis.\\\\ With the launch of the INTEGRAL satellite in October 2002, the observational part will hopefully be improved, and the construction of better resolution maps at 1.809 MeV is one of the main aims of the mission. From a theoretical point of view, we need the most up-to-date predictions in order to be able to interpret the forthcoming data.\\\\ In this paper, we address this latter part, and present new results for \\al production by rotating Wolf-Rayet stars and their contribution to the total amount observed in the Galaxy. ",
        "introduction": "\\al is a radioactive nuclide that can be produced by hydrostatic nucleosynthesis in H burning regions, by explosive nucleosynthesis and by spallation. Its half--life time in its ground state is of $7.2 \\times 10^5$ yr, and it is thus a good tracer of recent nucleosynthetic events in the Galaxy.\\\\ The maps obtained with COMPTEL (Diehl et al. 1995, Kn\\\"odlseder et al. 1999, Pl\\\"uschke et al. 2001) allowed to pin down the main contributors to the diffuse emission observed mainly in the galactic plane : massive stars (Prantzos \\& Diehl 1996, Kn\\\"odlseder 1999). These objects end their lives with a supernova explosion, during which \\al can be synthesized and will eventually be expelled (Heger et al. 2003). The more massive (and shorter lived) ones will also contribute during quiescent evolutionary phases through their winds. ",
        "conclusions": "We have presented new results concerning \\al production by WR stars from a new grid of models including rotation and updated physics, in particular the most recent prescriptions for mass loss rates.\\\\ Taking rotation into account globally leads to an enhancement of wind ejected mass of \\al by very massive stars.\\\\ Convolved with appropriate IMF and star formation rate indicator, these yields lead to a total mass of about 1.3 \\msun~of \\al originating from WR stellar winds. This value is in agreement with the conclusions drawn by Prantzos (2004) from recent measurements of the line flux ratio $^{60}{\\rm Fe}/^{26}{\\rm Al}$."
    },
    "0403/astro-ph0403256_arXiv.txt": {
        "abstract": "We use the statistics of strong gravitational lensing based on the Cosmic Lens All-Sky Survey (CLASS) data to constrain cosmological parameters in a spatially-flat, inverse power-law potential energy density, scalar-field dark energy cosmological model. The lensing-based constraints are consistent with, but weaker than, those derived from Type~Ia supernova redshift-magnitude data, and mildly favor the Einstein cosmological constant limit of this dark energy model. ",
        "introduction": "Recent cosmological measurements strengthen the evidence from Type~Ia supernova redshift-magnitude measurements (Riess et al.~1998; Perlmutter et al.~1999) that the energy density of the current universe is dominated by Einstein's cosmological constant $\\Lambda$, or by a dark energy term in the cosmic stress-energy tensor that only varies slowly with time and space and so acts like $\\Lambda$. These measurements include: (1) more recent Type~Ia supernova redshift-magnitude measurements (see, e.g., Knop et al.~2003; Barris et al.~2004); (2) the space-based Wilkinson Microwave Anisotropy Probe (WMAP) measurement of the cosmic microwave background (CMB) anisotropy, with some input from other measurements (see, e.g., Page et al.~2003; Spergel et al.~2003; Scranton et al.\\ 2003); and (3) other measurements of CMB anisotropy, which indicate the universe is close to spatially flat (see, e.g., Podariu et al.~2001b; Durrer, Novosyadlyj, \\& Apunevych 2003; Melchiorri \\& \\\"Odman 2003), in combination with the continuing strong evidence for low non-relativistic matter density (Chen \\& Ratra 2003b and references therein). See Peebles \\& Ratra (2003), Padmanabhan (2003), Bernardeau (2003), Steinhardt (2003), and Carroll (2004) for reviews of the current state of affairs.\\footnote{ Specific dark energy models and observational measurements are considered in Munshi, Porciani, \\& Wang (2003), Barreiro et al.~(2003), Mainini et al.~(2003), Lima, Cunha, \\& Alcaniz (2003), Silva \\& Bertolami (2003), Amendola et al.~(2003), Linder \\& Jenkins (2003), Makler, Oliveira, \\& Waga (2003), Bean \\& Dor\\'e (2003), {\\L}okas, Bode, \\& Hoffman (2003), Alam et al.~(2003), Choudhury \\& Padmanabhan (2003), Zhu \\& Fujimoto (2004), and Macci\\'o (2004), from which the earlier literature may be accessed.}  While Einstein's $\\Lambda$ was the first example of dark energy, nowadays much attention is focused on scalar field models in which the energy density slowly decreases with time and so behaves like a time-variable $\\Lambda$ (see, e.g., Peebles 1984; Peebles \\& Ratra 1988, 2003; Padmanabhan 2003; Steinhardt 2003; Carroll 2004). A simple scalar field dark energy model has scalar field $(\\phi)$ potential energy density $V(\\phi) \\propto \\phi^{-\\alpha}$ at low redshift, with $\\alpha > 0$ (see, e.g., Peebles \\& Ratra 1988; Ratra \\& Peebles 1988). Podariu \\& Ratra (2000), Waga \\& Frieman (2000), and Gott et al.~(2001) examine constraints on this model using Type~Ia supernova redshift-magnitude data. They find that a broad range of $\\alpha$ is consistent with the supernova data.\\footnote{ The proposed SNAP space mission (see http://snap.lbl.gov/, and Schubnell 2003 and Annis et al.~2003) will provide significantly tighter constraints on such models (Podariu, Nugent, \\& Ratra 2001a; Ericksson \\& Amanullah 2002; Caresia, Matarrese, \\& Moscardini 2003; Wang \\& Mukherjee 2003, and references therein). Mukherjee et al.~(2003a, 2003b), Spergel et al.~(2003), Caldwell \\& Doran (2003), Weller \\& Lewis (2003), Giovi, Baccigalupi, \\& Perrotta (2003), and references therein, discuss constraints on scalar field and related dark energy models from CMB anisotropy measurements; upcoming WMAP and other CMB data will improve these constraints.}  It is important that these dark energy models be tested by other independent methods. The redshift--angular-size test is one option. Indications from current data, while not as compelling as those discussed above, are consistent with a significant dark energy density at low redshift (see, e.g., Daly \\& Guerra 2002; Zhu \\& Fujimoto 2002; Chen \\& Ratra 2003a; Podariu et al.~2003; Jain, Dev, \\& Alcaniz 2003; Jackson 2003). Future higher-quality data should turn this into a much more precise cosmological test. The redshift-counts test also appears to be on the verge of becoming a very promising test (see, e.g., Newman \\& Davis 2000; Huterer \\& Turner 2001; Podariu \\& Ratra 2001; Levine, Schulz, \\& White 2002). Statistical analyses of strong gravitational lensing can be used to provide constraints on cosmological parameters. Fukugita, Futamase, \\& Kasai (1990) and Turner (1990) note that the rate of gravitational lensing increases rapidly with increasing $\\Lambda$. Ratra \\& Quillen (1992) and Waga \\& Frieman (2000) study gravitational lensing in the inverse power-law potential scalar field dark energy model.  The recently completed Cosmic Lens All-Sky Survey (CLASS) is the largest uniform survey for strong lensing (Myers et al.\\ 2003; Browne et al.\\ 2003). The survey has discovered 22 cases of multiple-imaging (that are induced by galaxy-scale lens potentials) out of $\\sim 16,500$ extragalactic radio sources. A subsample of 8958 sources containing 13 multiply-imaged sources satisfy well-defined observational selection criteria and is referred to as the CLASS statistical sample (Browne et al.\\ 2003). The CLASS statistical sample has been used to constrain cosmological parameters (see, e.g., Chae et al.\\ 2002; Chae 2003; Kuhlen, Keeton, \\& Madau 2004; Mitchell et al.\\ 2004) as well as to constrain global properties of galaxies (Chae 2003; Davis, Huterer, \\& Krauss 2003) and galaxy evolution (Chae \\& Mao 2003). The lensing-based constraints on cosmological parameters are consistent with those based on Type~Ia supernova magnitude-redshift data but have larger statistical errors.  In this work we use the CLASS statistical sample to constrain the inverse power-law potential scalar-field dark energy model (Peebles \\& Ratra 1988). In linear perturbation theory, a scalar field is mathematically equivalent to a fluid with time-dependent equation of state parameter $w = p/\\rho$ and speed of sound squared $c_s^2 = \\dot p/\\dot\\rho$, where $p$ and $\\rho$ are the pressure and energy density, and the dot denotes a time derivative (see, e.g., Ratra 1991). The XCDM parametrization of this dark energy model approximates $w$ as a constant, which is accurate during the radiation and matter dominated epochs but not in the current, scalar-field dark energy dominated epoch. This XCDM approximation thus leads to inaccurate predictions for the gravitational lensing considered here, which probes the low redshift universe. We emphasize, however, that unlike a lot of earlier work, we do not work in the XCDM approximation, instead we explicitly integrate the scalar-field dark energy equations of motion.  In $\\S$2 we summarize the data and method used. Results are presented and discussed in $\\S$3. ",
        "conclusions": "Figure 1 shows the CLASS lensing-based constraints on the parameters of the spatially-flat inverse power-law potential scalar field dark energy cosmological model. The likelihood is maximized for $\\Omega_{\\rm m,0} = 0.34$ and $\\alpha = 0$, i.e., a conventional cosmological constant. At 68\\% confidence, $\\alpha < 2.7$ and $0.18 < \\Omega_{\\rm m,0} < 0.62$.\\footnote{ If the SDSS measured VF (Sheth et al.\\ 2003; Mitchell et al.\\ 2004) were used instead of the SSRS2 inferred VF (Chae 2003), these ranges would be narrower and the maximum likelihood estimate of $\\Omega_{\\rm m,0}$ would be $\\sim 0.2$. See \\S 2 and Chae (2004, in preparation) for why we choose to use the SSRS2 VF.} However, at 95\\% confidence both $\\alpha = 8$ and $\\Omega_{\\rm m,0}=1$ are allowed. As mentioned in \\S 2, these results are based on the assumption that the comoving number density of early-type galaxies is unchanged from $z \\sim 1$. However, if there were fewer early-type galaxies at intermediate redshifts compared with the present epoch, the maximum likelihood estimate of $\\Omega_{\\rm m,0}$ would become lower and the confidence ranges for $\\alpha$ and $\\Omega_{\\rm m,0}$ would become narrower. The above results are consistent with, but not as constraining as, those derived from Type~Ia supernova redshift-magnitude data (Podariu \\& Ratra 2000; Waga \\& Frieman 2000). They are also consistent with, but more constraining than, those determined using measurements of angular size as a function of redshift (Chen \\& Ratra 2003a; Podariu et al.~2003).  It is interesting to note that various disparate data sets give consistent constraints on the inverse power-law potential energy density scalar-field dark energy model that weakly favor the conventional cosmological constant over a dynamical scalar field dark energy. However, current results are tentative and future much larger data sets are required to resolve this issue. Future lensing data (e.g.,\\ CLASS2; see \\S 6 of Chae 2003) would be valuable in this respect and are eagerly anticipated.  \\bigskip  We thank the anonymous referee for constructive comments. KHC and DWL acknowledge support from the ARCSEC of KOSEF. GC and BR acknowledge support from NSF CAREER grant AST-9875031 and DOE EPSCoR grant DE-FG02-00ER45824."
    },
    "0403/astro-ph0403583_arXiv.txt": {
        "abstract": "Primordial vector modes describe vortical fluid perturbations in the early universe. A regular solution exists with constant non-zero radiation vorticities on super-horizon scales. Baryons are tightly coupled to the photons, and the baryon velocity only decays by an order unity factor by recombination, leading to an observable CMB anisotropy signature via the Doppler effect. There is also a large B-mode CMB polarization signal, with significant power on scales larger than $l\\sim 2000$. This B-mode signature is distinct from that expected from tensor modes or gravitational lensing, and makes a primordial vector to scalar mode power ratio $\\sim 10^{-6}$ detectable. Future observations aimed at detecting large scale $B$-modes from gravitational waves will also be sensitive to regular  vector modes at around this level. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403060_arXiv.txt": {
        "abstract": "The second USNO CCD Astrograph Catalog, UCAC2 was released in July 2003. Positions and proper motions for 48,330,571 sources (mostly stars) are available on 3 CDs, supplemented with 2MASS photometry for 99.5\\% of the sources.  The catalog covers the sky area from $-90^{\\circ}$ to $+40^{\\circ}$ degrees declination, going up to $+52^{\\circ}$ in some areas; this completely supersedes the UCAC1 released in 2001. Current epoch positions are obtained from observations with the USNO 8-inch Twin Astrograph equipped with a 4k CCD camera. The precision of the positions are 15 to 70 mas, depending on magnitude, with estimated systematic errors of 10 mas or below. Proper motions are derived by utilizing over 140 ground-and space-based catalogs, including Hipparcos/Tycho, the AC2000.2, as well as yet unpublished re-measures of the AGK2 plates and scans from the NPM and SPM plates. Proper motion errors are about 1 to 3 mas/yr for stars to 12th magnitude, and about 4 to 7 mas/yr for fainter stars to 16th magnitude. The observational data, astrometric reductions, results, and important information for the users of this catalog are presented. ",
        "introduction": "The U.S.~Naval Observatory (USNO) operates the 8-inch Twin Astrograph (\\citet{DH}; \\citet{TACH}) currently from its Flagstaff station (NOFS). The program currently underway with this instrument is the USNO CCD Astrograph Catalog project; it is an ongoing program which started at the Cerro Tololo Interamerican Observatory (CTIO) in 1998. Completion of the all-sky, astrometric observations are expected in May 2004.  This second data release, the UCAC2, is a substantial increase in data volume and includes improvements of reduction techniques over the first release, UCAC1, as described in Paper I \\citep{U1}.  The UCAC2 encompasses the entire area of UCAC1 and completely supersedes it. The sky coverage has about doubled and new measurements were obtained from early epoch plates that are used in UCAC2 to significantly improve the proper motions with respect to the UCAC1 release.  The goal of this project is the densification of the reference frame at optical wavelengths; see also \\citet{spie}.  Toward this goal, UCAC provides about a factor of 30 more stars per square degree than the Tycho-2 catalog.  The precision of the UCAC observed positions comes close to the precision of Hipparcos positions at current epochs, and surpasses the precision of Tycho-2 positions at about 10th magnitude and fainter. The UCAC2, a compiled catalog, includes Hipparcos and Tycho observational data as well as virtually all ground-based catalogs used for the Tycho-2 proper motions.  Thus, within the sky area covered, UCAC2 supersedes the Tycho-2 astrometry for stars 10th magnitude and fainter, providing the most precise positions and proper motions available today for catalogs of comparable area coverage.  Users should note some UCAC limitations. Stars brighter than about R = 10, and in particular those brighter than R = 9, can suffer from overexposure effects and generally are based on 2 images of short exposures only. Their errors are higher and this is reflected in the catalog; they should be used with caution when the strictest astrometry is required. The UCAC observations provide only crude magnitudes in a single, non-standard bandpass (between V and R).  To make the catalog more useful to the astronomy community, the Two Micron All Sky Survey \\citep{2mass} J, H, and K$_{S}$ infrared magnitudes are included for the matched sources (99.5\\% of the total UCAC sources). UCAC2 does not provide any trigonometric parallaxes. Systematic errors in the UCAC2 positions are 5 to 10 mas; although very small, these are larger than in the Hipparcos Catalogue.  Along with UCAC2 superseding UCAC1, users should note that the UCAC1 was an observational catalog with attached, preliminary proper motions.  UCAC2 is a compiled catalog of positions and proper motions referred to a standard epoch (J2000.0); the mean CCD observational position is not published.  The level of completeness (about 80\\%) is the same for UCAC1 and UCAC2, avoiding all ``problem cases\" such as elongated images and blended images of close double stars. For the final release (UCAC3) likely both the mean observational data and the ``best\" compiled positions and proper motions will be published, with major improvements in completeness. ",
        "conclusions": "UCAC2 provides the most accurate positions and proper motions available today for most of the stars in the 9th to 16th magnitude range and the 86\\% of the sky covered so far. External, random errors are close to the quoted, internal errors. The 2MASS infrared photometry added to the data will be of benefit for the user, particularly for stellar statistics and galactic kinematics investigations. The average error of proper motions for the R = 13 to 16 magnitude stars is about 6 mas/yr, dramatically improved over the UCAC1, thanks to the inclusion of the Yellow Sky catalog. For brighter stars, with the inclusion of Hipparcos, Tycho, AGK2 and all catalogs used for the Tycho-2 construction, proper motion errors in the 1 to 2 mas/yr range could be achieved.  The high positional accuracy of UCAC at current epoch has been exploited in the minor planet community and has proven essential for occultation predictions \\citep{mpocc} \\footnote{see also \\url{http://www.asteroidoccultation.com} and \\url{http://mpocc.astro.cz}}. The goal, providing a densification of reference stars beyond the Hipparcos / Tycho-2 catalogs, has been achieved. The average density of UCAC2 is 1,360 stars per square degree, with a positional accuracy of stars to 14th magnitude close to the current epoch position errors of the Hipparcos Catalogue.  For most applications UCAC2 supersedes even Tycho-2 in the sky area covered for stars fainter than about 9th magnitude. However, UCAC2 is limited by remaining systematic errors on the 5 to 10 mas level, which -- although very small -- is significantly worse than for the Hipparcos Catalogue. Systematic errors in the UCAC2 proper motions have not been investigated in great detail yet. Comparisons with identified very distant sources indicate no obvious problem, with expected systematic errors on the 1 mas/yr level.  Using UCAC2 for determining a possible system offset (rotation) between HCRF and ICRF at current epochs via observations of counterparts of extragalactic radio sources is severely limited by the remaining systematic errors in UCAC2 positions. To overcome this, as part of the UCAC project, these sources are observed with deep CCD images.  The same fields are simultaneously observed with the UCAC astrograph on the east and west of the pier.  These special observations are not included in the UCAC2 release.  Mean positions derived from these additional observations will have much smaller systematic errors.  One more final data release of UCAC is planned after the completion of the all-sky survey. It is envisioned that the pixel data will be re-processed for UCAC3, which should slightly improve the astrometric accuracy. The main advantage will be a completeness level to over 99\\%, thus providing accurate positions for many known and new double stars. The astrograph could be used for future projects \\citep{smtel}; however, a bigger telescope is required to make significant progress in further densification efforts. USNO has plans for a dedicated, astrometric, robotic, wide-field telescope for an all-sky survey to about 20th magnitude with positions on the 5 to 10 mas level to R = 18$^{m}$ \\citep{urat}."
    },
    "0403/astro-ph0403132_arXiv.txt": {
        "abstract": "We present the discovery of a circumstellar dust disk surrounding AU Microscopium (AU Mic, GJ 803, HD 197481). This young M star at 10 parsec has the same age and origin as $\\beta$ Pictoris, another nearby star surrounded by a dust disk. The AU Mic disk is detected between 50 AU and 210 AU radius, a region where dust lifetimes exceed the present stellar age. Thus, AU Mic is the nearest star where we directly observe the solid material required for planet formation.  Since 85\\% of stars are M-type, the AU Mic disk provides new clues on how the majority of planetary systems might form and evolve. ",
        "introduction": "About 15\\% of nearby main-sequence stars exhibit an excess of far-infrared radiation that points to the existence of circumstellar dust grains ({\\it1}). Dust grains have short lifetimes, and their continued presence implies a source of replenishment.  The solar system has a disk-like distribution of dust that is continually replenished by the sublimation of comets and collisions between asteroids. We therefore infer that similar populations of undetected parent bodies produce dusty debris disks around infrared excess stars. Moreover, evidence for planets can be found by matching density variations in debris disks to theoretical models of how planets gravitationally perturb these disks ({\\it2}, {\\it3}, {\\it4}). In effect, circumstellar debris disks are a signpost for the existence of extrasolar planetary systems.  Direct images of debris disks are rare.  Starlight reflecting off optically thin debris disks is detected at optical and near-infrared wavelengths in only three cases - $\\beta$ Pictoris, HR 4796A and HD 141569 ({\\it5}, {\\it6}, {\\it7}). In three more cases - Vega, Fomalhaut, and $\\epsilon$ Eridani - debris disk structure is seen only at thermal infrared wavelengths ({\\it8}, {\\it9}). $\\beta$ Pic, HR 4796A and HD 141569 are relatively young ($<$20 Myr) main sequence stars that have the largest disk masses, and hence represent the more detectable of the debris disk systems.  $\\beta$ Pic has Galactic space motions in common with two M stars, AU Mic and AT Mic, that have ages $\\sim$20 Myr ({\\it10}). These stars and $\\beta$ Pic may be co-eval sister stars that have separated in space over time due to the small differences in their space motions.  A total of 17 stars may be members of this group with age 8$-$20 Myr ({\\it11}). AU Mic has significant infrared excess at 60 $\\mu$m ({\\it12}, {\\it13}) and recent sub-millimeter data reveal the the presence of cold (40 K) dust, a total dust mass roughly three times smaller than that of $\\beta$ Pic (Table 1), an absence of molecular gas, and a lack of grains within 17 AU of the star ({\\it14}). ",
        "conclusions": "The existence of morphologically similar dust disks around AU Mic and $\\beta$ Pic supports the hypothesis that these are sister stars born at the same time and location.  However, the two disks are not twins. The total mass of dust estimated from the spectral energy distributions is 3.3 times greater for $\\beta$ Pic relative to AU Mic (Table 1).  The relative brightnesses of the two disks in optical data are consistent with this result. To make the comparison, we imagine placing the $\\beta$ Pic dust disk around AU Mic. In Figure 2, we include the midplane surface brightness profile for $\\beta$ Pic using data from ({\\it16}) that is now scaled by factors that account for the AU Mic heliocentric distance and stellar luminosity.  We find that if the disk of $\\beta$ Pic surrounded AU Mic it would be about 1.5 mag arcsec$^{-2}$ brighter than what we measure for the AU Mic disk (Fig. 2). This corresponds to a factor of four greater scattering cross section of $\\beta$ Pic grains relative to AU Mic grains. If we assume that the two disks have exactly the same structure, grain properties, and viewing geometry then the AU Mic disk requires a dust mass that is four times smaller than that of $\\beta$ Pic.  Future observations of disk properties such as the inclination of AU Mic will elucidate the validity of these assumptions, but this result is consistent with the infrared dust luminosity.  The underlying grain properties are also likely to differ due to the weak radiation environment of an M star relative to an A star. AU Mic is 3.6 times less massive than $\\beta$ Pic, and 87 times less luminous (Table 1). For the AU Mic disk, the collision timescale at 100 AU radius is 0.2$-$1.8 Myr assuming a dust optical depth of $\\tau\\sim$10$^{-3} -10^{-4}$, respectively (Fig. S4).  At 200 AU, near the outer boundary of the detected disk, the collision timescale is 0.5$-$5.0 Myr.  Given an age of 8$-$20 Myr for AU Mic, most disk particles have undergone at least one collision.  However, as objects are shattered into smaller pieces, the radiation pressure force around AU Mic is too weak to remove the fragments ({\\it19}).  They can be removed by the system either by joining together to form larger objects, or by spiraling into the star by Poynting-Robertson (PR) drag.  The PR timescales at 100 AU are 0.2$-$1.8 Gyr for 1$-$10 $\\mu$m particles, respectively $-$ many times longer than the age of the system ({\\it1}).  For $\\beta$ Pic, on the other hand, grains a few microns and smaller are quickly ejected by radiation pressure and the disk mass diminishes over time ({\\it20}).  The AU Mic disk should preserve a larger population of sub-micron sized grains, and the mass of solid objects observed today should approximate the primordial disk mass.  In other words, most of the disk seen in our optical scattered light image may consist of primordial solid material.  Within $\\sim$50 AU of the star, the timescales for grain removal by collisions and PR drag become significantly shorter than the stellar age.  Primordial dust at the inner limit of our images (Fig. 1, 2) has mostly vanished, and the grains observed here, as well as those discovered as close as 17 AU from the star ({\\it14}), must be continually replenished by the collisional erosion of much larger objects such as comets and asteroids. The existence of planetesimals in this region lends plausibility to the argument that the same objects will form planets by accretion.  Given that AU Mic is only $\\sim$10 Myr old, we may be able to observe planets that are still in the process of accreting mass, or at least discern disk structure that is sculpted by planet-mass bodies.  Because AU Mic is closer to the Sun than $\\beta$ Pic, the 2$-$30 AU zone where terrestrial and gas giant planets might form can be resolved by current and future instrumentation (Fig. S5).  Planets around AU Mic may also be detected by indirect methods. The low stellar mass means that the star will display a significant astrometric reflex motion (2 milli-arcsec for a Jupiter-analog). The near edge-on orientation favors planet detection by transits of the stellar photosphere.  Finally, if a planet is detected by radial velocity techniques, then the near edge-on orientation gives the planet mass by constraining the sin($i$) ambiguity intrinsic to these measurements.  \\newpage \\noindent{\\bf References and Notes}  \\noindent1. D. E. Backman, F. Paresce, in {\\it Protostars and Protoplanets III}, E. H. Levy and J. I. Lunine, Eds. (University of Arizona Press, Tucson, 1993), pp. 1253-1304.  \\noindent2. F. Roques, H. Scholl, B. Sicardy, B. A. Smith, $Icarus$ {\\bf 108}, 37 (1994).  \\noindent3. J. -C. Liou, H. A. Zook, {\\it Astron. J.} {\\bf 118}, 580 (1999).  \\noindent4. L. M. Ozernoy, N. N. Gorkavyi, J. C. Mather, T. A. Taidakova {\\it Astrophys. J.} {\\bf 537}, L147 (2000).  \\noindent5. B. A. Smith, R. J.Terrile, $Science$ {\\bf 226}, 1421 (1984).  \\noindent6. G. Schneider et al., {\\it Astrophys. J.} {\\bf 513}, L127 (1999).  \\noindent7. M. D. Silverstone et al., {\\it Bull. Am. Astron. Soc} {\\bf 30}, 1363 (1998).  \\noindent8. W. S. Holland et al., $Nature$ {\\bf 392}, 788 (1998).  \\noindent9. J. S. Greaves et al., {\\it Astrophys. J.} {\\bf 506}, L133 (1998).  \\noindent10. D. Barrado y Navascues, J. R. Stauffer, I. Song, J-. P. Caillault, {\\it Astrophys. J.} {\\bf 520}, L123 (1999).  \\noindent11. B. Zuckerman, I. Song, M. S. Bessell, R. A. Webb,  {\\it Astrophys. J.} {\\bf 562}, L87 (2001).  \\noindent12. V. Tsikoudi, {\\it Astron. J.} {\\bf 95}, 1797 (1988).  \\noindent13. I. Song, A. J. Weinberger, E. E. Becklin, B. Zuckerman, C. Chen, {\\it Astron. J.} {\\bf 124}, 514 (2002).  \\noindent14. M. C. Liu, B. C. Matthews, J. P. Williams, P. G. Kalas, {\\it Astrophys. J.}, in press (2004).  \\noindent 15. Materials and methods are available as supporting material on {\\it Science Online}.  \\noindent16. P. Kalas, D. Jewitt, {\\it Astron. J.} {\\bf 110}, 794 (1995).  \\noindent17. To increase the signal-to-noise of the data shown in Fig. 1, we binned the data 3$\\times$3 pixels and then smoothed by a Gaussian function with $\\sigma$=0.5 pixel.   This smoothed image was used to find the maximum outer extent of the disk.  All other measurements were made using the unbinned and unsmoothed image shown in Fig. 1.  \\noindent18. P. Kalas, D. Jewitt,  {\\it Astron. J.} {\\bf 111}, 1347 (1996).  \\noindent19. R. Saija, et al. {\\it Mon. Not. R. Astron. Soc.} {\\bf 341}, 1239 (2003).  \\noindent20. P. Artymowicz, {\\it Astrophys. J.} {\\bf 335}, L79 (1988).  \\noindent21. M. S. Bessel,  F. Castelli,  B. Plez, {\\it Astron. Astrophys.} {\\bf 333}, 231(1998).  \\noindent22. F. Crifo, A. Vidal-Madjar, R. Lallement, R. Ferlet, M. Gerbaldi,  {\\it Astron. Astrophys.} {\\bf 333}, L29 (1997).  \\noindent23. J. D. Larwood, P. G. Kalas, {\\it Mon. Not. R. Astron. Soc.} {\\bf 323}, 402 (2001).  \\noindent24. W. R. F. Dent, H. J. Walker, W. S. Holland, J. S. Greaves, {\\it Mon. Not. R. Astron. Soc.} {\\bf 314}, 702 (2000).  \\noindent25. This work has been supported by the NASA Origins Program under grant NAG5-11769, and the NSF Center for Adaptive Optics, managed by the University of California at Santa Cruz under cooperative agreement No. AST-9876783.    BCM acknowledges support from NSF grant \\#0228963.  MCL acknowledges support from a Hubble Postdoctoral Fellowship (NASA Grant HST-HF-01152.01). The authors acknowledge the insightful contributions of two anonymous referees.  \\noindent {\\bf Supporting Online Material}\\\\ www.sciencemag.org\\\\ Materials and Methods\\\\ Figs. S1, S2, S3, S4 and S5\\\\ Table S1\\\\   \\newpage \\begin{figure} \\epsscale{0.8} \\plotone{fig1.ps} \\caption{ The disk surrounding AU Mic seen in optical scattered light.  North is up, east is left, and each side of this false-color image corresponds to 60$\\arcsec$.  The central dark region is produced by the 9.5$\\arcsec$ diameter focal plane occulting spot which is suspended by four wires and completely masks our direct view of the star.  This image represents 900 seconds total integration in the $R$ band and each pixel corresponds to 4 AU at the distance to AU Mic. Residual light evident near the occulting spot edge in the NE-SW direction is attributed to asymmetries in the point-spread function caused by instrumental scattering and atmospheric seeing. \\label{fig1}} \\end{figure}  \\begin{figure} \\epsscale{0.8} \\plotone{fig2.ps} \\caption{\\small Midplane surface brightness as a function of radius.  The midplane was sampled between 5$\\arcsec$ and 21$\\arcsec$ radius along a strip 1.2$\\arcsec$ wide. We show the mean value from two nights of data with two different PSF subtraction techniques.  The error bars represent one standard deviation of a single measurement. We fit the data between 6$\\arcsec$ and 16$\\arcsec$ radius with power laws that give indices -3.6 and -3.9 for the NW and SE midplanes of AU Mic, respectively.   The radial  profile for the NW midplane has a significant brightness enhancement at $\\sim$9$\\arcsec$ radius that is either intrinsic to the disk or a background object. We also plot the surface brightness of the $\\beta$ Pic disk from ({\\it16}), but with the surface brightness uniformly 3.0 mag arcsec$^{-2}$ fainter to simulate the existence of $\\beta$ Pic's disk around AU Mic at 9.9 pc. This scaling takes into account the fact that the absolute $R$-band magnitude of AU Mic is 5.6 mag fainter than $\\beta$ Pic (Table 1), and at a constant angular radius the $\\beta$ Pic disk is roughly a factor of ($d_{AU Mic} / d_{\\beta Pic})^{-3.6}$ = (9.9 / 19.3)$^{-3.6}$ =11.3 times brighter (i.e. 2.6 mag brighter;  see Eqn. 4 in {\\it18}). \\label{fig2} } \\end{figure}  \\clearpage  \\begin{deluxetable}{lllll} \\tabletypesize{\\small} \\tablecaption{Star (rows 1-8) and disk (rows 9-13) properties for AU Mic and $\\beta$ Pic. The stellar parameters for AU Mic are derived from data given by ({\\it10}, {\\it13}, {\\it21}). For $\\beta$ Pic's stellar parameters, we use ({\\it22}) and references therein.  The $\\beta$ Pic disk $R$-band surface brightness (SB) at 6$\\arcsec$ radius is given in ({\\it16}), while its maximum extent is given in ({\\it23}). \\label{tbl-1}} \\tablewidth{0pt} \\tablehead{ \\colhead{} & \\colhead{AU Mic} & \\colhead{$\\beta$ Pic} } \\startdata Spectral Type \t\t\t\t\t\t&M1Ve \t\t& A5V \\\\ Mass (M$_\\odot$)\t\t\t\t\t&0.5\t\t\t&1.8\\\\ T$_{eff}$ (K) \t\t\t\t\t\t&3500 \t\t& 8200\\\\ Luminosity (L$_\\odot$)\t\t\t\t&0.1\t\t\t&8.7\\\\ Distance (pc)  \t\t\t\t\t\t&9.9\t\t\t& 19.3 \\\\ $V$ (mag)\t\t\t\t\t\t&8.8\t\t& 3.9 \\\\ M$_V$ (mag)\t\t\t\t\t\t&8.8\t\t& 2.4 \\\\ $V-R$\t\t\t\t\t\t\t&0.88\t\t& 0.08\\\\ Disk SB (6$\\arcsec$)\t\t\t\t&20.1$\\pm$0.3\t&15.4$\\pm$0.3\\\\ SB fall-off\\tablenotemark{a}\t\t\t&-3.6 to -3.9\t&-3.8 to -4.1 \\\\ Max. Radius (AU)\t\t\t\t\t&210\t\t&1835  \\\\ $\\tau$=$L_{disk}$/$L_{bol}$\\tablenotemark{b}\t&6.1$\\times$10$^{-4}$ & 3$\\times$10$^{-3}$ \\\\ Total Dust Mass (g)\\tablenotemark{c}\t&6.6$\\times$10$^{25}$& 2.2$\\times$10$^{26}$ \\\\  \\enddata \\tablenotetext{a}{Value of exponent for a power-law fit to disk $R$-band surface brightness fall-off between 6-16$\\arcsec$ radius.  The shallower surface brightness profiles correspond to the NW and NE brightness profiles of AU Mic and $\\beta$ Pic, respectively.  The $\\beta$ Pic values are taken from ({\\it16}). } \\tablenotetext{b}{The fractional dust luminosity, assuming an optically thin disk, determined by taking the ratio of excess infrared luminosity to stellar bolometric luminosity. The values are obtained from  ({\\it14}) and ({\\it1}) for AU Mic and $\\beta$ Pic, respectively. } \\tablenotetext{c}{Dust mass from model fits to the spectral energy distributions taken from ({\\it14}) for AU Mic, and ({\\it24}) for $\\beta$ Pic. } \\end{deluxetable} \\clearpage"
    },
    "0403/astro-ph0403418_arXiv.txt": {
        "abstract": "We present imaging observations of the evolved star IRC+10216 in the CS J=14--13 line at 685.4 GHz and associated submillimeter continuum at $\\sim 2''$ resolution made with the partially constructed Submillimeter Array. The CS J=14--13 line emission from the stellar envelope is well resolved both spatially and spectrally. The strong central concentration of the line emission provides direct evidence that CS is a parent molecule that forms close to the stellar photosphere, in accord with previous images of the lower excitation CS J=2--1 line and inferences from unresolved observations of vibrationally excited transitions. The continuum emission is dominated by a compact, unresolved component, consistent with the photospheric emission, that accounts for $\\sim20\\%$ of the broadband 450~$\\mu$m flux. These are the first interferometer imaging observations made in the semi-transparent 450~$\\mu$m atmospheric window. ",
        "introduction": "The nearby carbon star IRC+10216 (CW Leo, IRAS 09452+1330) is surrounded by a dusty, expanding envelope with relatively simple geometry and kinematics that has been the subject of many studies of astrochemical processes. A rich chemistry develops in the envelope as gas is lost from the star at a high rate ($\\sim3\\times10^{-5}$~M$_{\\odot}$~yr$^{-1}$, Crosas \\& Menten 1997) and is modified by thermodynamic equilibrium and non-equilibrium reactions, photochemical reactions, ion-molecule reactions, and the condensation of dust grains. The range of physical conditions in the envelope make it especially well suited to observations of rotational emission lines from molecules, and numerous species have been identified in spectral scans at millimeter and submillimeter wavelengths (e.g. Cernicharo, Guelin \\&  Kahane 2000; Groesbeck, Phillips \\& Blake 1994, Avery et al. 1992).  The close proximity ($\\sim150$~pc) and copious mass loss of IRC+10216 allows studies of the circumstellar envelope on a wide range of size scales. On the largest scales, an extensive envelope of dust (Mauron \\& Huggins 1999) and gas (Fong, Meixner \\& Shah 2003) has been found to extend to $\\sim0.15$~pc ($\\sim3\\farcm3$) in radius. On much smaller scales, interferometric imaging at millimeter wavelengths has revealed the spatial distribution of molecular species within the central arcminute (e.g. Bieging \\& Tafalla 1993, Guelin, Lucas \\& Cernicharo 1993, Lucas \\& Guelin 1999). Parent molecules are found to be centrally peaked (e.g. SiS, HCN), while daughter molecules are found with shell distributions (e.g. CN, C$_2$H, HNC). High angular resolution is especially important for imaging lines of high excitation that emerge from the innermost regions of the envelope, where the temperatures and densities are highest.  Since the high excitation rotational lines of many common molecules lie at short submillimeter wavelengths, 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) opens new possibilities for probing physical conditions and chemistry by enabling imaging observations with arcsecond resolution. In this {\\em Letter}, we present observations of IRC+10216 in CS J=14--13 emission made with the partially constructed SMA (Ho, Moran \\& Lo 2004). These are the first interferometer images from observations in the atmospheric window centered near 450~$\\mu$m, where $\\sim$30\\% or better transmission occurs at high, dry sites like the summit of Mauna Kea a significant fraction of the time. They build on the pioneering efforts of Carlstrom et al. (1994), who obtained fringes on the nearby single baseline between the Caltech~Submillimeter~Observatory and the James~Clerk~Maxwell~Telescope. The results provide constraints on sulfur chemistry in the inner wind of the IRC+10216 prototype carbon-rich circumstellar envelope. ",
        "conclusions": "\\subsection{Spectrum at Image Center} Figure~\\ref{fig:cs_spectrum} shows the spectrum from the center of the image cube for the full USB correlator band, averaging groups of 8 channels, corresponding to a velocity width of 2.85 km~s$^{-1}$. The CS J=14--13 line is the prominent feature centered at $-26$~km~s$^{-1}$ and an approximate width at zero power of $\\sim30$~km~s$^{-1}$, values which are consistent with the systemic velocity of the source and the expansion velocity of the envelope, respectively. The line shape is difficult to interpret because spatial filtering by the interferometer results in missing flux that varies with velocity across the profile (see \\S\\ref{discussion:line}). The peak brightness temperature of the line at the full angular resolution of the data is well over 100~K. An additional feature appears in the spectrum near $-100$~km~s$^{-1}$ with a double peaked line shape and an approximate width consistent with the circumstellar expansion. A search of the JPL spectral line catalog suggests a tentative identification of this feature with the C$_3$H$_2$ $8(4,4)-7(3,5)$ line at 685.6125565~GHz. This species was previously detected in transitions at lower frequencies (Kawaguchi et al. 1995, Cernicharo et al. 2000). No alternative candidates line were found in the catalog of Cernicharo et al. (2000) that contains 1050 molecular species (Cernicharo, private communication). If the assignment is correct, then the line might be thought to emerge from the external shell where various carbon chain radicals are found (Guelin et al. 1993). This feature appears to be spatially compact, though, which perhaps renders the identification problematic.  \\subsection{Continuum Emission} Figure~\\ref{fig:cont} shows the continuum image obtained from the channels free of strong line emission. Visibilities from both sidebands, separated by 10~GHz, were combined in a multifrequency synthesis, resulting in an effective frequency of 680 GHz. The continuum peak position is consistent with previous determinations, in particular the 95~GHz observation from the IRAM PdBI (Guelin et al. 1993), which is marked by the cross. The peak continuum flux is 3.8~Jy, with an uncertainty dominated by systematic effects.  Previous bolometer observations of IRC+10216 at 450~$\\mu$m indicate that the SMA recovers approximately 20\\% of the broadband flux at this wavelength.  Sandell (1994) noted that IRC+10216 has long period variations in the submillimeter and found $19\\pm3$~Jy in an $18''$ aperture near the 635 day maximum.  Jenness et al. (2002) analyzed several years of SCUBA observations of IRC+10216 at 850~$\\mu$m and 450~$\\mu$m and found a period and maximum flux consistent with the earlier single pixel measurements.  The SMA observations were made approximately a month before the predicted maximum, when $\\sim95\\%$ of the peak flux would be expected. Remarkably, the CS J=14--13 line emission {\\em alone} accounts for $1\\%$ of the flux in the $\\sim68$~GHz wide filter used for the 450~$\\mu$m bolometer measurements. Given the many strong spectral lines excited in the inner envelope at these high frequencies, it is clear that spectral line emission makes an important contribution to the broadband ``continuum'' flux.  A large fraction of the continuum emission from the compact component detected by the SMA at 680~GHz likely comes from the stellar photosphere. Lucas \\& Guelin (1999) report a ``point source'' component in IRAM PdBI observations of $65\\pm1$~mJy at 89~GHz and $486\\pm7$~mJy at 242~GHz, identified as photospheric emission. Assuming an optically thick spectrum, $S_\\nu\\propto\\nu^2$, an extrapolation to 680~GHz gives 3.8~Jy, consistent with the measured value.  Thermal emission from dust in the inner envelope cannot contribute substantial additional flux in the small synthesized beam of the SMA at this high frequency.  \\subsection{CS J=14--13 Emission} \\label{discussion:line} Figure~\\ref{fig:cs_channels} shows images for a series of velocity intervals that span the CS J=14--13 line. Below each of the images is a plot of visibility amplitude vs. $(u,v)$ distance for the corresponding velocity interval to give an idea of spatial extent, which is not easy to ascertain from the images. The overall structure in the visibility amplitudes is consistent with simple expectations for a spherically expanding envelope. At the extreme velocities, the amplitude is approximately constant, as expected for the unresolved ``caps'' of blueshifted and redshifted emission, while at intermediate velocities, the fall-off of visibility amplitude with $(u,v)$ distance demonstrates that the emission is spatially extended and resolved. If the gas is expanding radially with approximately spherical symmetry, then the narrow velocity interval around the central velocity corresponds to a cross section through the envelope in the plane of the sky.  The central peak of CS J=14--13 emission distribution provides additional direct evidence that CS is ``parent'' molecule, in accord with the detection of ro-vibrational transitions in absorption in the infrared (Keady \\& Ridgeway 1993), the detection of vibrationally excited CS emission (Turner 1987, Lucas \\& Guelin 1999,  Highberger et al. 2000), and interferometric imaging of CS J=2--1 emission (Lucas et al. 1995). The CS molecules must originate close to the stellar photosphere, and they are lost from the expanding envelope through reactions that build more complex species, perhaps aided by shocks as suggested by Willacy \\& Cherchneff (1998). The J=14-13 emission distribution shows no indication of the extended $\\sim15''$ radius ring visible in the J=2-1 line, but emission on that size scale, if present, could not be detected on account of the small SMA field of view.  The reaction network for sulfur bearing carbon chains has been explored by Millar, Flores \\& Markwick (2001), who concluded that an initial CS abundance of $4\\times10^{-6}$ is needed to match observations of C$_3$S and C$_5$S. Models where CS is not a parent species, in which the CS fractional abundance rises from $<10^{-11}$ at radius $10^{16}$~cm, or $4\\farcs$5 at 150~pc, to a peak more than four orders of magnitude larger at radius $10^{17}$~cm, could not match observations of the longer carbon chains. The SMA observations provide a robust {\\em lower limit} to the CS abundance within a radius of $\\sim34$~R$_*$ ($2.25\\times10^{15}$~cm) that supports the parent species scenario. Under the assumption of a thermalized level population, the total CS column density is given by \\begin{equation} N(CS) = \\frac{3 k f}{8 \\pi^3 B \\mu^2} \\frac{\\exp{(hBJ_l(J_l+1)/kT)}}{J_l+1} \\frac{T + hB/3k}{1- \\exp{(-h\\nu/kT)}} \\int \\tau dv, \\end{equation} where the molecular constants are $\\mu=1.96$ Debye, $B=24.584$ GHz, and the filling factor $f$ may be estimated as (e.g. Scoville et al. 1986) \\begin{equation} f = \\frac{T_R^*/\\eta_c}{(h\\nu/k)/(\\exp{(h\\nu/kT)-1})(1 - \\exp{-\\tau}) }. \\end{equation} The lower limit to the line flux in a $\\sim2''$ beam is $\\sim 5000$~Jy~km~s$^{-1}$, or $\\sim 1250$~K~km~s$^{-1}$ in brightness units. Assuming small optical depth, $\\tau<<1$, a kinetic temperature equal to the excitation of the upper level of the observed transition (246.8~K), which minimizes the column density required to explain the observed line flux, and a spherical volume with molecular hydrogen density equal to the critical density of the transition of $3.2\\times10^8$~cm$^{-3}$, results in a lower limit to the CS fractional abundance of $\\sim3.4\\times10^{-9}$. The abundance close to the photosphere is likely substantially higher than this estimate; Highberger et al. (2000) derive $3-7\\times10^{-5}$ at a radius $\\sim14$~R$_*$ ($9\\times10^{14}$~cm) from their multi-transition single dish observations. Though the uncertainties are large, models such as those investigated by Millar et al. (2001) where the CS abundance rises from $<10^{-11}$ at radius $10^{16}$~cm are clearly in conflict with the high resolution images of the CS J=14--13 emission.  The SMA observations show directly that the CS molecules are formed close to the stellar photosphere.  Detailed radiative transfer models that account for the observed spatial distribution are needed for an accurate determination of the CS abundance in the inner envelope. The effect of pulsational shocks could be important in establishing the CS abundance as this species is injected into the wind, and careful modeling may provide insight and constraints on the shock chemistry.  Imaging the CS J=14--13 line in the inner envelope of IRC+10216 offers a first demonstration of the efficacy of the SMA in the semi-transparent 450~$\\mu$m atmospheric window. Calibration at these high frequencies should become considerably easier as more receivers for these frequencies are deployed at the SMA, the correlator bandwidth is increased to 2~GHz, and simultaneous operation of a lower frequency band enables phase transfer. Many more evolved stars, as well as other objects detected but with presently unresolved emission at short submillimeter wavelengths, are now accessible to imaging at the arcsecond scale."
    },
    "0403/astro-ph0403304_arXiv.txt": {
        "abstract": "We present the first results from the \\xmm Galactic Plane Survey (XGPS). In the first phase of the programme, 22 pointings were used to cover a region of approximately three square degrees between 19\\deg -- 22\\deg in Galactic longitude and $\\pm$0.6\\deg in latitude.  In total we have resolved over 400 point X-ray sources, at $\\geq 5 \\sigma$ significance, down to a flux limit of $\\sim2 \\times 10^{-14}$ \\ergseccm (2--10~keV). The sources exhibit a very wide range of spectral hardness,  with interstellar absorption identified as a major influence. The source populations detected in the soft (0.4--2 keV) band and hard (2--6 keV) band show surprisingly little overlap.  The majority of the soft sources appear to be associated with relatively nearby stars with active stellar coronae, judging from their high coincidence with bright stellar counterparts.  The combination of the XGPS measurements in the hard X-ray band with the results from  earlier surveys carried out by \\asca and \\chan reveals the form of the low-latitude X-ray source counts over 4 decades of flux. It appears that extragalactic sources dominate below $\\sim10^{-13}$ \\ergseccm (2--10~keV), with a predominantly Galactic source population present above this flux threshold. The nature of the faint Galactic population observed by \\xmm remains uncertain, although cataclysmic variables and RS CVn systems may contribute substantially. \\xmm observes an enhanced surface brightness in the Galactic plane in the 2--6 keV band associated with Galactic Ridge X-ray Emission (GRXE). The integrated contribution of Galactic sources plus the breakthrough of extragalactic signal accounts for up to 20 per cent of the observed surface brightness.  The XGPS results are consistent with the picture suggested from a deep \\chan observation in the Galactic plane, namely that the bulk of the GRXE is truly diffuse. ",
        "introduction": "With the current generation of X-ray astronomy missions, we are  for the first time able to carry out high sensitivity, coherent surveys of selected regions of the Galactic plane.  In particular, the \\xmm mirrors afford a large collecting area ($\\sim 4650 \\rm~cm^{-2}$ total geometric area) with good spatial resolution (FWHM $\\sim6''$ and HEW $\\sim15''$ on-axis) over a wide field of view ($30'$ diameter). In combination with the EPIC CCD cameras, this provides an excellent facility for surveying sky regions subtending many square degrees down to relatively faint flux levels in both the hard ($> 2$ keV) and soft ($< 2$ keV) X-ray bands.  The goal of the \\xmm Galactic Plane Survey (XGPS) is two-fold. The first objective is to study the properties of the Galactic X-ray source population at intermediate flux levels (down to $\\sim 2 \\times 10^{-14}$ \\ergseccm~in the 2--10~keV band  but an order of magnitude fainter in flux terms in the softer 0.4--2~keV band). The second is to search for extended, low X-ray surface brightness features including variations in the underlying diffuse Galactic Ridge X-ray Emission (GRXE; \\citeb{a255}; \\citeb{n317}; \\citeb{pasj38}; \\citeb{a404}; \\citeb{a491}; \\citeb{a505}).  The nature of the X-ray source population at high X-ray fluxes was established by early all-sky surveys and subsequent identification programmes, which revealed that the brightest sources in our Galaxy are predominantly X-ray binaries and supernova remnants. More sensitive surveys of the Galactic plane have since been made, including those made by {\\it ROSAT} (\\citeb{aa246}) and {\\it ASCA} (\\citeb{a134}) complemented by the serendipitous surveys carried out with the Einstein observatory (\\citeb{a278}). Together, these surveys have provided some glimpses of the X-ray source population at lower X-ray fluxes, and hence effectively at lower X-ray luminosities for Galactic objects, although the picture is far from complete. At soft X-ray energies ($<2$ keV) {\\it ROSAT} studies in particular have shown that coronal emission from relatively nearby active stars dominates (\\eg {\\citeb{aa246b}). Above 2 keV the characteristics of the harder population are far less well-defined, although it is clear that accreting binary sources (both X-ray binaries and cataclysmic variables) make a significant contribution.  To date the XGPS  survey has been targeted  at several locations in the Galactic segment between the Galactic Centre and the Scutum Spiral Arm. Here we report the results from the first phase of the XGPS (hereafter XGPS-I), which has entailed a total of 22 \\xmm pointings, covering a region of approximately three square degrees between 19\\deg -- 22\\deg in Galactic longitude and $\\pm$0.6\\deg in latitude. Over 400 discrete point-like X-ray sources have been detected in XGPS-I and in this paper we focus on the properties of this source population and the contribution these discrete sources make to the GRXE. In a second paper (Hands et al. 2004, in preparation) we will present the results of a search for low-surface brightness, spatially extended X-ray sources in the XGPS-I fields and also report on the properties of the underlying diffuse GRXE.   \\begin{table*} \\begin{center} \\caption{Observation log for the XGPS-I. \\label{tab:log}}  \\vspace{0.1in} \\begin{tabular}{lcllrrrr} \\hline  Field & & \\multicolumn{2}{c}{Field Centre (J2000)} & MOS & pn & MOS & pn\\\\ Name  & Observation Date & R.A.   & Dec. & exposure$^a$  & exposure$^b$ & fraction$^c$ & fraction$^c$ \\\\ \\hline  Ridge 1 & 2000-10-08 & 18 26 00.4 & -12 14 55.9 & 8393  & 5914 & 0.95 & 0.85 \\\\ Ridge 2 & 2002-09-21 & 18 26 48.4 & -11 52 48.7 & 13667 & 12046 & 1.00 & 1.00 \\\\ Ridge 3 & 2000-10-11 & 18 27 36.4 & -11 30 40.4 & 12044 & 9648 & 0.95 & 0.85 \\\\ Ridge 4 & 2000-10-12 & 18 28 17.0 & -11 07 53.0 & 9256  & 12998 & 0.94 & 0.42 \\\\ Ridge 5 & 2002-09-17 & 18 29 06.0 & -10 45 03.0 & 13667 & 12046 & 0.98 & 0.93 \\\\ XGPS 1  & 2001-03-08 & 18 25 04.6 & -11 50 00.4 & 7794  & 4797 & 1.00 & 1.00 \\\\ XGPS 2  & 2001-03-10 & 18 27 34.0 & -12 09 20.0 & 9144  & - & 1.00 & - \\\\ XGPS 3  & 2001-03-10 & 18 25 49.0 & -11 28 42.7 & 9144  &  - & 1.00 & - \\\\ XGPS 4  & 2001-03-10 & 18 28 19.4 & -11 48 05.2 & 9144  &  - & 1.00 & - \\\\ XGPS 5  & 2001-03-22 & 18 26 35.7 & -11 07 32.1 & 9994  & 7348 & 0.96 & 0.95 \\\\ XGPS 6  & 2001-03-22 & 18 29 05.9 & -11 26 57.4 & 8794  & 6148 & 0.44 & 0.23 \\\\ XGPS 7  & 2001-03-24 & 18 27 21.2 & -10 46 17.8 & 7637  & 4998 & 0.95 & 0.98 \\\\ XGPS 8  & 2002-09-29 & 18 29 50.8 & -11 05 41.2 & 13167 & 11546 & 1.00 & 1.00 \\\\ XGPS 9  & 2001-03-26 & 18 28 06.4 & -10 25 10.0 & 11894 & 9248 & 0.92 & 0.62 \\\\ XGPS 10 & 2001-10-03 & 18 30 36.2 & -10 44 29.3 & 9767  & 8146 & 0.97 & 0.83 \\\\ XGPS 11 & 2002-09-19 & 18 29 43.9 & -10 24 09.6 & 8409  & 6788 & 0.92 & 0.81 \\\\ XGPS 12 & 2002-09-27 & 18 28 59.1 & -10 06 50.9 & 7367  & 5746 & 1.00 & 1.00 \\\\ XGPS 13 & 2003-04-10 & 18 31 25.5 & -10 24 32.5 & 6666  & 5047 & 0.73 & 0.66 \\\\ XGPS 14 & 2002-03-11 & 18 30 29.3 & -10 02 47.1 & 7269  & 4998 & 1.00 & 1.00 \\\\ XGPS 15 & 2002-03-27 & 18 29 36.6 & -~9 42 41.1 & 7274  & 4998 & 0.82 & 0.96 \\\\ XGPS 16 & 2002-03-15 & 18 32 06.5 & -10 01 51.8 & 7762  & 5486 & 1.00 & 1.00 \\\\ XGPS 17 & 2002-03-29 & 18 31 13.9 & -~9 41 39.4 & 7274  & 4998 & 0.99 & 0.94 \\\\ \\hline &            &           &           &       & \\\\ \\multicolumn{4}{l}{$^a$ Total exposure for the MOS 1 camera (s)} \\\\ \\multicolumn{4}{l}{$^b$ Total exposure for the pn camera (s)} \\\\ \\multicolumn{4}{l}{$^c$ Fraction of the exposure time used in producing images} \\\\  \\end{tabular} \\end{center}  \\end{table*} ",
        "conclusions": "The XGPS-I survey, which covers approximately three square degrees of the Galactic Plane near $l=20$\\deg, has resulted in a catalogue containing over 400 discrete X-ray sources. The measured X-ray source counts trace the source population down to a limiting flux of $\\sim 2 \\times 10^{-14}$ \\ergseccm~in the 2--10 keV band at which point the source density is between 100--200 sources per square degree. Consistent with an earlier \\chan study, the source counts at this flux are predominately due to extragalactic sources, despite the fact that the fluxes of extragalactic objects are significantly suppressed by absorption in the Galactic plane. However, the conclusion of the present work is that at fluxes above $10^{-13}$  \\ergseccm (2--10 keV) Galactic source populations do come to the fore.  The Galactic source population observed between $10^{-13}$ and $10^{-12}$ \\ergseccm~could comprise largely CVs and RS CVn systems with X-ray luminosities in the range $10^{30-32}$ \\ergsec~but the details remain uncertain on the basis  of the X-ray information alone. Extensive programmes to identify and characterise  optical/infra-red counterparts are required, although this will be taxing given the high obscuration and high object density in the Galactic plane.  The present work demonstrates that the strategy of the XGPS programme, namely the use of shallow observations to give relatively wide angle coverage is close to optimum in terms of maximising the number of Galactic source detections."
    },
    "0403/astro-ph0403074_arXiv.txt": {
        "abstract": "{ We have designed and tested a new relativistic Lagrangian hydrodynamics code, which treats gravity in the conformally flat approximation to general relativity. We have tested the resulting code extensively, finding that it performs well for calculations of equilibrium single-star models, collapsing relativistic dust clouds, and quasi-circular orbits of equilibrium solutions. By adding in a radiation reaction treatment, we compute the full evolution of a coalescing binary neutron star system.  We find that the amount of mass ejected from the system, much less than a percent, is greatly reduced by the inclusion of relativistic gravitation.  The gravity wave energy spectrum shows a clear divergence away from the Newtonian point-mass form, consistent with the form derived from relativistic quasi-equilibrium fluid sequences.} ",
        "introduction": "It has long been recognized that coalescing binary neutron star (NS) systems are a leading candidate to be the first observed source of gravity waves (GW).  With LIGO, GEO, and TAMA all taking scientific data, and VIRGO in the commissioning stage, it is growing increasingly important to have quantitatively accurate predictions of the GW signals we expect to measure during the merger process.  Besides their use in aiding detections, these predictions are crucial for determining important physical information about the mass, radius, and equation of state (EOS) of NS from GW observations.  Calculations of binary NS coalescence have been performed for many years, beginning with studies in  Newtonian gravity.  It was recognized all along, however, that general relativity (GR) will play an important role during the merger, since the characteristic gravitational fields and velocities are squarely within the relativistic regime. As a result, increasingly sophisticated gravitational formalisms have been used in hydrodynamical calculations, starting with post-Newtonian treatments \\cite{FR1,FR2,FR3,APODR,SNO}, many of which were based on a formalism developed by Blanchet et al. \\cite{BDS} which includes all lowest-order 1PN effects as well as lowest order dissipative effects from gravitational radiation reaction losses. More recently, calculations have been performed in full general relativity \\cite{SU1,SU2,STU}. Unfortunately, the PN approximation breaks down during the merger when higher-order relativistic effects grow significant, and fully relativistic calculations typically introduce numerical instabilities which limit the amount of time for which a calculation will remain accurate.  A middle ground is provided by the conformally flat (CF) approximation, developed originally by Wilson et al. \\cite{WMM}, which includes much of the non-linearity inherent in GR, but results in a set of coupled, non-linear, elliptic field equations, which can be evolved stably.  We assume that the spatial part of the GR metric is equal to the flat-space form, multiplied by a conformal factor which varies with space and time, the metric taking the form \\begin{equation} ds^2=-(N^2-B_i B^i)dt^2-2B_i dt dx^i+A^2 \\delta_{ij}dx^i dx^j. \\end{equation} While this approach cannot reproduce the exact GR solution for a general matter configuration, it is exact for spherically symmetric systems, and yields solutions which agree with those calculated using full GR to within a few percent for many systems of interest \\cite{MMW}. ",
        "conclusions": ""
    },
    "0403/astro-ph0403597_arXiv.txt": {
        "abstract": "We analyse the anisotropy of the cosmic microwave background (CMB) in hyperbolic universes possessing a non-trivial topology with a fundamental cell having an infinitely long horn. The aim of this paper is twofold. On the one hand, we show that the horned topology does not lead to a flat spot in the CMB sky maps in the direction of the horn as stated in the literature. On the other, we demonstrate that a horned topology having a finite volume could explain the suppression of the lower multipoles in the CMB anisotropy as observed by COBE and WMAP. ",
        "introduction": "A fundamental problem in cosmology is the large-scale geometry of the Universe, in particular its spatial curvature and topology. Since the Einstein gravitational field equations are differential equations, they determine the local properties of space-time, but not the global structure of the Universe at large. In the so-called concordance model of cosmology it is assumed that the Universe is at large scales spatially flat and possesses the trivial topology, implying that it has infinite volume. In the framework of inflationary scenarios, these properties are supposed to be determined by the initial conditions at the Big Bang. However, since we are lacking a theory of quantum gravity, the initial conditions cannot be derived from first principles. Instead, one can analyse the recent astronomical data, in particular on the cosmic microwave background (CMB) radiation, in order to deduce the curvature and the topology of the Universe at large scales.   In 1992, COBE \\cite{Smoot_et_al_1992} made the spectacular discovery of the temperature fluctuations of the CMB and thus provided important clues about the early Universe and its time evolution. Expanding the observed temperature fluctuations $\\delta T(\\hat n)$ across the microwave sky into spherical harmonics $Y_l^m(\\hat n)$, yields the expansion coefficients $a_{lm}$ which in turn lead to the {\\it multipole moments} \\begin{equation} \\label{Eq:C_l} C_l \\; := \\; \\frac 1{2l+1} \\, \\sum_{m=-l}^l \\, | a_{lm} |^2 \\end{equation} and the {\\it angular power spectrum} $\\delta T_l^2 := l(l+1) C_l /(2\\pi)$. In particular, COBE \\cite{Smoot_et_al_1992} detected in the angular power spectrum of the CMB a low quadrupole moment $C_2$ corresponding to a strange suppression of power on large angular scales. It was soon realized that although the standard cosmological models are in agreement with the CMB anisotropy on small and medium scales, they fail to match the loss of power on large angular scales, especially those corresponding to the quadrupole moment. This observation was one of the motivations to study non-trivial topologies (see the reviews \\cite{Lachieze-Rey_Luminet_1995,Levin_2002}). In particular, compact hyperbolic universes were studied by several authors \\cite{Bond_Pogosyan_Souradeep_1998,Bond_Pogosyan_Souradeep_1999a,% Bond_Pogosyan_Souradeep_1999b,Cornish_Spergel_Starkman_1998,Aurich_1999,% Cornish_Spergel_1999,Inoue_Tomita_Sugiyama_1999,Aurich_Steiner_2000}. For these compact hyperbolic universes, it was shown that the non-trivial topology leads indeed to a suppression of $C_l$ for small values of $l$. Thus the observed loss of power on large angular scales was interpreted as a clear hint to a non-trivial topology of our Universe.  The first findings of NASA's explorer mission ``Wilkinson Microwave Anisotropy Probe'' (WMAP) \\cite{Hinshaw_et_al_2003} have tremendously increased our knowledge of the temperature fluctuations of the CMB, since WMAP has measured the anisotropy  of the CMB radiation over the full sky with high accuracy. The WMAP data confirm  not only COBE's measurement of the low quadrupole moment $C_2$, but display in the temperature (auto-) correlation function $C(\\vartheta)$ $\\,(\\hat n \\cdot \\hat n' = \\cos \\vartheta)$ \\begin{equation} \\label{Eq:C_theta} C(\\vartheta) \\; \\simeq \\; \\frac 1{4\\pi} \\, \\sum_{l=2}^\\infty \\, (2l+1) \\, C_l \\, P_l(\\cos\\vartheta) \\end{equation} very weak correlations at wide angles, $70^\\circ \\lesssim \\vartheta \\lesssim 150^\\circ$, see the dashed curve in figure \\ref{Fig:C_theta_concordance}. (Note that the monopole and dipole are not included in the sum (\\ref{Eq:C_theta}).) At the largest angles, $\\vartheta \\gtrsim 160^\\circ$, the WMAP-data display even a ``correlation hole'', i.\\,e.\\ negative values of $C(\\vartheta)$. In figure \\ref{Fig:C_theta_concordance} we also show as a dotted curve the theoretical prediction according to the concordance model using the best-fit values for the cosmological parameters as obtained by WMAP \\cite{Hinshaw_et_al_2003}. It is seen that the concordance model, i.\\,e.\\ the best-fit $\\Lambda$CDM model, does not reproduce the experimentally observed suppression at $\\vartheta \\gtrsim 60^\\circ$ and the observed correlation hole.  Since the correlation function $C(\\vartheta)$ emphasizes large angular scales and thus the low $l$-range, it is an ideal indicator function to search for a fingerprint of a possible non-trivial topology of the Universe.   \\begin{figure}[htb] \\begin{center} \\hspace*{-12pt}\\begin{minipage}{9cm} \\includegraphics[width=9cm]{psplots/C_theta_concordance.ps} \\put(-1,4){$\\vartheta$} \\put(-274,150){$C(\\vartheta)$} \\put(-276,130){$[\\mu \\hbox{K}^2]$} \\end{minipage} \\vspace*{-10pt} \\end{center} \\caption{\\label{Fig:C_theta_concordance} The correlation function $C(\\vartheta)$ from the WMAP-data (dashed curve) in comparison with the concordance model (dotted curve). } \\end{figure}   The evidence for a low quadrupole and missing power on large scales is, however, currently heavily discussed. The low value of the quadrupole obtained by the WMAP team \\cite{Hinshaw_et_al_2003} has been criticized by the way how the foregrounds \\cite{Bennett_et_al_2003b} are taken into account which are mainly caused by free-free, synchrotron and dust emission. Alternative reconstructions of the true cosmological signal from the five frequency bands measured by WMAP are discussed in \\cite{Tegmark_deOliveira_Costa_Hamilton_2003,Eriksen_Banday_Gorski_Lilje_2004} which lead to a higher quadrupole moment. The observation that the plane of the quadrupole and two of the three planes of the octopoles are aligned towards the ecliptic might be seen as a hint of an unknown source or sink of CMB radiation in the outer solar system or as an unrecognized systematic \\cite{Schwarz_Starkman_Huterer_Copi_2004}. The influence of the approximation of the likelihood function used for the angular power spectra analysis is investigated in \\cite{Slosar_Seljak_Makarov_2004}, where for the lowest multipoles an exact likelihood function estimation is presented. The statistical significance of the low value of the quadrupole moment is also discussed in \\cite{Efstathiou_2003} who found the WMAP results to be consistent with the concordance $\\Lambda$CDM model. The statistic is limited by the one-sky realization we are able to measure, also called cosmic variance. To circumvent this problem it is suggested to utilize the fact that the CMB polarization is sourced by the local temperature quadrupole. On the one hand, the polarization signal contains information about the quadrupole moment at the reionization epoch. This could lead to a probability for the observed quadrupole moment of the order of $10^{-4}$ \\cite{Skordis_Silk_2004} compared with the concordance model. On the other hand, future measurements of the linear polarization of the CMB towards clusters of galaxies can betray the local quadrupole moment at the location of the cluster \\cite{Kamionkowski_Loeb_1997,Sazonov_Sunyaev_1999,Seto_Sasaki_2000} thus circumventing the cosmic variance limit.   Furthermore, we would like to mention that ``the question whether the geometry of the three-dimensional space of astronomy might be non-Euclidean'' \\cite{Chandrasekhar_1986} has already been posed by Schwarzschild \\cite{Schwarzschild_1992,Schwarzschild_1900,Schwarzschild_1998} in 1900, fifteen years before the founding of general relativity! He stated the problem as follows (see p.\\,32 in \\cite{Schwarzschild_1992}): ``As must be known to you, during this century [meaning the 19th century] one has developed non-Euclidean geometry (besides Euclidean geometry), the chief examples of which are the so-called spherical and pseudo-spherical spaces. We can wonder how the world would appear in a spherical or pseudo-spherical geometry with possibly a finite radius of curvature \\dots One would then find oneself, if one will, in a geometrical fairyland; and one does not know whether the beauty of this fairyland may in fact be realized in nature.'' Schwarzschild ``actually estimated limits to the radius of curvature of the three-dimensional space with the astronomical data available at his time and concluded that if the space is hyperbolic its radius of curvature cannot be less than 64 light years and that if the space is spherical its radius of curvature must at least be 1600 light years'' \\cite{Chandrasekhar_1986}. In an Addendum to his article ``On the permissible scale of curvature of space'' \\cite{Schwarzschild_1900} he mentioned also the possibility of spaces with non-trivial topology by referring to Clifford-Klein space forms. And he emphasized that such spaces do not necessarily lead to infinite universes as commonly assumed, even in the case of Euclidean or hyperbolic geometry. Schwarzschild concluded that the only condition imposed by astronomical data is that the volume of the Universe must be larger than the visible system of stars.  It is the purpose of this paper to investigate two models of the Universe whose global topology is not the universal covering space of their spatial geometry. The spatial geometry of the models is the pseudo-spherical or hyperbolic space, i.\\,e.\\ it has negative curvature. The corresponding hyperbolic universes are non-compact and possess an infinitely long horn. The first model was introduced in 1976 by Sokolov and Starobinskii \\cite{Sokolov_Starobinskii_1976} and is described by a fundamental cell ${\\cal F}$ having infinite volume. The second model, called the Picard model \\cite{Picard_1884}, has also an infinitely long horn, but nevertheless possesses a finite volume. The Picard model has been investigated in detail in the context of quantum chaos \\cite{Matthies_1995,Steil_1999,Then_2003}.  The Sokolov-Starobinskii model has been studied already in \\cite{Sokolov_Starobinskii_1976} and \\cite{Levin_Barrow_Bunn_Silk_1997} (see also the review \\cite{Levin_2002}). In \\cite{Levin_Barrow_Bunn_Silk_1997} it has been claimed that the periodic horn topology produces a flat spot in the temperature map of the CMB even when multiple images of astronomical sources are unobservable. The flat spot, i.\\,e.\\ a large flat region in the CMB sky map corresponding to a negligibly small metric perturbation, is argued in \\cite{Levin_Barrow_Bunn_Silk_1997} to result from the exponential decay of the eigenmodes in the direction of the horn.  We shall show in the following that the flat spot found in \\cite{Levin_Barrow_Bunn_Silk_1997} is a consequence of a too low wavenumber cut-off $k_c$. For the cut-off chosen in \\cite{Levin_Barrow_Bunn_Silk_1997}, the eigenmodes cannot produce indeed a perturbation in the horn at the position which corresponds to the distance of the surface of last scattering (SLS). As a consequence, the sky maps computed in \\cite{Levin_Barrow_Bunn_Silk_1997} do not show temperature fluctuations in the horn. However, if the cut-off is sufficiently increased, we shall demonstrate that there is no suppression  in the horn and therefore no flat spot.  The universes described by the Sokolov-Starobinskii and Picard model, respectively, possess negative spatial curvature in contrast to the concordance model corresponding to a spatially flat universe. Therefore the question arises whether a negatively curved universe is realized in nature. In \\cite{Aurich_Steiner_2002b,Aurich_Steiner_2003} we have analysed the CMB data and the magnitude redshift relation of supernovae type Ia in the framework of quintessence models and have shown that these data are consistent with a nearly flat hyperbolic geometry of our Universe if the optical depth $\\tau$ to the SLS is not too big. The restriction comes from the large amplitude of the fluctuations at large scales in the CMB. However, it was shown in \\cite{Conversi_Melchiorri_Mersini_Silk_2004} that by replacing the quintessence component having a rest frame sound velocity of $c_s = c$ by a generalized dark matter component \\cite{Hu_1998} with a vanishing rest frame sound velocity $c_s = 0$, the amplitude in the CMB anisotropy is much smaller at large scales. Thus with such a generalized dark matter component universes with negative curvature are permissible even for larger optical depths $\\tau$.  Recently, an ellipticity analysis of the CMB maps has been reported \\cite{Gurzadyan_et_al_2003a,Gurzadyan_et_al_2003b,Gurzadyan_et_al_2004} which gives further support to a hyperbolic spatial geometry of the Universe. In these analyses hot and cold anisotropy spots in the CMB maps have been studied in terms of shape for various temperature thresholds. Analysing with the same algorithm the COBE-DMR, BOOMERanG 150 GHz and WMAP maps, an ellipticity of the anisotropy spots has been found of the same average value (around 2) for these experiments. The WMAP data confirm the effect for scales both smaller and larger than the horizon at the SLS. This suggests that the effect is not due to physical effects at the SLS, and can arise after, while the photons are moving freely in the Universe.  Finally, we would like to mention that recently a model was presented \\cite{Luminet_Weeks_Riazuelo_Lehoucq_Uzan_2003} possessing positive curvature and a finite volume with the shape of the Poincar\\'e dodecahedral space. The authors of ref.\\ \\cite{Luminet_Weeks_Riazuelo_Lehoucq_Uzan_2003} calculated the CMB multipoles for $l=2,3$ and 4, set the overall normalization factor to match the WMAP data at $l=4$ and examined the prediction for $l=2$ and 3. They found a weak suppression of the power at $l=3$ and strong suppression at $l=2$ in agreement with the WMAP observations. However, in ref.\\ \\cite{Luminet_Weeks_Riazuelo_Lehoucq_Uzan_2003} only the modes up to $k_c=30$ have been used, and it thus remains the question about how this low cut-off affects the integrated Sachs-Wolfe (ISW) contribution. Our experience shows that increasing the cut-off usually enhances the ISW contribution. ",
        "conclusions": "A large part of this paper was devoted to the question of the existence of flat spots in the CMB sky maps for universes with a horned topology. By the example of two such universes, the Sokolov-Starobinskii model and the Picard model, we showed that the infinitely long horn does not lead to flat spots, i.\\,e.\\ to a suppression of the CMB fluctuations in the horn, if the wavenumber cut-off $k_c$ is chosen sufficiently large. The flat spots reported earlier \\cite{Levin_Barrow_Bunn_Silk_1997,Levin_2002} are thus seen as the result of taking not enough modes into account in the expansion of the metric perturbation. We conclude that the CMB sky maps do not reveal a signature of the horned topology in form of flat spots.  Another main point of this paper was to show that a universe with a horned topology, but with a finite volume, such as the Picard model, can explain the loss of power at large angular scales in the CMB anisotropy, as observed by COBE and WMAP. However, if the volume is infinite, as in the Sokolov-Starobinskii model, we have demonstrated that the low quadrupole moment is not reproduced. There is, however, some controversy of how serious one has to take this low value. The small values of the first few multipole moments $C_l$ have, on the one hand, not been taken so serious to require some new physics, but rather have been considered as a manifestation of cosmic variance or an unsufficient consideration of Galactic emission \\cite{Efstathiou_2003} or caused by the local supercluster \\cite{Abramo_Sodre_2003}. From this point of view the Sokolov-Starobinskii model is a viable model. On the other hand, the suppression has been taken as a hint to new physics such as new information on the inflation potentials \\cite{Cline_Crotty_Lesgourgues_2003,% Contaldi_Peloso_Kofman_Linde_2003,Feng_Zhang_2003}. However, these potentials have to be fine-tuned such that the arising power spectra are suppressed around the present day cosmological horizon \\cite{Contaldi_Peloso_Kofman_Linde_2003}. Such a fine-tuning is not necessary in the case of a non-trivial topology in a universe with negative spatial curvature due to the Mostow rigidity theorem \\cite{Mostow_1973,Prasad_1973}. Thus if the curvature scale is fixed by the densities $\\Omega_x$, all side-lengths of the fundamental cell ${\\cal F}$ are determined, and in turn the comoving wavenumber $k$ at which the spectrum is suppressed, not because of the initial power spectrum but simply because of the absence of modes below the lowest wavenumber $k_1$.  In models with a non-trivial topology, certain points at the surface of last scattering can be identical due to the periodicity condition. These matching points are located on pairs of circles with the same radius. Along two such circles the temperature fluctuations $\\delta T$ produced by the naive Sachs-Wolfe effect are the same, which is called the circles-in-the-sky signature \\cite{Cornish_Spergel_Starkman_1998b}. In \\cite{Cornish_Spergel_Starkman_Komatsu_2003} a search for such circles in the WMAP data was carried out for nearly back-to-back circles, i.\\,e.\\ for circles whose centers have a distance greater than 170$^\\circ$ and whose radii are  greater than 25$^\\circ$ on the sky. They found no signature and rule out all topologies having such circles. However, it should be kept in mind that only the naive Sachs-Wolfe contribution leads to identical temperature fluctuations. The integrated Sachs-Wolfe contribution arises on the photon path to the observer, which is not identified for the observer and the ``copy'' of the observer. The same is valid for the Doppler contribution, since the observer and its copy see another projection of the velocity, in general. Furthermore, there are a lot of other secondary contributions to the temperature fluctuations. Nevertheless, it is claimed in \\cite{Cornish_Spergel_Starkman_Komatsu_2003} that the naive Sachs-Wolfe contribution is strong enough for the identification of circles-in-the-sky. The Picard model is not ruled out by this work, since there are no nearly back-to-back circles. For the model with $\\Omega_{\\hbox{\\scriptsize mat}} = 0.3$ and $\\Omega_\\Lambda = 0.65$ with $x_3=1.6$, we obtain 40 pairs, where the largest distance of the centers is at 145$^\\circ$ which is not covered by the study in \\cite{Cornish_Spergel_Starkman_Komatsu_2003}. The model with $\\Omega_{\\hbox{\\scriptsize mat}} = 0.35$ and $\\Omega_\\Lambda = 0.6$ has only 32 circle pairs. The number of paired circles increases if the observer is posited higher up in the horn, i.\\,e.\\ at larger values of $x_3$. For the extreme position at $x_3=5$ one obtains 275 circle pairs with distances up to 168.6$^\\circ$ for $\\Omega_{\\hbox{\\scriptsize mat}} = 0.3$ and $\\Omega_\\Lambda = 0.65$. These circles could have been detected in \\cite{Cornish_Spergel_Starkman_Komatsu_2003}. However, for more generic observers, which are not sitting extremely high in the horn, the separation of the circles is too small such that the model cannot yet be ruled out.  In conclusion, we would like to emphasize again that the Picard model studied in this paper is in nice agreement with the observed suppression of power on large scales in the angular power spectrum of the CMB (see figures \\ref{Fig:Picard_gen_ran_m30_q00_l65_h65_C_theta}, \\ref{Fig:Picard_gen_ran_m35_q00_l60_h65_C_theta}, \\ref{Fig:Picard_gen_Spitze_ran_m30_q00_l65_h65_C_Theta_cusp} and \\ref{Fig:Picard_gen_ran_m30_q00_l65_h65_Merge_C_Theta}). This is in contrast to the concordance model which does not reproduce the experimentally observed suppression at $\\vartheta \\gtrsim 60^\\circ$ and the observed correlation hole at $\\vartheta \\gtrsim 160^\\circ$. If future observations will confirm the WMAP data but with smaller errors, this can be interpreted as a clear hint to a non-trivial topology of our Universe having negative spatial curvature and a finite volume."
    },
    "0403/astro-ph0403242_arXiv.txt": {
        "abstract": "We present the first determination of the 15$\\mu$m luminosity function of galaxies from the European Large Area ISO survey (ELAIS) southern fields. We have adopted a new criterion to separate the quiescent, non-evolving and the starburst, evolving populations based on the ratio of mid-infrared to optical luminosities. Strong evolution is suggested by our data for the starburst galaxy population, while normal spiral galaxies are consistent with no evolution. The starburst population must evolve both in luminosity and in density with rates of the order $L(z) \\propto (1+z)^{3.5}$ and $\\rho(z) \\propto (1+z)^{3.8}$ up to $z \\sim 1$. The evolutionary parameters of our model have been tested by comparing the model predictions with other observables, like source counts at all flux density levels (from 0.1 to 300 mJy) and redshift distributions and luminosity functions at high-$z$ (0.7 $< z <$ 1.0 from HDF-N data). The agreement between our model predictions and the observed data is remarkably good. We use our data to estimate the star-formation density of the Universe up to z=0.4 and we use the luminosity function model to predict the trend of the star-formation history up to $z=1$. ",
        "introduction": "\\label{intro_sec} The extragalactic background light shows that the emission from galaxies at infrared and sub-millimeter wavelengths is an energetically significant component of the Universe. This emission originates from star-formation activity and active galactic nuclei. The precise contribution from each type of activity is still debated. It is thus important to our understanding of galaxy and AGN formation to study those populations that emit a substantial amount of light at infrared (IR) wavelengths in their rest frame. In particular, data from deep {\\it Infrared Space Observatory} (ISO) surveys at 15 $\\mu$m (i.e. \\citealp{1999A&A...351L..37E}; \\citealp{1999ApJ...517..148F}; \\citealp{2001MNRAS.325.1173L}; \\citealp{2003A&A...407..791M}) seem to require strong evolution for galaxies emitting in the infrared wavebands. This result, supported also by the detection of a substantial cosmic Infrared Background in the 140 $\\mu$m - 1 mm range \\citep{1996A&A...308L...5P,1998ApJ...508...25H,1999A&A...344..322L}, has stimulated the development of several evolutionary models for IR galaxies (i.e. \\citealp{2001ApJ...549..745R}, \\citealp{2001A&A...378....1F}, \\citealp{2001ApJ...556..562C}, \\citealp{2003ApJ...587...90X}). All these models fit with different degrees of success the IR/sub-millimeter source counts and the cosmic Infrared Background within the present uncertainty limits, but suffer of parameter degeneracy and none of them is based on a local luminosity function (LLF) obtained from 15-$\\mu$m data, being all extrapolated from different IR wavelengths (12, 25, 60 $\\mu$m).  So far, complete spectroscopic samples of 15-$\\mu$m sources have been obtained only in small fields (i.e. HDF-N: \\citealp{1999usis.conf.1023A}; HDF-S: \\citealp{2002MNRAS.332..549M}, \\citealp{2003A&A...403..501F}), too small and too deep to allow a detailed study of the local luminosity function.  The ELAIS survey is the largest open time project conducted by ISO \\citep{2000MNRAS.316..749O}, mapping an area of $\\approx$ 12 deg$^2$ at 15 $\\mu$m and 90 $\\mu$m.  The final, band-merged ELAIS catalogue has recently been completed \\citep{MRR03}. The spectroscopic data is most complete in the southern fields and in this paper we present an analysis of the 15-$\\mu$m luminosity function derived from these data. This is the first determination of the 15-$\\mu$m luminosity function and its evolution, constrained by all the available observables in this band (source counts from IRAS to the deepest ISOCAM flux densities; redshift distributions at low and high-$z$ using both data from ELAIS and data from the deeper HDF-N survey). The model fitting the LF is then used to estimate the star-formation history of the Universe.  This paper is structured as follows. In Section \\ref{sample_sec} we present our data sample. In Section \\ref{kcorr_sec} we discuss the adopted IR and optical K-corrections. In Section \\ref{fdl_sec} we show the method used to compute the 15-$\\mu$m LF and present the results. In Section \\ref{llf_section} we compare our LLF determination with previous ones. In Section \\ref{evolution_sec} we discuss the evolution rates derived from our data, and compare the model predictions with the observable constraints and with other literature models. In Section \\ref{concl_sec} we present our conclusions.  Throughout this paper we will assume $H_0 = 75$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m = 0.3$ and $\\Omega_{\\Lambda} = 0.7$, unless explicitly stated. ",
        "conclusions": "\\label{concl_sec}  We have presented the first direct determination of the 15-$\\mu$m luminosity function and its cosmic evolution for galaxies.  As previously found by other authors, three populations of sources give rise to the 15-$\\mu$m emission: the normal, the starburst and the AGN populations, characterized by different cosmic evolution. In this work we have analysed the galaxy component only (quiescent plus actively starforming). The contribution of the AGN component is discussed in a separate paper (\\citealp{2002MNRAS.332L..11M}; Matute et al. in preparation).  The analysis is based on data from the ELAIS southern fields survey. The sample is composed by $\\sim$150 galaxies in the redshift interval 0.0$<$z$\\le$0.4 and covers a large flux density range intermediate between the IRAS and the deep ISOCAM surveys (0.5$\\le{S}\\le{50}$ mJy). Differently from other authors, we have adopted in this work the $L_{15{\\mu}m}/L_{\\rm R}$ ratio as a criterion to separate the quiescent, non-evolving and the starburst, evolving populations. This criterion, suggested by the existing correlation between $L_{15{\\mu}m}/L_{\\rm R}$ and the amount of activity in galaxies, is a posteriori supported by the results of the V/V$_{max}$ analysis on the two populations defined on this basis.  The main results of our analysis are:  \\begin{enumerate}  \\item{In the ML analysis we have simultaneously fitted both the evolution rates and the shape parameter of the local LF for both the spiral and the starburst populations. We have assumed that the spiral population does not evolve, while we have let to evolve the starburst population both in luminosity and in density. Since the two populations sample different luminosity ranges, we have obtained an accurate determination for the faint end of the LLF for the quiescent component, while the knee and the $\\sigma$ parameters of the LLF are better constrained for the starburst one. The evolution found for the active population is $\\sim(1+z)^{3.5}$ in luminosity and $\\sim(1+z)^{3.8}$ in density, up to $z_{break}\\sim1$.}   \\item{Our total 15-$\\mu$m LLF is in agreement with previous determinations derived from the IRAS data. On the contrary, the LLFs for different populations derived by different authors have not the same level of consistency. While in our subdivision the quiescent population is expected to dominate locally over all the luminosity range, in other models (i.e. \\citealp{2003ApJ...587...90X}) the starburst population dominates locally at high luminosities, leading to a large discrepancy in the model predictions.}   \\item{To test the evolution parameters with higher degree of confidence, we have compared our model predictions with all the observables existing in literature, over all the $z$ and flux ranges (source counts, luminosity functions, $z$-distributions). Our best-fitting model well reproduces all the observables. In the critical interval 1${\\lsimeq}S{\\lsimeq}$3 mJy, where the source counts from different surveys show the larger discrepancies, our model is intermediate between the data from ELAIS-S1 \\citep{2002MNRAS.335..831G} and the data from the deep surveys \\citep{1999A&A...351L..37E}. On the other hand, in the flux range 3${\\lsimeq}S{\\lsimeq}$10 mJy, our model is well consistent with existing data, differently from the \\cite{2003ApJ...587...90X} model, whose predicted differential sources counts are at least a factor of 3 higher than the data.}  \\item{Using the evolutionary model found for the 15-$\\mu$m galaxies and the data points from the $1/V_{max}$ LF analysis, we have estimated the star-formation rate density. The redshift range sampled by our data ($0.0 < z \\le0.4$) is of particular interest, since it provides an estimate of the star-formation at relatively low redshift, but not so local to be affected by clustering and local dishomogeneities. We find $\\dot \\rho$=(0.025$\\pm$0.007) hM$_\\odot$yr$^{-1}$Mpc$^{-3}$ and $\\dot \\rho$=(0.043$\\pm$0.020) hM$_\\odot$yr$^{-1}$Mpc$^{-3}$ at the two mean redshifts $z=0.12$ and $z=0.27$, respectively. At $z \\lsimeq 0.4$ our model predictions are well consistent with other estimates derived from UV, optical and Mid-IR data. At higher redshift our model predictions are significantly higher than the UV extinction corrected data and lower by about a factor of two than the estimates derived from radio data by \\cite{2000ApJ...544..641H}.}   \\end{enumerate}"
    },
    "0403/gr-qc0403080_arXiv.txt": {
        "abstract": "In order to detect the rare astrophysical events that generate gravitational wave (GW) radiation, sufficient stability is required for GW antennas to allow long-term observation. In practice, seismic excitation is one of the most common disturbances effecting stable operation of suspended-mirror laser interferometers. A straightforward means to allow more stable operation is therefore to locate the antenna, the ``observatory'', at a ``quiet'' site. A laser interferometer gravitational wave antenna with a baseline length of 20\\,m (LISM) was developed at a site 1000\\,m underground, near Kamioka, Japan. This project was a unique demonstration of a prototype laser interferometer for gravitational wave observation located underground. The extremely stable environment is the prime motivation for going underground. In this paper, the demonstrated ultra-stable operation of the interferometer and a well-maintained antenna sensitivity are reported. ",
        "introduction": "First-generation ground-based gravitational wave antennas (LIGO-I~\\cite{ligo1, ligo2}, VIRGO~\\cite{virgo1, virgo2}, GEO\\,600~\\cite{geo1, geo2, geo3}, TAMA\\,300~\\cite{tama1, tama2, tama3}) are expected to come on-line early in this decade as a global network searching for astrophysical gravitational wave radiation. At present, some of the detectors are already operating intermittently, hoping to observe the spacetime strain of the universe.  The aim of these international projects is to directly detect gravitational radiation, faint ripples in the spacetime fabric. There are several kinds of expected astrophysical sources, including chirping gravitational waves from inspiraling compact star binaries, burst signals from supernovae explosions, and the stochastic background radiation. The expected event rate of these sources is, however, quite low even if the uncertainty of the population estimate~\\cite{eventrate1, eventrate2, eventrate3} is taken into account, so, to avoid missing these rare and faint signals, stable operation of the detector, keeping the duty cycle and the detector sensitivity high and also keeping the data quality high, are indispensable requirements for a gravitational wave observatory. In general, the technologies used in a laser interferometer are based on an ultra-high precision measurement pursuing extremely high sensitivity, so the instruments are very sensitive to almost all kinds of noise, disturbances and drifts. The noise source that most commonly disturbs stable operation of suspended-mirror laser interferometers is seismic excitation. The most promising solution for this problem is to avoid the source of these disturbances by selecting a quiet environment for a detector site.  The goal of this project (LISM, Laser Interferometer gravitational-wave Small observatory in a Mine) is to demonstrate stable operation of the laser interferometer and to obtain high quality data for searching for gravitational waves at a well-suited observatory site.  The 20-m baseline laser interferometer was originally developed for various prototyping experiments~\\cite{20m1,20m2,20m3,20m4} at the campus of The National Astronomical Observatory of Japan, in Mitaka, a suburb of Tokyo from 1991 to 1998. In 1999, it was moved to the Kamioka mine in order to perform long-term, stable observations as LISM. In this paper, the merit of going underground and the demonstrated stable operation of the antenna are reported. The results of data analysis and a GW search will appear as a separate article~\\cite{coincidence}. ",
        "conclusions": "The prototype laser interferometer gravitational wave antenna LISM with an arm length of 20\\,m was moved into a deep underground laboratory, a first for a gravitational wave observatories, and a long-term observational run was performed. The goal of this project was to confirm that the laser interferometer can operate stably and provide high quality data, making the most of the stable environment in the underground laboratory. The interferometer sensitivity achieved was $1\\x10^{-18}\\rm\\,m/\\sqrt{Hz}$ around 800\\,Hz in displacement, which corresponds to $5\\x10^{-20}\\rm/\\sqrt{Hz}$ in strain. At higher frequencies, above that floor sensitivity, LISM was quantum noise limited by shot noise. With this sensitivity spectrum, LISM can detect gravitational wave events emitted from coalescence of 1.4--1.4\\,$M_{\\odot}$ equal-mass binary systems a few\\,kpc away by matched filtering analysis with a SNR of 10. The operational duty cycle of the interferometer exceeded 99.8\\% owing to the stable, low-disturbance environment and the self-recovering automation systems. If we adopt the expected SNR of the matched filtering analysis as an index for stability of the interferometer sensitivity, this SNR duty cycle was about 90\\%. This means that the interferometer sensitivity was kept within a 3\\,dB window for 90\\% of the observational period. According to these results, we conclude that the laser interferometer gravitational wave antenna was operated stably enough for a long term observational run, and the underground environment is suitable as a gravitational wave antenna site."
    },
    "0403/astro-ph0403524_arXiv.txt": {
        "abstract": "We have mapped the ultracompact \\hii\\ region, G5.89-0.39, and its molecular surroundings with the Submillimeter Array at $2\\arcseconds.8 \\times 1\\arcseconds.8$ angular resolution in 1.3 mm continuum, \\siofive, and eight other molecular lines. We have resolved for the first time the highly energetic molecular outflow in this region. At this resolution, the outflow is definitely bipolar and appears to originate in a 1.3 mm continuum source. The continuum source peaks in the center of the \\hii\\ region. The axis of the outflow lines up with a recently discovered O5V star. ",
        "introduction": "\\label{section:int}  G5.89-0.39 (G5.89) is a shell-like ultracompact (UC) \\hii\\ region \\citep{woo89b} which was recently found to contain an O5V star via near-IR imaging \\citep{fel03}. \\citet{aco98} have used proper motion measurements of the expansion of G5.89 to determine its distance, $\\sim 2$ kpc, its size, 0.01 pc, and its dynamical age, 600 years. Associated with G5.89 is a molecular outflow \\citep{har88} with a mass of 77 \\msun\\ and an energy of $5 \\times 10^{47}$ ergs \\citep[hereafter, AWC]{aco97}. While many \\uchii\\ regions have been observationally associated with molecular outflows \\citep{sne90,she96a}, the question of whether an \\uchii\\ region can itself be or contain the source of a bipolar outflow remains in dispute.  Bipolar outflows are generally understood to be a signpost of ongoing accretion, while the presence of an \\uchii\\ region has been understood by some as a sign that accretion has ceased \\citep{gar99}. Recent observational and theoretical results have shown that the presence of an \\uchii\\ region does not necessarily shut off accretion \\citep{ket02a,ket02b}. But since massive stars and their \\uchii\\ regions tend to form in very crowded fields, high resolution imaging is necessary to associate the origin of any particular bipolar molecular outflow with an \\uchii\\ region, and not a distinct, nearby protostar. This determination is possible with and well suited to the capabilities of the new 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. For more on the SMA and its specifications see \\citet{ho03}.} (SMA) on Mauna Kea.  We have used the SMA to map the region of G5.89-0.39 in the \\siofive\\ line, known to trace outflows and this outflow in particular \\citep[AWC]{cod99}. Because of the spectral capabilities of the SMA, we were simultaneously able to map the region in nine spectral lines and broadband continuum. The aim of the project was to improve on the previous $18\\arcseconds$ resolution map in \\siofi\\ of AWC, and also to detect in dust and molecular lines any other massive protostars in the vicinity which might be sources of the outflow. ",
        "conclusions": "\\label{section:discussion}  We believe the origin of the outflow is the 1.3 mm continuum source for two reasons. First, the slightly extended 1.3 mm continuum peak lies on the projected outflow axis, halfway between the peaks of the two outflow lobes. Second, the 1.3 mm continuum source seems to show a physical connection to both outflow lobes in the form of the extensions of the continuum along the outflow axis. Apparently the dust is being swept up in both lobes of the outflow (as seen in L1157 by \\citet{gue03}), and thus it seems that the outflow is emanating from the 1.3 mm continuum source.  The question of whether the outflow actually originates {\\it inside} the \\uchii\\ region is somewhat thornier. One possibility is that the 1.3 mm continuum source lines up with the \\uchii\\ region by chance, and is a distinct high-mass protostar with its own dusty envelope, not physically connected to the \\uchii\\ region, but part of the same young cluster. Alternatively, the dust and free-free continuua could be from a single source, one dusty, hot molecular core surrounding an embedded \\uchii\\ region.  In this case, it seems the outflow would have to come from inside the \\uchii\\ region in order for both sides of the outflow to be seen in the dust continuum. We cannot distinguish between these two cases, and neither seems unlikely. However, we can say that the known O5 star is a bad candidate to be the source of the outflow since it is not equidistant from the outflow lobes.  We cannot eliminate any model chronologies by comparing timescales. The outflow has a dynamical age of 1600 years, compared to just 600 for the \\uchii\\ region. This might lead one to believe that the outflow is truly older than the \\uchii\\ region. However, any ``quenching'' of the \\uchii\\ region by infalling material would tend to make the dynamical age of the \\hii\\ region an underestimate of its true age, perhaps by a great deal \\citep{ket02b}. The gas in the outflow was probably accelerated before the \\uchii\\ region began to expand, but not necessarily before the \\uchii\\ region was created.  The central unknown about the outflow is its inclination. The fact that the two lobes hardly overlap at all on the plane of the sky and the fact that the \\siofi\\ emission is extended perpendicular to the outflow axis might lead one to believe that the outflow has a wide opening angle, and lies primarily in the plane of the sky. But these two facts could just as easily be interpreted as a highly collimated outflow, like those seen around low-mass proto-stars \\citep{bac91}, whose axis is largely along the line of sight. Including only the line-of-sight velocity, the outflow is already quite energetic, so adding a large component in the plane of the sky would make this a remarkable outflow, indeed. Some studies of regions of massive star formation have found poor collimation of outflows \\citep{ric00,rid01}. But \\citet{beu02a} found that 15 outflows with apparently low collimation were consistent with the very high collimations seen in low mass cases, with the discrepancy due to low resolution. Higher resolution interferometer studies have confirmed the presence of highly collimated outflows in some regions of high-mass star formation \\citep{beu02b}.  The velocity structure of the \\siofi\\ is intriguing and may be a clue as to the outflow orientation. In both lobes, the spectra peak at low relative velocity and tail off towards much higher relative velocity. In the case of outflows from low-mass protostars, the SiO emission is concentrated immediately behind the shock in the highest velocity molecular gas \\citep{van98}. So perhaps the fact that the low relative velocity peaks in the \\siofi\\ emission, only $\\pm 5$ \\kms, is an indication that the outflow is mainly in the plane of the sky. This outflow, however, is far more energetic than an outflow from a low-mass protostar. There could be highly excited SiO emission which follows behind the shock more closely than the \\siofi\\ emission. We plan to observe \\sioeight, and lines of several CO isotopes with the SMA allowing calculations of excitation and abundance in the outflow to see if higher excitation gas occurs at higher line of sight velocities, as we might expect for an outflow largely along the line of sight. These observations will also have higher angular resolution which should clarify the outflow opening angle.  On an equally hypothetical note, consider the fact that while some dust emission appears to come from the outflow, the free-free emission shows no preference for the direction of the outflow, and is more extended perpendicular to the outflow. Recent results have shown that massive protostars ($\\sim 10$\\msun) can have disk-like structures akin to the disks seen in low-mass protostars \\citep{zha98b,zha02}. If such a disk were to photoevaporate, as was modeled by \\citet{hol94}, that would naturally produce free-free emission which is extended perpendicular to the outflow, and possibly ring-like. However, if the outflow does not originate in the \\uchii\\ region, the more traditional limb-brightened-shell interpretation of the free-free emission may make more sense. Future work to investigate the geometry and the velocity structure of the ionized gas could be carried out with the SMA using low quantum number radio recombination lines (RRLs) of hydrogen, including \\htwonea, \\htwsixa, \\htha, and \\hthonea. These lines suffer less from pressure broadening than higher quantum number RRLs, and could be used to probe the velocity structure of the densest ionized gas in the \\uchii\\ region."
    },
    "0403/astro-ph0403238_arXiv.txt": {
        "abstract": "{ The accretion column in magnetic Cataclysmic Variables may have a not negligible Thomson optical depth. A fraction of the thermal radiation from the post--shock region may therefore be scattered -- and then polarized -- before escaping the column. Moreover, part of the thermal radiation is reflected -- and again polarized -- by the White Dwarf surface. We show that X--ray polarimetry can provide valuable, and probably unique informations on the geometry and physical parameters of the accretion column by calculating, by means of Monte Carlo simulations, the expected polarization properties of magnetic CVs as a function of the geometrical parameters (assuming a cylindrical geometry) and the Thomson optical depth of the column. We find that degrees of polarization as high as about 4\\% can be present, and  apply our calculations to the archetypal magnetic CV, AM~Herculis. ",
        "introduction": "The magnetic field in some subclasses of Cataclysmic Variables (notably in Polars and in at least a fraction of intermediate Polars; see Warner 1995 for a complete overview on CVs) is strong enough to channel the accreting matter along the field lines. For a dipolar field, this means that the accretion occurs via accreting columns on one or both the magnetic poles; the disalignement between the spin and magnetic axes results in pulsed emission.  Hard X--rays are then produced by optically thin thermal line and continuum emission in the so--called post--shock region, where temperatures can reach values as large as several tens of keV (e.g. Frank et al. 1992; Cropper et al. 1999). Thermal emission is expected to be unpolarized (see next section); however, the Thomson optical depth in the accretion column, while probably less than unity, may be not negligible. Indeed, Hellier et al. (1998) found significant broadening in the iron K$\\alpha$ line of several magnetic CVs, which they interpreted as due to Compton broadening. Thomson scattered radiation is polarized, provided that the geometry is not spherical. It is therefore to be expected that the hard X--ray emission in magnetic CVs is polarized, with the net polarization degree increasing with the Thomson optical depth of the accretion column.  In this paper we calculate, by means of Monte Carlo simulations, the polarization properties of the accretion column in magnetic CVs. In Sec. 2 the numerical code is described, while the results (including an application to the archetypal Polar, AM~Herculis) are presented in Sec.~3. Results are then summarized, and observational perspectives discussed, in Sec.~4.   \\begin{figure*}[t] \\hbox{ \\includegraphics*[width=8.0cm,height=8.0cm]{tau_0.1.ps} \\hspace{0.5cm} \\includegraphics*[width=8.0cm,height=8.0cm]{tau_0.2.ps} } \\hbox{ \\includegraphics*[width=8.0cm,height=8.0cm]{tau_0.5.ps} \\hspace{0.5cm} \\includegraphics*[width=8.0cm,height=8.0cm]{tau_1.ps} } \\caption{The degree of polarization, P,  as a function of $\\mu=cos\\theta$ for four values of $H/r_C$. In each panel, five values of $\\tau_T$ are shown: 0.05, 0.1, 0.3, 0.5 and 1 (in increasing order of the absolute value of $P$). } \\label{polittico_1} \\end{figure*}  \\begin{figure*}[t] \\hbox{ \\includegraphics*[width=8.0cm,height=8.0cm]{tau_2.ps} \\hspace{0.5cm} \\includegraphics*[width=8.0cm,height=8.0cm]{tau_3.ps} } \\hbox{ \\includegraphics*[width=8.0cm,height=8.0cm]{tau_5.ps} \\hspace{0.5cm} \\includegraphics*[width=8.0cm,height=8.0cm]{tau_10.ps} } \\caption{The same as in the previous figure, for other four values of $H/r_C$. } \\label{polittico_2} \\end{figure*}  \\begin{figure*}[t] \\hbox{ \\includegraphics*[width=8.0cm,height=8.0cm]{tau_flux_0.1.ps} \\hspace{0.5cm} \\includegraphics*[width=8.0cm,height=8.0cm]{tau_flux_0.2.ps} } \\hbox{ \\includegraphics*[width=8.0cm,height=8.0cm]{tau_flux_0.5.ps} \\hspace{0.5cm} \\includegraphics*[width=8.0cm,height=8.0cm]{tau_flux_1.ps} } \\caption{The flux (in arbitrary units)  as a function of $\\mu=cos\\theta$ for four values of $H/r_C$. In each panel, five values of $\\tau_T$ are shown: 0.05, 0.1, 0.3, 0.5 and 1 (in increasing order of anisotropy). } \\label{polittico_flux_1} \\end{figure*}  \\begin{figure*}[t] \\hbox{ \\includegraphics*[width=8.0cm,height=8.0cm]{tau_flux_2.ps} \\hspace{0.5cm} \\includegraphics*[width=8.0cm,height=8.0cm]{tau_flux_3.ps} } \\hbox{ \\includegraphics*[width=8.0cm,height=8.0cm]{tau_flux_5.ps} \\hspace{0.5cm} \\includegraphics*[width=8.0cm,height=8.0cm]{tau_flux_10.ps} } \\caption{The same as in the previous figure, for four other values of $H/r_C$. } \\label{polittico_flux_2} \\end{figure*}  \\medskip ",
        "conclusions": "We have calculated the polarization properties of the accretion column in magnetic CVs. Polarization arises because a fraction of the thermal radiation can be Thomson scattered before escaping the column. The polarization degree can be as high as $\\sim$4\\%, and depends on the angle between the magnetic field and the line--of--sight (which of course varies with the spin phase), which is often well known (e.g. Cropper 1988). It depends also on the Thomson optical depth, $\\tau_T$, and on the ratio between the radius and height of the accretion column. If one of these two parameters can be independently estimated (e.g. $\\tau_T$ from the iron line broadening, e.g. Hellier et al. 1998), than the other can in principle be deduced by polarization measurements. In particular, as the polarization degree is negative (positive) for $H/r_C$ less (greater) than one, while the polarization of the reflection component (whose importance increases with energy) is always negative, an increase (decrease) of the polarization degree with energy is expected in the former (latter) case.  In this paper we have discussed the continuum emission. However, iron lines provide a significant fraction of the total X--ray flux. Recombination lines emitted in the accretion column may suffer not only Compton scattering but also resonant scattering, so the effective optical depth in the line is larger than in the continuum. Line photons are therefore likely to be more polarized than continuum photons. On the contrary, the fluorescent neutral line emitted by the White Dwarf surface is unpolarized, at least in the line core.  Degrees of polarization of the order of one percent are within the detection capabilities of the new generation X--ray polarimeters based on the photoelectric effect (Costa et al. 2001) when coupled with large enough X--ray telescopes (Costa et al. 2003), at least for the brightest magnetic CVs in high state (when their flux can be as high as several millicrabs; for instance the 2--10 keV flux of AM~Herculis in high state is about 10$^{-10}$ erg cm$^{-2}$ s$^{-1}$, Matt et al. 2000).  Moreover, the spectral resolution should be good enough to search for different polarization degrees in the emission iron lines. Bright magnetic CVs should therefore be added to the traditional lists of targets for future polarimetric missions."
    },
    "0403/astro-ph0403712_arXiv.txt": {
        "abstract": "The Space Telescope Imaging Spectrograph has measured the absolute flux for  Vega from 0.17--1.01~$\\mu$m on the \\emph{HST} White Dwarf flux scale. These data are saturated by up to a factor of 80 overexposure but retain linearity to a  precision of 0.2\\%, because the charge bleeds along the columns and is recovered  during readout of the CCD. The S/N per pixel exceeds 1000, and the resolution $R$ is about 500. A $V$~magnitude of 0.026$\\pm$0.008 is established  for Vega; and the absolute flux level is $3.46\\times10^{-9}$~erg cm$^{-2}$ s$^{-1}$ at 5556~\\AA. In the regions of Balmer and Paschen lines, the STIS equivalent widths differ from the pioneering work of Hayes in 1985 but do agree with predictions of a Kurucz model atmosphere, so that the STIS flux distribution is preferred to that of Hayes. Over the full wavelength range, the model atmosphere calculation shows excellent agreement with the STIS flux distribution and is used to extrapolate predicted fluxes into the IR region. However, the IR fluxes are 2\\% low with respect to the standard Vega model of Cohen. \\emph{IUE} data provide the extension of the measured STIS flux distribution from 0.17 down to 0.12~$\\mu$m. The STIS relative flux calibration is based on model atmosphere calculations of pure hydrogen WDs, while the Hayes flux calibration is based on the physics of laboratory lamps and black body ovens. The agreement to 1\\% of these two independent methods for determining the relative stellar flux distributions suggests that both methods may be correct from 0.5--0.8~$\\mu$m and adds confidence to claims that the fluxes relative to 5500~\\AA\\ are determined to better than 4\\% by the pure hydrogen WD models from 0.12 to 3~$\\mu$m. ",
        "introduction": "The most commonly used flux calibration for the fundamental standard Vega ($\\alpha$~Lyr, HD~172167, HR~7001) is the compilation of Hayes (1985), while Megessier (1995) suggests an increase of 0.6\\% in the Hayes value of $3.44\\times10^{-9}$ erg cm$^{-2}$ s$^{-1}$ \\AA$^{-1}$ at 5556~\\AA. The Hayes flux for Vega is traceable to fundamental standard lamps pedigreed by NBS (now NIST), which is charged with maintaining fundamental standards of physical units. The Hayes flux compilation is based on direct comparisons of the star to calibrated lamps and black body ovens. Most ground based estimates of absolute stellar flux are traceable to the flux of Vega.  With the advent of extensive space based observations in the far-UV, a need for standard flux candles arose at wavelengths below the atmospheric cutoff. A few rocket experiments established some UV standard stars with $\\sim$10\\% precision (e.g.\\  Strongylis \\& Bohlin 1979). To establish standards with better accuracy, a technique based on model atmospheres for simple, pure hydrogen WDs was suggested by D.~Finley and J.~Holberg for calibrating the \\emph{IUE} satellite spectrophotometry, (cf.\\ Fig.~8 in Bohlin, et~al.\\ 1990). The hot WD standard stars G191B2B, GD~153, and GD~71 at $V=12$--13~mag were established by computing LTE model atmosphere flux distributions for pure hydrogen atmospheres, which were then normalized to precisely measured $V$~magnitudes from A.~Landolt (see Bohlin 2000). EUV observations below the Lyman limit demonstrated that any effects of interstellar extinction in these three stars is less than 1\\% longward of the 912~\\AA\\ Lyman limit. These three WD standards are now based on NLTE models by Ivan Hubeny (Bohlin 2003) and are complemented by three solar analog stars for \\emph{HST} calibrations extending into the IR (Bohlin, Dickinson, \\& Calzetti 2001).  The Vega observations with STIS are discussed in Section~2. Section~3 establishes the best estimate of the V band magnitude of Vega, while Section~4 presents the first direct comparison between the WD and the standard lamp methods for establishing absolute flux. A third method for defining the flux of Vega at all wavelengths is based on modeling of its stellar atmosphere, as discussed in Section~5. Model atmosphere codes are now sophisticated enough to warrant detailed comparisons between observation and theory. Because of the excellent agreement of a model (Kurucz 2003, Castelli and Kurucz 1994) with the STIS observations up to 1~$\\mu$m, the model should provide a good estimate of the flux beyond 1~$\\mu$m, as discussed in Section~6. ",
        "conclusions": "As a result of the STIS observations of Vega, a V mag of 0.026 is established for the improved bandpass function of Cohen et~al. \\ (2003b). The Kurucz \\ (2003) model of Vega agrees so well with the STIS observations that the model itself defines the standard star fluxes longward of 4200 \\AA. These improvements, combined with the more realistic bandpass function that is used to translate the V magnitudes of the WD standards into relative fluxes, have resulted in WD fluxes that are uniformly 0.5\\% fainter than found by Bohlin \\ (2000). The WD standard star model flux distributions must be updated in the CALSPEC archive, along with the new Vega spectrum. Because these files have not been updated since 00Apr4, the changes quantified by Bohlin (2003) as a function of wavelength for the switch from LTE to NLTE are also included in the new *mod*.fits files. The updates for the CALSPEC secondary flux standards are being prepared and will reflect the above changes plus a number of new STIS observations, better corrections for the changes with time and temperature, and improved CTE corrections for the CCD data. In addition, a suite of NICMOS grism calibration observations are being obtained to resolve the IR differences between the pure hydrogen WDs and the solar analogs and to extend the \\emph{HST} spectrophotometric coverage to 2.5~$\\mu$m. Unfortunately, Vega is too bright for NICMOS grism observations."
    },
    "0403/astro-ph0403462_arXiv.txt": {
        "abstract": "We have performed \\HST\\ imaging of a sample of 23 high-redshift ($1.8<z<2.75$) Active Galactic Nuclei, drawn from the \\combo\\ survey. The sample contains moderately luminous quasars ($M_B \\sim -23$). The data are part of the \\gems\\ imaging survey that provides high resolution optical images obtained with the Advanced Camera for Surveys in two bands (\\V\\ and \\z), sampling the rest-frame UV flux of the targets. To deblend the AGN images into nuclear and resolved (host galaxy) components we use a PSF subtraction technique that is strictly conservative with respect to the flux of the host galaxy. We resolve the host galaxies in both filter bands in 9 of the 23 AGN, whereas the remaining 14 objects are considered non-detections, with upper limits of less than 5~\\% of the nuclear flux. However, when we coadd the unresolved AGN images into a single high signal-to-noise composite image we find again an unambiguously resolved host galaxy. The recovered host galaxies have apparent magnitudes of $23.0<\\mathrm{\\V}<26.0$ and $22.5<\\mathrm{\\z}<24.5$ with rest-frame UV colours in the range $-0.2<(\\mathrm{\\V}-\\mathrm{\\z})_\\mathrm{obs}<2.3$. The rest-frame absolute magnitudes at 200~nm are $-20.0<M_{200~\\mathrm{nm}}<-22.2$. The photometric properties of the composite host are consistent with the individual resolved host galaxies. We find that the UV colors of all host galaxies are substantially bluer than expected from an old population of stars with formation redshift $z\\le5$, independent of the assumed metallicities. These UV colours and luminosities range up to the values found for Lyman-break galaxies (LBGs) at $z=3$. Our results suggest either a recent starburst, of e.g.\\ a few per cent of the total stellar mass and 100~Myrs before observation, with mass-fraction and age strongly degenerate, or the possibility that the detected UV emission may be due to young stars forming continuously. For the latter case we estimate star formation rates of typically $\\sim$$6\\,\\mathrm{M}_\\odot\\;\\mathrm{yr}^{-1}$ (uncorrected for internal dust attenuation), which again lies in the range of rates implied from the UV flux of LBGs. Our results agree with the recent discovery of enhanced blue stellar light in AGN hosts at lower redshifts. ",
        "introduction": "\\label{sec:intro} Around redshifts of $z\\sim 2$--3, luminous quasars were orders of magnitude more numerous than today. Although the physics of how active galactic nuclei evolve is still not understood, several links between galaxy and quasar evolution have emerged over recent years.  The observational confirmation of supermassive black holes in the nuclei of all galaxies with a substantial bulge component \\citep[e.g.][]{gebh00} makes every such galaxy a potential AGN host. The strong evolution of the AGN space density could therefore be related to the availability of accretion fuel in the host galaxies, or to the frequency of AGN triggering events.  Gravitational interaction and major or minor merging of galaxies have long been suggested as important factors in driving nuclear activity in galaxies. Confirming any of these as the dominant process has proved difficult, mainly because the morphological characteristics found for relatively nearby AGN host galaxies are so diverse. Furthermore, the properties of the hosts in the `heyday' of quasars ($z \\ga 2$) are still elusive, a consequence of the substantial observational difficulties. The contrast between the bright nuclear point sources and the surrounding galaxy increases dramatically beyond $z\\sim 1$ as a result from both surface brightness dimming and waveband shifts towards the rest frame UV.  The last years have seen numerous attempts to resolve the host galaxies of high-redshift quasars. Owing to the observational challenges of detecting distant host galaxies the observational effort for each object is large, and the observed samples have consequently been very small, of the order of $\\la 5$ per target group. While radio-loud quasars appear to be very extended and have been resolved out to $z \\sim 4$ \\citep[e.g.,][]{lehn92,carb98,hutc99,kuku01,hutc03,sanc03}, this is not the case for the large majority of radio-quiet quasars.  At high redshifts two constraints dominate observational studies of host galaxies: On one side, very good seeing conditions are required to maximize the spatial contrast of the compact nuclear source compared to the extended host galaxy. On the other, large telescope apertures are preferrential to trace faint quasar hosts to as far away from the nucleus as possible. Thus significant progress had to wait for 8m-class telescopes at very good sites with active optics systems -- with a very high light collecting power but atmospheric seeing limitations -- and for the \\HST\\ and its high space-based sensitivity, combined with unprecedented spatial resolution, but limited size that might miss light from faint outer structures of the hosts. Some host galaxies of radio-quiet quasars at $z\\simeq 2$ have now been resolved both in the near infrared \\citep{aret98,kuku01,ridg01,falo04} and in the optical domains \\citep{hutc02}, showing these objects to be moderately luminous, corresponding to present-day $L^\\star$ or slightly brighter.  However, host galaxy colours have been unavailable, precluding estimates of the mass-to-light ratio ($M/L$).  Thus, without colours the observed luminosities, and their evolution with redshift, cannot be mapped to the mass evolution if young stars contribute a major fraction of AGN host's light. This is important as several high-luminosity quasars at $z \\ga 2$ appear to be located in very UV-luminous host galaxies \\citep{lehn92,aret98,hutc02}. Also, at low redshifts there is a link between nuclear activity and enhanced global star formation in the host galaxies. \\citet{kauf03} reported that SDSS spectra of local Seyfert~2 galaxies show a significant contribution from young stellar populations, and that this trend is strongly correlated with nuclear luminosity.  In a multicolour study of low-$z$ QSO hosts \\citep{jahn04a} as well as at intermediate redshifts \\citep[see below]{sanc04a} we found that hosts of elliptical morphology can be significantly bluer than the bulk of inactive ellipticals. These results indicate that in the recent past the star formation activity in galaxies hosting an AGN may be different from normal galaxies. The details are far from understood. Clearly more information is required to investigate the relation of starformation and AGN activity, their common cause or causal order and the evolution of these properties with redshift.  The new generation of wide-field imaging mosaics obtained with the Hubble Space Telescope (\\HST), especially in the conjunction with deep AGN surveys, has opened a new observational avenue towards AGN host galaxy studies. Here we present first results on AGN within the \\gems\\ project \\citep{rix04}, the largest \\HST\\ colour mosaic to date. In the present paper we investigate the presence of rest-frame ultraviolet light in a substantial sample of $z>1.8$ AGN, all with nuclear luminosities near $M_B = -23$. In a companion paper \\citep{sanc04a} we study rest-frame colours and morphological properties of a sample of intermediate-redshift ($z \\la 1$) AGN.  The paper is organised as follows. We first describe the sample selection and properties together with a summary of the observational data (Sect.~\\ref{sec:data}). We then comment on the decomposition of the nuclear and galaxy contribution, including a brief summary of the extensive simulations that we use to estimate measurement errors (Sect.~\\ref{sec:analysis}).  In Sect.~\\ref{sec:results} we present the measured host galaxy magnitudes and describe our treatment of non-detections.  We move on to discuss the results in Sect.~\\ref{sec:discussion}, followed by our conclusions in Sect.~\\ref{sec:conclusions}.  We use $H_0=70$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_m=0.3$ and $\\Omega_\\Lambda = 0.7$ throughout this paper. All quoted magnitudes are zeropointed to the AB system with ZP$_\\mathrm{F606W}=26.493$ and ZP$_\\mathrm{F850LP}=24.843$. ",
        "conclusions": "\\label{sec:conclusions}  We performed the hitherto largest study of host galaxies properties of a complete sample of high-redshift AGN. We detected the hosts and extract colour information in 9 of the 23 AGN, and we also achieved a statistical detection of the host in the remaining 14 from a stack analysis. The UV luminosities can be interpreted in three ways: either as contribution from a passively evolving population of relatively young stars, forming typically 0.5~Gyrs ago, as a mix between a population of old (e.g.\\ $t_{age}\\sim t_\\mathrm{Hubble}$) stars and a small contribution of a recently formed young population (e.g.\\ 0.1\\%--10\\% in mass at an age of 100~Myrs or 1/10th of this for age 10~Myrs), or as an indicator of ongoing star formation at a level of $\\sim$2--15\\,$\\mathrm{M}_\\odot\\:\\mathrm{yr}^{-1}$ (uncorrected for internal dust attenuation). While the first possibility is very simplistic and appears unphysical, the UV colours actually favour the two burst interpretations; but the possibility of on-going star formation cannot be completely ruled out from our data.  In the framework of combined old and young populations, it is remarkable how similar the host galaxy colours are within the sample, and, unless different mass--age combinations conspire, hence the estimated stellar mass fractions and ages. The derived young population mass fractions and ages are also very similar to the values estimated in our companion \\gems\\ study of AGN at $z\\la 1$ \\citep{sanc04a}, where we find abnormally blue rest-frame $U-V$ colours for a substantial fraction of host galaxies, particularly the most luminous AGN in the sample. Even more, these colours and ages are in turn very close to the mean values obtained from our ground-based low-$z$ multicolour sample \\citep{jahn04a}. While the stellar population diagnostics of \\citet{kauf03} are not immediately convertible into our simple colour indicators, their impressive and highly significant results point in exactly the same direction.  While the results from all these redshift regimes are similar and point to a connection of nuclear activity and the presence of young stars, and that mass fractions of young stars are similar, we always find that a larger range of absolute masses is involved, showing as a range in UV luminosity. Here we find a variation of a factor of $\\ga$40 for a given nuclear luminosity.  Our host galaxy colours span a range that reaches the colours of Lyman Break Galaxies for a few very luminous hosts, while, as mentioned, the colours of the majority are somewhat redder than these. A comparison of optical/UV properties to the general population of high redshift galaxies would be very illuminating, but large statistical samples only become available in the near future, e.g.\\ from the \\goods\\ project.  This persistent trend to find AGN to be associated with blue stellar colours is intriguing and suggests a close connection between enhanced star formation and nuclear activity. Additional support for such a connection comes from the detection of submm CO emission in a number of extremely luminous high-redshift QSOs and radio galaxies \\citep{omon03}, although the current sensitivity of submm telescopes is insufficient to perform this test for less luminous AGN at high $z$.  While the fact that there is a relation can hardly be denied, its physical origin remains obscure. Is the enhancement of star formation a prerequisite for nuclear activity? Is it a simultaneously occurring phenomenon, caused by the same trigger? Or is it a consequence of the AGN? Galaxy merging and interaction are clearly two possible candidates to connect these two phenomena, but neither the only ones nor are the involved physics understood. Much additional data will be required, in particular those helping reliably to reconstruct the star formation history in high-redshift galaxies, before any firm conclusions can be drawn."
    },
    "0403/astro-ph0403181_arXiv.txt": {
        "abstract": "We revisit the calculation of the abundance of primordial black holes (PBHs) formed from primordial density perturbations, using a formation criterion derived by Shibata and Sasaki which refers to a metric perturbation variable rather than the usual density contrast. We implement a derivation of the PBH abundance which uses peaks theory, and compare it to the standard calculation based on a Press--Schechter-like approach. We find that the two are in reasonable agreement if the Press--Schechter threshold is in the range $\\Dc \\simeq 0.3$ to $0.5$, but advocate use of the peaks theory expression which is based on a sounder theoretical footing. ",
        "introduction": "Primordial black holes (PBHs) may have formed during the early Universe, and if so can have observational implications at the present epoch, either from effects of their Hawking evaporation or from a contribution to the present dark matter density \\cite{carr,lims}. That there is no unambiguous observational evidence of PBHs is a significant constraint on some possible types of early Universe physics.  In particular, they are the only known way of constraining the density perturbation spectrum on extremely short scales, and indeed until fairly recently provided the most powerful upper limit on the spectral index of perturbations with an exactly power-law power spectrum.  However, the abundance of PBHs formed from a given initial power spectrum remains uncertain. The traditional calculation takes the same form as the Press-Schechter calculation much used in large-scale structure studies \\cite{PS}, where the density field is smoothed on a mass scale $M$ (in this application taken to be at the time of horizon crossing), and those regions where the density contrast exceeds a threshold value $\\Dc$ are assumed to form PBHs with mass greater than $M$. However the correct value for the threshold is quite uncertain. The `standard' value of 1/3 for a radiation-dominated Universe was derived by Carr \\cite{carr} (see also Ref.~\\cite{harr}), but was probably only ever intended as an order-of-magnitude estimate. Subsequently, Niemeyer and Jedamzik~\\cite{nj} carried out numerical simulations of the collapse of isolated regions and found the threshold for PBH formation, in terms of the relative excess mass within the horizon, to be $\\Delta M / M_{{\\rm h}}= 0.7$.  However, Shibata and Sasaki~\\cite{ss} have pointed out that they formulate their initial data after horizon crossing, and hence their criterion cannot be related to the initial perturbations produced by, for instance, a period of inflation.  More recently, Shibata and Sasaki~\\cite{ss} devised a new approach to the formation of individual PBHs, seeking to find criteria on the metric perturbation rather than the density field, and in a form which can be applied to superhorizon initial perturbations. They were able to specify a criterion in terms of whether the initial central value of a particular metric perturbation variable $\\psi$ exceeds a threshold value. In this paper, we investigate the implications of this result for the abundance of PBHs formed. ",
        "conclusions": "We have provided a new calculation of the abundance of PBHs generated by primordial density perturbations. By using a metric perturbation variable rather than the density contrast, a PBH formation criterion can be applied directly to the initial perturbation spectrum. Within this formalism, we have found that the PBH mass spectrum is best computed using the theory of peaks, rather than the standard Press--Schechter-like calculation.  Given the considerable uncertainties involved, our results do not lead to any drastic revision of the PBH formation rate, but do put the calculation on a sounder theoretical footing. Our mass function can be fairly well approximated by that of the standard calculation in the region of interest ($\\Omega_{{\\rm PBH}} \\sim 10^{-20}$), if the threshold density $\\Dc$ is taken in the range $0.3$ to $0.5$. This range of threshold values is however significantly lower than the value $\\Dc \\simeq 0.7$ suggested by the simulations of Niemeyer and Jedamzik \\cite{nj}, and in fact encompasses the value $\\Dc = 1/3$ used in the earliest PBH literature. However, we advocate that anyone using our results adopts the peaks theory expression for the mass function given by Eq.~(\\ref{omegapk}).    \\vspace*{10pt}  After completion of this paper Ref.~\\cite{yok} was brought to our attention. This paper uses the constraints on the metric perturbation variable $\\psi$ from Ref.~\\cite{ss} to calculate the PBH abundance, but does not use the peaks formalism."
    },
    "0403/astro-ph0403339_arXiv.txt": {
        "abstract": "s{ A sample of radio-loud active galactic nuclei (AGN) at 2cm is studied to test the isotropic distribution of radio sources in the sky and their uniform distribution in space. The sample is complete flux-density limits of 1.5Jy for positive declinations and 2Jy for declinations between $0^{\\circ}<\\delta< -20^{\\circ}$. The active galactic nuclei sample comprises of 133 members. Application of the two-dimensional Kolmogorov-Smirnov test shows that there is no significant deviation from the isotropic distribution in the sky, while the generalised $V/V_{\\rm m}$ test shows that the space distribution of AGN is not uniform at high confidence level (99.9\\%). This is indicative of a strong positive evolution of AGN with cosmic epoch implying that AGN (or jet activity phenomena) were more populous at high redshifts. It is shown that the evolution depends strongly on luminosity: low-luminosity QSOs show a strong positive evolution, while high-luminosity counterparts do not seem to show any evolution with cosmic epoch.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403613_arXiv.txt": {
        "abstract": "A number of methods of flare prediction rely on classification of physical characteristics of an active region, in particular optical classification of sunspots, and historical rates of flaring for a given classification. However these methods largely ignore the number of flares the active region has already produced, in particular the number of small events. The past history of occurrence of flares (of all sizes) is an important indicator to future flare production. We present a Bayesian approach to flare prediction, which uses the flaring record of an active region together with phenomenological rules of flare statistics to refine an initial prediction for the occurrence of a big flare during a subsequent period of time. The initial prediction is assumed to come from one of the extant methods of flare prediction. The theory of the method is outlined, and simulations are presented to show how the refinement step of the method works in practice. ",
        "introduction": "Solar flares influence local `space weather,' and as a result there is a demand for accurate flare prediction. Unfortunately no reliable deterministic method of predicting a flare is known, and existing methods are probabilistic in nature.  A number of methods discussed in the literature are based on a commonly used white-light classification of sunspots, and the correlation between classification and flare occurrence. The McIntosh classification (McIntosh 1990) categorizes a group of sunspots into one of 60 classes, based on three parameters. Historical flare rates for each of the classifications were used by McIntosh (1990) as the basis of an `expert system' for flare prediction. The system, called Theophrastus (the associated code is called THEO), also incorporates additional information including dynamical properties of spot growth, rotation and shear, magnetic topology inferred from sunspot structure, magnetic classification, and previous flare activity. The method is apparently somewhat subjective, involving rules of thumb incorporated by a human expert. A second approach using the McIntosh classification was presented by Bornmann and Shaw (1994). In this case multiple linear regression was used to determine the effective contribution of each of the McIntosh parameters to the rate of flaring, based on historical records of flaring. Codes based on the methods of McIntosh (1990) and Bornmann and Shaw (1994) are used by the Ionospheric Prediction Service (IPS) of Australia to issue flare predictions at their Learmonth and Culgoora observatories.\\footnote{See http://www.ips.gov.au.} Recently Gallagher, Moon and Wang (2002) implemented a system using historical averages of flare numbers for McIntosh classifications to predict a rate for an active region, and then converted this to a probability of flaring in a day using the assumption of Poisson statistics. This prediction is given as part of the Big Bear Solar Observatory Active Region Monitor (ARM).\\footnote{See http://beauty.nascom.nasa.gov/arm/latest/.} Finally the US National Oceanic and Atmospheric Administration (NOAA) issues flare probability forecasts for active regions which include input from THEO.\\footnote{See http://www.sec.noaa.gov/ftpdir/latest/daypre.txt.}  A shortcoming of methods relying on correlations of flaring with active region classification based on historical records is that they ignore the important information of how many flares the active region of interest has already produced. The system of McIntosh (1990) incorporates information about previous activity, but it is unclear how objectively this is done, and the information is limited to the number of large flares already produced by the given active region. In the flare prediction literature, the tendency of a region which has produced large flares in the past to produce large flares in the future is called persistence, which is recognised as one of the most reliable predictors for large flare occurrence in 24-hour forecasts (e.g.\\ Neidig, Weiborg, \\& Seagraves 1989). In this paper we argue that the history of occurrence of all flares (large and small) observed in a given active region is an important indicator as to how the region will flare in the future, and should be used in any prediction. A related criticism of methods based on classification and historical records is that a given classification may embrace active regions with a variety of flaring rates. If an active region has a flaring rate differing from the average historical rate for its class then the predictions will be in error.  Studies of solar flare statistics provide simple phenomenological rules describing flare occurrence. It is well known that flares follow a power-law size distribution, where by size we mean e.g.\\ peak flux in soft X-ray. More formally the flare frequency-size distribution $N(S)$ (i.e.\\ the number of events per unit size $S$ and per unit time) may be written \\begin{equation}\\label{eq:pldist} N(S)=AS^{-\\gamma} \\end{equation} where $A$ and $\\gamma$ are constants. The exact power-law index $\\gamma$ depends on the choice of the quantity $S$, but typically it is found to be in the range 1.5 to 2 (e.g.\\ Crosby, Aschwanden, \\& Dennis 1992). The power law index $\\gamma$ appears to be the same in different active regions~\\cite{whe00}, although there is some evidence that it varies with the solar cycle~\\cite{bai93}. A second simple rule concerns the way flares occur in time. Studies of the rate of occurrence of soft X-ray flares in individual active regions suggest that events occur as a Poisson process in time (e.g.\\ Moon et al.\\ 2001), although many active regions exhibit changes in the mean rate of events (Wheatland 2001).  In this paper we show how the observed record of flaring in an active region may be used together with the phenomenological rules of flare statistics to objectively refine an initial flare prediction. The initial prediction may be based on the McIntosh classification, or may come from any other prediction method which does not consider the flare data. The new method is envisaged to work as follows. When an active region appears at the east limb of the Sun, the best guess as to its future flare productivity comes from one of the conventional prediction methods. However, as the active region produces flares, the observed flare statistics are used to adjust the prediction for future flaring. After many flares have been observed, the prediction for future flaring may be dominated by the contribution from the observed data. This process --- refining a probability estimate based on new data --- is naturally performed using Bayes's theorem (e.g.\\ Sivia 1996; Jaynes 2003).  The layout of the paper is as follows. In \\S\\,2 a simple approach to flare prediction using only the past record of flaring from an active region [previously presented in Wheatland (2001)] is reiterated. In \\S\\,3 the new method of prediction, combining existing methods and information from observed flare statistics, is described. In \\S\\,4 simulations are presented showing how the method uses the observed flaring record, and in \\S\\,5 the results are discussed. ",
        "conclusions": "Existing methods of solar flare prediction do not make complete use of an important source of information: the time history of flares already observed in the active region of interest, in particular frequently occurring small events. A new method for flare prediction is presented which exploits the observed history of flaring from an active region to improve an initial prediction, which e.g.\\ may come from one of the existing methods. To make the example concrete we may think of the initial prediction coming from from the McIntosh sunspot classification, which is a common basis for prediction. This background information provides an initial estimate for the expected flaring rate through a prior distribution $\\Lambda_{\\rm MC}(\\lambda_1)$, which represents the probability that the flaring rate above a (small) size $S_1$ is $\\lambda_1$, given historical rates of occurrence of flares for the given McIntosh class. Bayes's theorem is then used to estimate the probability $\\epsilon$ of observing a large flare (above size $S_2$) in a given period of time, based on this prior information and on the sequence of flares already produced by the active region, and assuming simple phenomenological rules describing the occurrence of flares. In this paper the basic theory behind the inference of $\\epsilon$ based on observed data is presented. The inclusion of background information [i.e.\\ the construction of the priors $\\Lambda_{\\rm MC}(\\lambda_1)$] is yet to be done.  The method relies on event sizes following the phenomenological law~(\\ref{eq:pldist}). Some studies of very small extreme ultraviolet events (`nanoflares') suggest that their thermal energies follow a steeper distribution than energies of large events (e.g.\\ Krucker and Benz 1998; Parnell and Jupp 2000), although this remains controversial (e.g.\\ Aschwanden and Parnell 2002). From the point of view of the prediction method presented here, the uncertainty over the low-size end of the distribution is irrelevant provided events significantly larger than nanoflares are used. In any case the observed distributions from many active regions may be examined as a check on Equation~(\\ref{eq:pldist}). A related point is that the distribution~(\\ref{eq:pldist}) requires a cutoff at large sizes on energetics grounds, and neglect of this cutoff will lead to the number of large flares being overestimated. A cutoff will be incorporated before the method is applied to real data.  The choice of the quantity $S$ has not been addressed, although a good choice is likely to be important to the method. Most flare forecasting deals with soft X-ray events, in particular prediction of GOES (Geostationary Observational Environmental Satellite) M and X class events (events with peak fluxes greater than $10^{-5}$W/m$^2$ and $10^{-4}$W/m$^2$ respectively in the 1-8 Angstrom band observed by the satellites). A practical motivation for this is that flare soft X-ray emission causes disturbances of the ionosphere which affect shortwave radio communication, and there is a need to predict these occurrences. A disadvantage of using GOES events is that they are not ideal for flare statistics e.g.\\ because of problems with event selection due to the large background in soft X-ray (see Wheatland 2001).  A number of other issues also need to be considered before the method is implemented with real data. A point neglected so far is that active regions evolve, so that predictions based on the traditional methods also change with time. For example, an active region evolves through McIntosh classifications (e.g.\\ Bornmann, Kalmbach, Kulhanek, and Casale 1990). Changes in background information such as this should be incorporated through changes in the prior, and this question will be considered in more detail in future work. A related point concerns the construction of the prior distributions for rate. It is likely that the McIntosh classification will be used, although other possibilities will be considered. The problem is then to determine the probability of a given McIntosh class having a given rate, based on observed flaring sequences in the historical record for active regions of that class. The details of this calculation will be addressed in future work.  Finally, as with all methods of forecasting, it is essential to test the reliability of the method. It is straightforward to compare, after the fact, the number of predicted and the number of observed events for a large sample of active regions. The method presented here will be implemented and tested in this way, and the results compared with existing methods of prediction."
    },
    "0403/astro-ph0403425_arXiv.txt": {
        "abstract": "We propose a dynamical mechanism for capturing stars around a massive black hole (MBH), which is based on the accumulation of a very dense cluster of compact stellar remnants near the MBH. This study is motivated by the presence of $\\sim\\!10$ young massive stars ($\\Ms\\!\\sim\\!3$--$15\\,\\Mo$, spectral types $\\sim$B9V--O8V) less than $0.04$ pc from the MBH in the Galactic center (GC). Their existence in the extreme environment so close to an MBH is a challenge for theories of star formation and stellar dynamics. We show that young stars, which formed far from the MBH and were then scattered into eccentric orbits, repeatedly cross a cluster of stellar black holes (SBHs), where they may undergo rare direct three-body exchanges with an MBH-SBH {}``binary''. The interaction between two objects of comparable mass ejects the SBH and captures the star on a tight orbit around the MBH. Such captures can naturally explain some trends observed in the orbits of the young stars. We derive the capture cross-section, validate it by Monte Carlo simulations, and calculate the number of captured stars in the GC using the currently uncertain estimates of the numbers of SBHs in the inner 0.04 pc and of young stars in the inner few parsecs of the GC. We find that under favorable conditions three-body exchange can account for $\\sim\\!25$\\% of the observed stars, mostly at the fainter end of the observed range. We discuss additional effects that possibly increase the capture efficiency. Future detections of the dark mass around the MBH and deeper surveys of the central parsecs will establish whether or not there are enough SBHs and young stars there for exchange captures to singly account for the central young stars. We estimate that there are also $\\sim\\!35$ lower mass stars ($\\Ms\\!\\sim\\!1$--$3\\,\\Mo$, $\\sim$G2V--A0V) in the inner $0.04$ pc similarly captured by exchanges with neutron stars (NSs). Ongoing replacement of compact remnants by main-sequence stars (SBHs by NS progenitors, NSs by white dwarf progenitors) may regulate the accumulation of compact remnants near the MBH. ",
        "introduction": "Deep near-infrared photometric (Krabbe et al. \\citeyear{Kra95}; Genzel et al. \\citeyear{Gen03a}), spectroscopic (Genzel et al. \\citeyear{Gen97}; Eckart, Ott \\& Genzel \\citeyear{Eck99}; Figer et al. \\citeyear{Fig00}; Gezari et al. \\citeyear{Gez02}; Ghez et al. \\citeyear{Ghe03a}) and astrometric (Ghez et al. \\citeyear{Ghe03b}; Sch\\\"odel et al. \\citeyear{Sch03}) observations of the dense stellar cusp around the massive black hole (MBH) in the Galactic center (GC) reveal a centrally concentrated distribution of young massive stars, $\\Ns\\sim\\!10$ stars within the central $\\rs\\sim\\!0.04$ pc, $\\sim\\!40$ within $\\sim\\!0.1$ pc, spanning a mass range of $\\sim\\!3$--$15\\,\\Mo$ (spectral types $\\sim$B9V--O8V) with a median mass of $\\Ms\\!\\sim\\!10\\,\\Mo$, radius of $\\Rs\\!\\sim\\!4.5\\,\\Ro$ and main sequence life time of $\\ts\\!\\sim\\!3\\!\\times\\!10^{7}$ yr% \\footnote{These are rough inferences. The separation of the young stars from the old population is uncertain since at present only the brightest star has been spectroscopically identified (Ghez et al. \\citeyear{Ghe03a}).% }.  Orbital solutions obtained for eight of the stars (Ghez et al. \\citeyear{Ghe03b}; Sch\\\"odel et al. \\citeyear{Sch03} ) tentatively suggest some trends in their orbital properties: a lower bound on the apoapse of $\\sim\\!0.01$ pc (Ghez et al. \\citeyear{Ghe03b}) and higher than random orbital eccentricities (Sch\\\"odel et al. \\citeyear{Sch03}).  None of the solutions proposed so far for the puzzle of the young stars (Genzel et al. \\citeyear{Gen03a}; Ghez et al. \\citeyear{Ghe03a}) are satisfactory. These fall into three categories: exotic modes of star formation near the MBH; rejuvenation of old stars in the local population; or dynamic migration from farther out, where stars can form. Even if shock cooling of molecular gas by cloud-cloud collisions could trigger star formation near the MBH (Morris \\citeyear{Mor93}), molecular clouds would either form stars or be tidally disrupted well outside of the inner 0.1 pc (Vollmer \\& Duschl \\citeyear{Vol01}). Growth and rejuvenation by mergers (Genzel et al. \\citeyear{Gen03a}) are not expected to be efficient in high velocity collisions near an MBH. Tidal heating by the MBH requires that the stars approach the MBH much closer than they are observed to do (Alexander \\& Morris \\citeyear{Ale03a}). The young stars are too short-lived to have formed far from the MBH and then migrated inward by mass segregation or dynamical friction. The migration can be accelerated if the stars are associated with a massive {}``anchor'': an extremely dense young cluster (Portegies Zwart, McMillan \\& Gerhard \\citeyear{Por03}; Kim \\& Morris \\citeyear{Kim03}), a very massive binary companion (Gould \\& Quillen \\citeyear{Gou03}), or a $10^{3}$--$10^{4}\\,\\Mo$ black hole (Hansen \\& Milosavljevi\\'c \\citeyear{Han03}). However, these scenarios must assume the existence of very rare, or even hypothetical objects, or else they cannot bring the stars close enough to the MBH. Such processes may possibly explain the separate population of very massive and luminous {}``He stars'' $0.1$--$0.5$ pc from the MBH (Krabbe et al. \\citeyear{Kra95}), which we do not attempt to model here.  Our model is based on the fact that $10^{4}$--$10^{5}$ stellar black holes (SBHs) of mass $\\sim\\!7$--$10\\,\\Mo$ are estimated to exist within $\\sim\\!1$ pc of the MBH in the GC, where they have been accumulating by dynamical friction over the lifetime of the Galaxy ($t_{H}\\!\\sim\\!10\\,\\mathrm{Gyr}$ ) from a {}``collection basin'' $\\sim\\!10$ pc wide (Morris \\citeyear{Mor93}; Miralda-Escud\\'{e} \\& Gould \\citeyear{Mir00}). Numerical simulations of the evolution of the GC (Freitag \\citeyear{Fre03}) confirm that the SBHs sink to the center on a short timescale of a few gigayears, settle into a centrally concentrated distribution where the enclosed number scales as $\\ns(<\\! r)\\!\\propto r^{5/4}$ (Bahcall \\& Wolf \\citeyear{Bah77}), and dominate the stellar mass there. ",
        "conclusions": "The efficiency of a {}``billiard ball'' recoil depends strongly on the mass ratio of the colliding objects. The dynamical evolution of a stellar system around an MBH naturally provides a dense concentration of targets whose masses are well matched for stopping and capturing unbound young stars like those observed very near the MBH in the GC. The attempt to predict the number of captured stars is limited by the uncertainty in the mass and number distribution of the SBHs and in the number and orbital properties of young stars (spectral types $\\sim$B9V--O8V) in the inner few parsecs of the GC. We show that under favorable conditions, capture can account for $\\sim\\!25$\\% of the observed young stars.  Additional effects that were not taken into account here may increase the capture efficiency. We considered only capture by a single direct exchange in the point-mass approximation. However, a few weaker interactions may also lead gradually to capture during the star's lifetime. Further study is needed to estimate the contribution of multiple scatterings to the capture cross section. Tidal energy extraction in captures near the tidal limit also increases the capture cross section there. A possible outcome of internal mixing by a strong tidal interaction is an extended main-sequence lifetime and higher luminosity (Maeder \\& Meynet \\citeyear{Mae00}). This would further increase the number of captured luminous stars (Eq. \\ref{eq:iNs}).  In the context of direct exchange capture, the stars with the smallest apoapse offer an opportunity to study the long term effects of a strong tidal interaction. Interestingly, S2, the star with the smallest apoapse, is also the brightest (Ghez et al. \\citeyear{Ghe03b}; Sch\\\"odel et al. \\citeyear{Sch03}). The orbital eccentricity distribution reflects the mass ratio between the star and the SBH. Thus, a spectral determination of the masses of the captured stars together with a statistical analysis of their eccentricities can probe the poorly known mass function of SBHs. The total number of captured stars over the lifetime of the Galaxy (assuming steady state), $\\sim\\!2\\Ns t_{H}/\\ts\\!\\sim\\!3\\!\\times\\!10^{4}$ (for $40$ stars of $10\\,\\Mo$ in 0.1 pc), is of the same order as the number of SBHs in the inner parsec. Thus, if the dynamical friction timescale at $\\left\\langle a_{0}\\right\\rangle $ (where the SBHs are ejected to) is not much smaller than the age of the Galaxy, the continual replacement of SBHs by NS progenitors and of NSs by white dwarf progenitors (Fryer \\& Kalogera \\citeyear{Fry01}) may regulate the build-up of the dense cusp of compact stellar remnants."
    },
    "0403/astro-ph0403080_arXiv.txt": {
        "abstract": "{We present photometry and spectroscopy of the afterglow of \\grb. VLT spectra of the afterglow show damped \\lya\\ (DLA) absorption and low- and high-ionization lines at a redshift $z$=3.3718$\\pm$0.0005. The inferred neutral hydrogen column density, log N(\\ion{H}{i})=21.90$\\pm$0.07, is larger than any (GRB- or QSO-) DLA \\ion{H}{i} column density inferred directly from \\lya\\ in absorption. From the afterglow photometry, we derive a conservative upper limit to the host-galaxy extinction: A$_{\\rm V}$$<$0.5 mag. The iron abundance is [Fe/H]=--1.47$\\pm$0.11, while the metallicity of the gas as measured from sulphur is [S/H]=--1.26$\\pm$0.20. We derive an upper limit on the H$_2$ molecular fraction of 2$N$(H$_2$)/(2$N$(H$_2$)+$N$(\\ion{H}{i}))$\\lsim$10$^{-6}$.  In the \\lya\\ trough, a \\lya\\ emission line is detected, which corresponds to a star-formation rate (not corrected for dust extinction) of roughly 1 M\\subsun\\ yr$^{-1}$.  All these results are consistent with the host galaxy of \\grb\\ consisting of a low metallicity gas with a low dust content.  We detect fine-structure lines of silicon, \\ion{Si}{ii}*, which have never been clearly detected in QSO-DLAs; this suggests that these lines are produced in the vicinity of the GRB explosion site. Under the assumption that these fine-structure levels are populated by particle collisions, we estimate the \\ion{H}{i} volume density to be n$_{\\ion{H}{i}}=10^2-10^4$ cm$^{-3}$. HST/ACS imaging 4 months after the burst shows an extended AB(F606W)=28.0$\\pm$0.3 mag object at a distance of 0\\farcs14 (1kpc) from the early afterglow location, which presumably is the host galaxy of \\grb. ",
        "introduction": "\\label{sec:introduction}  Damped \\lya\\ (DLA) absorbers, conventionally detected in Quasi-Stellar Object (QSO) spectra, are absorption-line systems that have a column density of N(\\ion{H}{i}) $\\ge$ 2$\\times 10^{20}$ atoms cm$^{-2}$, as determined from the damping wings of the \\lya\\ line \\citep[e.g.][]{1986ApJS...61..249W,1989ApJ...344..567T}.  DLA systems are believed to contain the bulk of the neutral hydrogen at high redshift and to be the gas reservoir from which the stars at the present epoch are produced \\citep[e.g.][]{1987txra.symp..309W,1991ApJS...77....1L}. Numerous high-resolution spectroscopic studies have extracted detailed information about the metallicity \\citep[e.g.][]{prochaska}, the kinematics \\citep{1997ApJ...487...73P,1998A&A...337...51L}, and the dust and H$_2$ contents \\citep{2000A&A...364L..26P,ledoux} of DLA systems as a function of redshift.  Despite intensive searches, only a handful of DLA counterparts have been detected so far \\citep[see][]{2002ApJ...574...51M}; linking DLA systems with galaxy types has therefore proven difficult: some advocate large, disk-forming galaxies \\citep[e.g.][]{1995ApJ...454..698W,1997ApJ...487...73P}, others suggest they are faint, gas-rich dwarfs \\citep{1998ApJ...495..647H}.  Gamma-ray burst (GRB) afterglows are, just as QSOs, bright and distant sources. For instance, the spectacular GRB\\,990123 was detected at the 9$^{\\rm th}$ visual magnitude \\citep{1999Natur.398..400A} while it was located at $z$=1.6 \\citep{1999Natur.398..389K,1999Sci...283.2075A}. However, the afterglow brightness in general fades very rapidly in time (roughly flux $\\propto$ time$^{-1}$). The current afterglow redshifts range from $z$=0.169 \\citep{2003GCN..2020....1G} to $z$=4.5 \\citep{2000A&A...364L..54A}.  Moreover, GRBs are associated with massive-star formation: the discovery of a supernova (SN) spectrum similar to that of SN1998\\,bw \\citep{1998Natur.395..670G} superimposed on the GRB\\,030329 afterglow spectrum \\citep{stanek,hjorth030329} provided strong evidence that at least some of the long-duration ($\\gsim$ 2s) GRBs are caused by the core collapse of massive stars \\citep{1993ApJ...405..273W,1999ApJ...524..262M}.  The discovery of a damped \\lya\\ (DLA) absorption line at the burst redshift in the spectra of several GRB afterglows \\citep{2001A&A...370..909J,fynbo926spectrum,hjorth020124} is consistent with the massive-star progenitor scenario: they indicate a high neutral hydrogen column density origin in the host galaxy, presumably a star-forming region.  However, the signal-to-noise ratio at the location of the DLA absorption line in the spectra is fairly low in these cases, much lower than for typical QSO-DLAs.  We here present afterglow spectra of the high-redshift \\grb, which unambiguously demonstrate a GRB-DLA, with a column density exceeding that of any (QSO- or GRB-) DLA measured so far using \\lya\\ in absorption. These spectra allow us to measure the metallicity of the host environment and obtain an upper limit on the molecular fraction, i.e.  measurements that are routinely performed for QSO-DLAs, but that are still unique for GRB hosts.  Although the GRB-DLA sample is still very small, we compare them with QSO-DLAs in two aspects: their \\ion{H}{i} column density and their metallicity.  \\grb\\ was detected on 23 March 2003 at 21:57 UT by HETE \\citep{2003GCN..1956....1G} with a fluence of 1.1$\\times$10$^{-6}$ ergs cm$^{-2}$ (30-400 keV), and a duration of 26 seconds. Following the HETE localization, the optical counterpart was discovered 7.6 hours after the burst at R.A. 11\\hour06\\min09\\fs38, Decl. --21\\degree46\\arcmin13\\farcs3 (J2000) by \\citet{2003GCN..1949....1G}, with a brightness of R=18.7. Our team reported a preliminary redshift of $z$=3.372 \\citep{2003GCN..1953....1V}, which is currently the third highest redshift for a GRB \\citep{1998Natur.393...35K,2000A&A...364L..54A}.  This paper is organized as follows: in Sect. \\ref{sec:observations}, we describe the data reduction of both the spectroscopic and imaging observations. In Sect. \\ref{sec:photometry}, we present the light curves and infer an upper limit on the rest-frame optical extinction. We measure the equivalent widths of the absorption lines and determine the burst redshift in Sect. \\ref{sec:redshift}. An \\ion{H}{i} column density model is fitted to the damped \\lya\\ line in Sect. \\ref{sec:column}, and we analyze the spectra in more detail in Sect. \\ref{sec:metallicityandh2} to derive the metallicity and an upper limit on the molecular hydrogen (H$_2$) fraction. The detection of \\lya\\ in emission is presented in Sect. \\ref{sec:lya}, and we report on the detection of the probable host galaxy of \\grb\\ in HST/ACS imaging data in Sect. \\ref{sec:hst}. In the final section, we close with a general discussion of all these results.  \\begin{table}[tp] \\centering \\caption[]{Log of UT4/FORS2 spectroscopic observations}\\label{tab:spectroscopy} \\null\\vspace{-1.0cm} $$ \\begin{array}{cllccc} \\hline \\noalign{\\smallskip} \\rm UT \\, date & \\rm grism(filter) & \\rm coverage & \\rm \\lambda / \\Delta\\lambda & \\rm exptime & \\rm seeing \\\\ \\rm March\\,03 & & \\rm (nm) & & \\rm (min) & (\\arcsec) \\\\ \\hline \\rm 25.050 & \\rm 300V        & 330-660  & 440  & 3\\times10  & 1.1 \\\\ \\rm 25.077 & \\rm 300I(OG590)  & 600-1100 & 660  & 3\\times10  & 0.7 \\\\ \\rm 26.213 & \\rm 1400V       & 456-586  & 2100 & 4\\times30  & 0.8 \\\\ \\rm 26.306 & \\rm 1200R(GG435) & 575-731  & 2140 & 4\\times30  & 0.9 \\\\ \\hline \\end{array} $$ \\end{table}   \\begin{table}[tbp] \\centering \\caption[]{Log of imaging observations}\\label{tab:imaging} \\null\\vspace{-1.35cm} $$ \\begin{array}{rccccc} \\hline \\noalign{\\smallskip} \\rm UT \\, date & \\rm magnitude ^{\\mathrm{a}} & \\rm filter & \\rm exptime & \\rm seeing & \\rm tel./instr.^{\\mathrm{b}} \\\\ \\rm (2003) & & & \\rm (min) & (\\arcsec) & \\\\ \\hline \\rm Mar\\, 24.302 & 20.39 \\pm 0.06 & \\rm  B &  12 &  2.4 & \\rm USNO\\,1m  \\\\ \\rm Mar\\, 24.310 & 19.68 \\pm 0.05 & \\rm  V &   8 &  2.5 & \\rm USNO\\,1m  \\\\ \\rm Mar\\, 24.316 & 18.75 \\pm 0.03 & \\rm  R &   8 &  2.6 & \\rm USNO\\,1m  \\\\ \\rm Mar\\, 24.323 & 18.20 \\pm 0.04 & \\rm  I &   8 &  2.4 & \\rm USNO\\,1m  \\\\ \\rm Mar\\, 24.992 & 20.64 \\pm 0.26 & \\rm  R &8.33 &  1.1 & \\rm CAHA\\,2.2m  \\\\ \\rm Mar\\, 25.027 & 20.56 \\pm 0.08 & \\rm  R &  54 &  2.0 & \\rm Danish  \\\\ \\rm Mar\\, 25.030 & 21.44 \\pm 0.03 & \\rm  V &   1 &  1.1 & \\rm FORS2  \\\\ \\rm Mar\\, 25.033 & 21.45 \\pm 0.03 & \\rm  V &   1 &  1.0 & \\rm FORS2  \\\\ \\rm Mar\\, 25.091 & 22.29 \\pm 0.04 & \\rm  B &   1 &  0.9 & \\rm FORS2  \\\\ \\rm Mar\\, 25.092 & 21.42 \\pm 0.02 & \\rm  V &   1 &  0.8 & \\rm FORS2  \\\\ \\rm Mar\\, 25.094 & 20.51 \\pm 0.02 & \\rm  R &   1 &  0.8 & \\rm FORS2  \\\\ \\rm Mar\\, 25.095 & 20.02 \\pm 0.02 & \\rm  I &   1 &  0.8 & \\rm FORS2  \\\\ \\rm Mar\\, 25.099 & 20.04 \\pm 0.10 & \\rm  I &  15 &  2.0 & \\rm Danish  \\\\ \\rm Mar\\, 25.129 & 20.57 \\pm 0.03 & \\rm  R &  25 &  2.0 & \\rm Danish  \\\\ \\rm Mar\\, 25.150 & 21.55 \\pm 0.12 & \\rm  V &  15 &  1.8 & \\rm Danish  \\\\ \\rm Mar\\, 25.250 & 20.83 \\pm 0.13 & \\rm  R & 130 &  3.0 & \\rm SARA   \\\\ \\rm Mar\\, 26.023 & 22.17 \\pm 0.08 & \\rm  V &  25 &  1.2 & \\rm Danish  \\\\ \\rm Mar\\, 26.041 & 21.27 \\pm 0.05 & \\rm  R &  25 &  1.0 & \\rm Danish  \\\\ \\rm Mar\\, 26.064 & 20.82 \\pm 0.08 & \\rm  I &  25 &  1.0 & \\rm Danish  \\\\ \\rm Mar\\, 26.107 & 19.86 \\pm 0.04 & \\rm  J &   6 &  0.7 & \\rm NACO  \\\\ \\rm Mar\\, 26.120 & 17.93 \\pm 0.07 & \\rm  K &2.25 &  0.7 & \\rm NACO  \\\\ \\rm Mar\\, 26.162 & 22.12 \\pm 0.02 & \\rm  V &   1 &  0.6 & \\rm FORS2  \\\\ \\rm Mar\\, 26.164 & 22.09 \\pm 0.03 & \\rm  V &   1 &  0.6 & \\rm FORS2  \\\\ \\rm Mar\\, 26.257 & 22.22 \\pm 0.03 & \\rm  V &   1 &  0.7 & \\rm FORS2  \\\\ \\rm Mar\\, 26.259 & 22.18 \\pm 0.04 & \\rm  V &   1 &  0.8 & \\rm FORS2  \\\\ \\rm Mar\\, 26.267 & 21.50 \\pm 0.07 & \\rm  R & 160 &  2.2 & \\rm USNO\\,1m  \\\\ \\rm Mar\\, 26.306 & 18.13 \\pm 0.18 & \\rm  K &3.35 &  0.5 & \\rm UKIRT  \\\\ \\rm Mar\\, 26.351 & 22.32 \\pm 0.04 & \\rm  V &   3 &  0.8 & \\rm FORS2  \\\\ \\rm Mar\\, 26.351 & 19.38 \\pm 0.05 & \\rm  H &3.35 &  0.6 & \\rm UKIRT  \\\\ \\rm Mar\\, 26.353 & 21.43 \\pm 0.02 & \\rm  R &   3 &  0.8 & \\rm FORS2  \\\\ \\rm Mar\\, 26.356 & 20.90 \\pm 0.03 & \\rm  I &   3 &  0.8 & \\rm FORS2  \\\\ \\rm Mar\\, 26.379 & 20.09 \\pm 0.03 & \\rm  J &3.35 &  0.7 & \\rm UKIRT  \\\\ \\rm Mar\\, 27.041 & 22.01 \\pm 0.14 & \\rm  R &  30 &  1.0 & \\rm Danish \\\\ \\rm Mar\\, 27.156 & 22.94 \\pm 0.09 & \\rm  V &  24 &  0.7 & \\rm Gemini\\,S  \\\\ \\rm Mar\\, 27.176 & 22.15 \\pm 0.15 & \\rm  R &22.5 &  0.6 & \\rm Gemini\\,S  \\\\ \\rm Mar\\, 27.196 & 21.57 \\pm 0.09 & \\rm  I &22.5 &  0.6 & \\rm Gemini\\,S  \\\\ \\rm Mar\\, 28.159 & 22.36 \\pm 0.09 & \\rm  R &  36 &  1.2 & \\rm Danish \\\\ \\rm Mar\\, 28.398 & 19.06 \\pm 0.20 & \\rm  K &   7 &  0.5 & \\rm UKIRT  \\\\ \\rm Mar\\, 28.455 & 20.34 \\pm 0.08 & \\rm  H &   7 &  0.7 & \\rm UKIRT  \\\\ \\rm Apr\\,  3.282 & 24.30 \\pm 0.13 & \\rm  R &  88 &  1.3 & \\rm Gemini\\,S  \\\\ \\rm Jul\\,  5.979 & > 25.0 \\,(3\\sigma) & \\rm  I &  15 &  0.7 & \\rm FORS2 \\\\ \\rm Jul\\,  5.993 & > 25.6 \\,(3\\sigma) & \\rm  R &  15 &  0.8 & \\rm FORS2 \\\\ \\rm Jul\\, 20.958 & 28.0  \\pm 0.3  & \\rm V &  32 &      &  \\rm HST \\\\ \\hline \\end{array} $$ \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] The magnitudes have {\\it not} been corrected for Galactic extinction, and the errors do {\\it not} include the uncertainty in the absolute calibration (see text). \\item[$^{\\mathrm{b}}$] Combinations of telescopes and instruments: USNO Flagstaff Station 1.0m with 2k$\\times$2k Tek CCD, 0$\\farcs$68/pixel; Calar Alto 2.2m and CAFOS with 1k$\\times$1k SITe CCD, 0$\\farcs$29/pixel; ESO/Danish 1.54m and DFOSC with 2k$\\times$4k EEV CCD, binned to 0$\\farcs$78/pixel; Yepun and FORS2 with two MIT CCDs of 4k$\\times$2k, binned to 0$\\farcs$25/pixel; Yepun and NACO with the Aladdin InSb 1k$\\times$1k detector and S54 camera, 0$\\farcs$054/pixel; SARA 0.9m and Apogee Ap7 CCD camera with 512$\\times$512 array, 0$\\farcs$7/pixel; UKIRT and UFTI with a 1k$\\times$1k HgCdTe array, 0$\\farcs$091/pixel; Gemini South and acquisition camera with 1K$\\times$1K EEV CCD, binned to 0$\\farcs$23/pixel; HST/ACS and WFC detector with two 2k$\\times$4k SITe CCDs, 0$\\farcs$05/pixel. \\end{list} \\end{table} ",
        "conclusions": "\\label{sec:discussion}  Our observations show that \\grb\\ occurred behind a very high \\ion{H}{i} column density, in an environment (immediate and host-galaxy combined) having a low molecular hydrogen fraction (f$\\lsim$10$^{-6}$), a low metallicity ([S/H]=--1.26$\\pm$0.20) and a low dust content ($\\kappa$=0.02). For the DLA host of GRB\\,020124, \\citet{hjorth020124} also find evidence for a large \\ion{H}{i} column density with a low reddening.  The inferred low dust content may be interpreted as a selection bias: GRBs that would occur in a dusty host galaxy would be harder to detect because they would then be fainter. However, in this case one would expect to observe many GRB afterglows with considerable extinction in the optical, for which there is no clear evidence \\citep{2001ApJ...549L.209G,2002MNRAS.330..583L}. In apparent contradiction with this is the detection of several host galaxies in the radio and sub-mm \\citep[e.g.][]{frail222,2003MNRAS.338....1B,2003ApJ...588...99B}, suggesting that at least some GRB hosts are dusty, as expected when most of the star formation in the universe occurs in submm-bright galaxies \\citep[see][]{2002MNRAS.329..465R}.  Dust destruction \\citep[e.g.][]{2001ApJ...563..597F,2002ApJ...569..780D,2003ApJ...585..775P}, which has been proposed to account for the apparent discrepancy between the low optical extinction and large (X-ray) gas column densities \\citep{2001ApJ...549L.209G}, could play a role, but is not required by our data: the reduced metallicity and hence the low dust-to-gas ratio in the host of \\grb\\ is sufficient to explain the combination of a large \\ion{H}{i} column density with a low optical extinction \\citep[see also][]{hjorth020124}.  \\lya\\ in emission is detected for \\grb, and we inferred a star-formation rate of about 1 M\\subsun\\ yr$^{-1}$, which is in good agreement with the SFR value that we obtained from the UV continuum emission of the host galaxy. \\citet{fynbolya} note that \\lya\\ is commonly observed in all GRB host galaxies at high redshift for which it could be detected.  In contrast, only 25\\% of the Lyman-break galaxies are \\lya\\ emitters with an equivalent width EW $>$ 20 \\AA\\ \\citep{2003ApJ...588...65S}.  \\citet{fynbolya} suggest that this difference is due to GRB hosts having a low metallicity and a low dust content, consistent with our observations of \\grb\\ and with those of GRB\\,020124 \\citep{hjorth020124}.  We note that QSO-DLAs also have a low metallicity and a low dust content, but they rarely show \\lya\\ in emission. However, since most galaxy counterparts of QSO-DLAs are very faint, \\lya\\ in emission is not expected to be detected in most cases with the current detection limits \\citep[see][]{1999MNRAS.305..849F}. The low-dust inference for GRB\\,020124 \\citep{hjorth020124} and \\grb\\ is different from the results of \\citet{2003ApJ...585..638S}, who find evidence for a high dust content in three GRB host galaxies.  From the fine-structure lines \\ion{Si}{ii}* \\lam\\lam 1309,1533 , we estimated the \\ion{H}{i} volume density of the gas producing this absorption: n$_{\\ion{H}{i}}=10^2-10^4$ cm$^{-3}$, under the assumption that these fine-structure levels are populated by collisions, and not through direct excitation by infra-red photons (which is not an important excitation mechanism in the case of \\ion{Si}{ii}*), or fluorescence \\citep[see][]{2002MNRAS.329..135S}.  This volume density is higher than inferred for QSO-DLA environments \\citep{2002MNRAS.329..135S}, but typical of Galactic molecular clouds \\citep[e.g.][]{1999osps.conf....3B,2002ApJ...565..174R}. As these lines has never been clearly detected up to now in QSO-DLAs, the detection of these {Si}{\\sc ii}* lines in the \\grb\\ spectrum suggests an origin in the vicinity of the GRB place of birth (e.g. the star-forming region in which it exploded). Combining the measured \\ion{H}{i} column density with the order of magnitude estimate of the \\ion{H}{i} volume density, we obtain a size (diameter) of $\\sim$ 5pc (taking n$_{\\ion{H}{i}}=10^3$ cm$^{-3}$) and a mass of $\\sim$ 2$\\times$10$^3$ M\\subsun\\ for the \\ion{Si}{ii}* absorbing region.  With the volume density so high, one would expect hydrogen molecules to be present, which, surprisingly, we do not detect.  We obtain a rather strong upper limit on the mean molecular fraction of the gas in the GRB environment and the host galaxy: f$\\equiv$2$N$(H$_2$)/(2$N$(H$_2$)+$N$(\\ion{H}{i}))$\\lsim$10$^{-6}$. This could be explained by the low metallicity of the gas \\citep[see][]{ledoux}, but it may also be that the molecules in the GRB environment have been dissociated by the strong GRB UV/X-ray emission \\citep[e.g.][]{2002ApJ...569..780D}. In the latter case, however, the UV/X-ray flash would also ionize the neutral gas in the GRB vicinity \\citep[see][]{2002ApJ...569..780D}, which would make the high \\ion{H}{i} column density detection improbable. Therefore, a large fraction of the \\ion{H}{i} column density may not be located close to the GRB explosion site, but elsewhere in the host galaxy, while the high volume density \\ion{Si}{ii}* region (and the expected molecular hydrogen), is located in the vicinity of the burst. In this case, the disks of GRB host galaxies need to be much denser than the Galactic disk, as 7 random sight lines through the disk toward the location of the Earth would not result in 5 \\ion{H}{i} column densities above 10$^{21}$ cm$^{-2}$ \\citep[see Fig.  5 of][]{1990ARA&A..28..215D}, as is observed for GRB sightlines (see Fig. \\ref{histo}). Finally, the population of the \\ion{Si}{ii}* levels may have been partly caused by fluorescence of photons from the GRB itself, in which case the volume density estimate above is a strict upper limit.  The \\ion{H}{i} column density that we inferred toward \\grb\\ is higher than that of any (QSO- or GRB-) DLA measured using \\lya\\ in absorption. It is generally assumed that the apparent \\ion{H}{i} column density limit of N(\\ion{H}{i})$\\sim$10$^{22}$ atoms cm$^{-2}$ for QSO-DLAs is due to an observational bias against the detection of such high-column density systems, as these would obscure the background QSO if they contain some dust \\citep[e.g.][]{1984ApJ...278....1O,1993ApJ...402..479F}.  However, a radio-selected QSO survey for DLA systems by \\citet{2001A&A...379..393E} did not uncover a previously unrecognized population of N(\\ion{H}{i})$>10^{21}$ cm$^{-2}$ DLA systems in front of faint QSOs. An alternative scenario was proposed by \\citet{2001ApJ...562L..95S}: the lack of high \\ion{H}{i} column density systems could be due to the conversion of \\ion{H}{i} to H$_2$ as the neutral gas density increases. This picture is consistent with observations of Galactic molecular clouds \\citep[e.g.][]{1999osps.conf....3B}. In \\grb, however, we do not find any evidence for the presence of H$_2$ in addition to \\ion{H}{i} to support this scenario.  Future GRBs with possibly even larger \\ion{H}{i} column densities than that toward \\grb\\ could provide further constraints to the existence of a rapid conversion of \\ion{H}{i} to H$_2$ at high \\ion{H}{i} column densities.  We compared the metallicities and \\ion{H}{i} column densities of the (still very small) sample of GRB-DLAs with QSO-DLAs, and we found both quantities to be higher in GRB-DLAs than in QSO-DLAs.  This is not surprising, as GRBs are now known to probe massive-star forming regions \\citep{stanek,hjorth030329} where the gas density and the metallicity are higher than along random QSO sight lines through foreground galaxies.  A KS test applied to the column densities shows that the probability that the GRB- and QSO-DLA samples are drawn from the same parent distribution is very low (0.0006). On the other hand, two GRB afterglows have very low column densities.  A large sample of high-resolution spectra of GRB afterglows could provide statistical information about the distribution of the gas in high-redshift star-forming regions, in addition to the evolution of the metallicity and dust and H$_2$ contents of GRB host galaxies.  Such a sample can be created in the years to come thanks to rapid and accurate GRB localizations from future satellite missions such as Swift\\footnote{see http://swift.gsfc.nasa.gov/} and EXIST\\footnote{see http://exist.gsfc.nasa.gov/}."
    },
    "0403/astro-ph0403205_arXiv.txt": {
        "abstract": "We present the $UBVRI$ CCD photometry in the region of the open cluster NGC 2421. Radius of the cluster is determined as $\\sim$ 3$^\\prime$.0 using stellar density profile. Our Study indicates that metallicity of the cluster is $Z \\sim$ 0.004. The reddening $E(B-V) = 0.42\\pm$0.05 mag is determined using two colour $(U-B)$ versus $(B-V)$ diagram. By combining the 2MASS $JHK$ data with the optical data we determined $E(J-K) = 0.20\\pm$0.20 mag and $E(V-K) = 1.15\\pm$0.20 mag for this cluster. Colour-excess diagrams show normal interstellar extinction law in the direction of the cluster. We determined the distance of the cluster as 2.2$\\pm$0.2 Kpc by comparing the ZAMS with the intrinsic CM diagram of the cluster. The age of the cluster has been estimated as 80$\\pm$20 Myr using the stellar isochrones of metallicity $Z = 0.004$. The mass function slope $x = 1.2\\pm0.3$ has been derived by applying the corrections of field stars contamination and data incompleteness. Our analysis indicate that the cluster NGC 2421 is dynamically relaxed. ",
        "introduction": "The investigation of young open star clusters provides us a powerful tool to understand the structure and history of star formation in our Galaxy. In order to fully exploit the information provided by open clusters we must know their accurate ages, distances, reddenings, metal abundances and stellar contents. For this, multicolour CCD photometric observations have proved to be very useful. With the development of more accurate stellar models it has been possible to provide a better estimate of the cluster ages simply by comparing theoretical isochrones with the observed CCD colour-magnitude (CM) diagrams. So, In this paper we have considered an open cluster NGC 2421 with the aim of presenting new accurate CCD photometry. From this photometry we select photometric members and derive several fundamental parameters, such as distance, interstellar reddening, metallicity and age as well as luminosity and mass function.  The young open cluster NGC 2421 = C0734$-$205 ($\\alpha_{2000} = 07^{h}36^{m}16^{s}$, $\\delta_{2000}=-20^{d}36^{\\prime}44^{\\prime\\prime}$; $l = 236^{\\circ}.24$, $b = 0^{\\circ}.08)$ is classified as a Trumpler class I2m by Ruprecht (1966). This cluster was first studied by Moffat \\& Vogt (1975) photoelectrically and derived a distance of about 1.87 Kpc having $E(B-V)=$ 0.47$\\pm$0.05 mag and age less than 10$^{7}$ years. Ramsay \\& Pollacco (1992) also studied this cluster using CCD photometry and found a colour excess $E(B-V)=$ 0.49$\\pm$0.03 mag but a distance of 2.75 Kpc. To our knowledge no other studies have been carried out for the cluster NGC 2421 so far.  The layout of the paper is as follows. In Sec. 2 we briefly describe the observations and data reduction strategies as well as comparison with the previous photometry. Sec. 3 is devoted on the detailed analysis of the present photometric data for the determination of cluster parameter. Finally, Sec. 4 summarizes the main results of the paper.  \\begin{figure*} \\centering \\hspace{0.7cm}\\psfig{figure=ME150fig1.ps,width=10cm,height=10cm} \\caption{Finding chart of the stars in the cluster NGC 2421. The (X, Y) coordinates are in pixel units corresponding to 0$^{\\prime\\prime}$.72 on the sky. Direction is indicated in the map. Filled circles of different sizes represent brightness of the stars. Smallest size denotes stars of $V$$\\sim$20 mag. Open circle having centre at 'C' in the chart represent the cluster size.} \\end{figure*} ",
        "conclusions": "We have investigated the area of open cluster NGC 2421 using $UBVRI$ CCD and 2MASS $JHK_s$ data. The main results of our analysis are the following.  \\begin{enumerate} \\item The radius of the cluster is determined as 3$^\\prime$.0 which corresponds to 1.9 pc at the distance of the cluster. \\item We estimated the abundance of the cluster stars as $Z =$ 0.004 using excess in $(U-B)$. The $(U-B)$ versus $(B-V)$ colour-colour diagram yields $E(B-V) = 0.42\\pm$0.05. The analysis of $JHK$ data in combination with the optical data provide $E(J-K) = 0.20\\pm$0.20 mag and $E(V-K) = 1.15\\pm$0.20 mag. Our analysis shows that interstellar extinction law is normal towards the cluster direction. No stars found which are having near-IR fluxes due to the presence of circumstellar material around them. \\item A ZAMS fitting procedure gives a distance of 2.2$\\pm$0.2 Kpc for this cluster which is also supported by the value of 2.3$\\pm$0.3 Kpc determined by us using the optical and near-IR data. An age of 80$\\pm$20 Myr is determined using the isochrones of $Z =$ 0.004 given by Bertelli et al. (1994). \\item The mass function slope $x = 1.2\\pm0.3$ is derived by considering the corrections of field star contamination and data incompleteness. Our analysis indicate that the cluster NGC 2421 is dynamically relaxed and one plausible reason of this relaxation may be the dynamical evolution of the cluster. \\end{enumerate}"
    },
    "0403/astro-ph0403033_arXiv.txt": {
        "abstract": "{ We report the results of a non-LTE Fe, O and Mg abundance analysis of the carbon-nitrogen-rich ultra-metal-poor giants CS\\,29498--043 and CS\\,22949--037. The abundance of oxygen has been derived from measurements of the oxygen triplet at 7771--5 \\AA\\ in high resolution spectra obtained with KeckI/HIRES and the forbidden line [O\\,{\\sc i}] 6300 \\AA\\ detected in the TNG/SARG spectra of CS\\,29498-043. Detailed non-LTE analysis of Fe lines has provided reliable stellar parameters which, however, do not resolve the oxygen abundance conflict as derived from the O\\,{\\sc i} 7771-5 \\AA\\ triplet and the [O\\,{\\sc i}] 6300 \\AA\\ forbidden lines. We obtained the following oxygen abundance: for CS\\,22949--037 [O/Fe] =3.13, 1.95; and for CS\\,29498--043; [O/Fe]=3.02, 2.49, based on the O\\,{\\sc i} 7771--5 \\AA\\ triplet and the [O\\,{\\sc i}] 6300 \\AA\\ forbidden line, respectively. A similar conflict appears to exist between the forbidden resonance line Mg\\,{\\sc i} 4571 \\AA\\ and several subordinate lines, such as Mg\\,{\\sc i} 5172 and 5183 \\AA. Our analysis demonstrates the failure of standard plane--parallel atmosphere models to describe the physical conditions in the line-forming regions of these ultra-metal-poor giants. ",
        "introduction": "Surveys of metal-poor stars (Beers et al. 1992) and their abundance studies (Ryan, Norris \\& Beers 1996; Norris, Ryan \\& Beers\\,1999) are aimed at investigating the chemical evolution of the Galaxy and the nucleosynthetic yields of supernovae. Despite the numerous studies in this field, the current situation with abundance trends of various $\\alpha$-elements is very confusing. Observations of Stephens \\& Boesgaard (2002) demonstrate that the [$\\alpha$/Fe] ratios for  Ca, Si, Ti and Mg do not show a flat {\\itshape plateau} at [Fe/H] $< -1$, as was claimed in previous studies. The results obtained by Idiart \\& Thevenin\\,(1999) and Stephens \\& Boesgaard (2002) cannot be called  {\\it consistent} with the analysis presented by Carretta et al.\\,(2002) and many others (see McWilliam 1997). The situation with oxygen is far from being resolved (Israelian, Garc\\'\\i a L\\'opez \\& Rebolo \\,1998; Israelian et al.\\,2001; Nissen et al.\\ 2002; Takeda 2003; Fulbright \\& Johnson\\ 2003), while a new debate over the sulfur abundance in metal-poor stars has already emerged (Israelian \\& Rebolo 2001; Takada-Hidai et al.\\ 2002; Nissen et al.\\ 2003). These and many other studies clearly show  that  [$\\alpha$/Fe] $>$ 0 for the great majority of  metal-poor stars in the Galaxy. However, it is hard to speak of any trend when the abundance ratios in many stars computed by different authors disagree by more than 0.3--0.4 dex. There are many aspects to this serious problem, and we shall not discuss them further here. Abundance analysis of ultra-metal-poor stars with [Fe/H] $< -$3 gives rise to even more enigmas into this field.  It is well known that the chemical composition of the atmospheres of halo dwarfs is not altered by any internal mixing and therefore provides a good opportunity to constrain Galactic chemical evolution models. Unfortunately, most of the known ultra-metal-poor stars are not dwarfs but giants (Beers et al. 1992), which pose two serious problems. First, the atmospheric parameters of giants are more uncertain and second, their surfaces can be polluted by enriched material that has been either dredged from the stellar interior or transferred from a companion star. Oxygen is a key element in this scheme as it can help to distinguish between pristine and pollution origins of other elements and also show which of the aforementioned processes was dominant. There have been intensive investigations of the oxygen abundances in halo stars over the last five years. Abundances derived from near-UV OH lines in metal-poor dwarf stars (Israelian et al.\\,1998; Boesgaard et al. 1999; Israelian et al.\\,2001) show that the [O/Fe] ([O/Fe] = log(O/Fe)$_\\star$--log (O/Fe)$_\\odot$) ratio increases from 0 to 1 between [Fe/H] = 0 and $-3$. The abundances derived from low-excitation OH lines agreed well with those derived from high-excitation lines of the O\\,{\\sc i} triplet at 7771--5 \\AA\\ (Israelian et al. 1998, 2001; Boesgaard et al. 1999 ; Nissen et al.\\,2002). It seems that even the [O\\,{\\sc i}] forbidden line at 6300 \\AA\\ supports the ``quasi-linear'' trend of [O/Fe] (Nissen et al.\\,2002) when standard 1D atmospheric models are employed. While some authors claim a good agreement between the forbidden line [O\\,{\\sc i}] 6300 \\AA\\ and the near-IR triplet (Mishenina et al. 2000; Nissen et al.\\,2002), others suggest the opposite (Carretta, Gratton, \\& Sneden 2000). In a recent study Takeda (2003) found that the disagreement between the triplet and the forbidden line tends to be larger for cool giants. Fulbright \\& Johnson (2003) support this conclusion in a detailed study of 55 subgiants and giants. These authors conclude that it is impossible to resolve the disagreement in the two indicators without adopting an ad hoc temperature scale that is incompatible with standard temperature scales such as IRFM and H$\\alpha$. In fact, the [O/Fe] trend obtained based on the ad hoc scale (see Fig.\\ 13 of Fulbright \\& Johnson 2003) is identical to the trends presented by Israelian et al. (2001) and Nissen et al. (2002). There is no {plateau} at [O/Fe] = 0.5. Despite considerable observational effort the trend of the [O/Fe] ratio in the halo is still unclear. However, the latest studies (Israelian et al. 2001; Nissen et al. 2002; Takeda 2003; Fulbright \\& Johnson 2003) suggest that dwarfs provide more reliable and consistent abundances than giants. It is not clear how the 3D effects will resolve this conflict since the latter do not predict an agreement between different oxygen abundance indicators in dwarfs (Asplund \\& Garc\\'\\i a P\\'erez 2002).  McWilliam et al.\\,(1995) were the first to carry out a detailed spectroscopic analysis of CS\\,22949-037 and to confirm that the star is very metal-poor with an $\\alpha$-element excess. Furthermore,  Depagne et al.\\ (2002) performed a more detailed investigation of this object and found a large excess of oxygen ([O/Fe] = 2.0) and sodium ([Na/Fe] = 2.1). Zero-heavy-element supernovae models with fall back have been invoked in order to interpret the elemental abundance ratios in this star. Aoki et al.\\, (2002) have presented a detailed analysis of another ultra-metal-poor giant CS\\,29498--043 with a very high abundance excess of [Mg/Fe] = 1.81. Both,  CS\\,22949--037 and CS\\,29498--043 exhibit a large overabundance of N and C but show no significant enhancement of neutron-capture elements. It is possible that the surfaces of these stars have been polluted by enriched material, either dredged from the star's inner core or transferred from a companion star. The abundances of $\\alpha$-elements could be used to discriminate in favour of one of these hypotheses or to confirm a pristine origin. The detailed comparison of elemental abundances may provide important constraints on the properties of the first supernova progenitors.   In this article we present observations of the oxygen triplet in CS\\,22949--037 and CS\\,29498-043, as well as the detection of the  forbidden line [O\\,{\\sc i}] 6300 \\AA\\ in CS\\,29498--043. The stellar parameters and the abundances of iron, oxygen and magnesium were derived in non-LTE. We report a significant discrepancy between the abundances derived from the oxygen triplet and the forbidden line. The conflict cannot be resolved under any circumstances, at least for CS\\,22949-037. A similar conflict was found for Mg. We question the validity of standard plane--parallel models of atmospheres employed in the present analysis. ",
        "conclusions": "Observations with Keck I/HIRES have revealed strong lines of the oxygen near-IR triplet in the spectra of the ultra-metal-poor giants CS\\,22949--037 and CS\\,29498--043. The forbidden line of oxygen with EW = 60$\\pm$10m\\AA\\ was observed with TNG/SARG. A detailed non-LTE analysis of Fe has been carried out and a new set of the atmospheric parameters have been obtained. Our analysis suggests that the gravities of metal-poor giants derived from the LTE Fe analysis are strongly underestimated because of the neglect of non-LTE effects.  The oxygen abundance in CS\\,22949--037 and CS29498--043 derived from the triplet and the forbidden line differ by 1.18 and 0.53 dex, respectively. This disagreement cannot be explained by a non-LTE effects, quality of the data and/or uncertainties in stellar parameters. Other mechanisms must be invoked in order to explain this puzzle. A similar discrepancy was found for Mg when comparing the abundances obtained from the resonance line 4571 \\AA\\ and several strong subordinate lines. Based on the present analysis we propose that the Kurucz (1992) models are not reliable for these ultra-metal-poor giants."
    },
    "0403/astro-ph0403519_arXiv.txt": {
        "abstract": "We present VLT-UVES echelle spectra of the \\Lya~emission and absorption in five radio galaxies at redshifts z=2.55--4.1. Together with data from our pilot study, we have a sample of 7 such systems with radio source sizes $\\sim 1-90$\\kpc~with which to address the origin of the absorbing gas. Echelle resolution again reveals that some systems with $N_{\\rm{HI}}>10^{18}$\\psqcm~in lower resolution data in fact consist of several weaker absorbers with $N_{\\rm{HI}}<10^{15}$\\psqcm. We identify two groups of HI absorbers: strong absorbers with $N_{\\rm{HI}} \\simeq 10^{18}-10^{20}$\\psqcm~and weaker systems with $N_{\\rm{HI}} \\simeq 10^{13}-10^{15}$\\psqcm. There are none at intermediate $N_{\\rm{HI}}$.  The strong absorbers may be a by-product of massive galaxy formation or could instead represent material cooling behind the expanding bow-shock of the radio jet, as simulated by Krause. New observations are required to discriminate between these possibilities. We argue that the weaker absorbers with $N_{\\rm{HI}} \\simeq 10^{13}-10^{15}$\\psqcm~are part of the \\Lya~forest, as their rate of incidence is within a factor of 2--4 of that in the IGM at large. Such column densities are consistent with models of a multi-phase proto-intracluster medium at $z>2$. ",
        "introduction": "High-redshift ($z>2$) radio galaxies (HzRGs) are among our best probes of massive galaxy formation. The K-z relation demonstrates that their supermassive black holes ($10^{9}-10^{10}$\\Msun) reside in the most massive elliptical galaxies at their epoch (e.g. Jarvis et al.~2001a; Willott et al.~2003). Measurements of galaxy clustering and of galaxy over-densities around individual sources, imply that they inhabit proto-cluster environments which will evolve into rich clusters at $z=0$ (for a review of radio source clustering and proto-clusters, see R\\\"{o}ttgering et al.~2003).  Deep within these gravitational potential wells and surrounding the HzRGs to radii of tens to several hundred \\kpc, are luminous ($>10^{44}$\\ergps~in \\Lya) emission line nebulae. Their properties bear witness to the complex gaseous processes in nascent massive galaxies, viz. gravitational collapse and shock heating, cooling, bursts of star formation and feedback in the form of superwinds; the latter can inject substantial quantities of energy and metals into the gas, enriching it to solar or super-solar metallicity out to tens of \\kpc~(e.g. Villar-Martin et al.~2001). The presence of a powerful radio source has a profound impact on such an environment: within the dense interstellar medium of the galaxy, the relativistic radio jet drives shocks into the clumpy gas, leading to extreme kinematics ($FWHM > 1000$\\kmps; Villar-Martin, Binette and Fosbury~1999), additional star formation and the emission line signatures of shock ionisation (Best et al.~2000; De Breuck et al.~2000). At radii beyond the radio source, the gas is kinematically quiescent (reflecting purely gravitational motions) and the continuum from the quasar nucleus is the main source of ionisation (Villar-Martin et al.~2003).  New insight into the gaseous haloes of HzRGs came with the discovery by R\\\"{o}ttgering et al.~(1995) and van Ojik et al.~(1997) that the majority of small radio sources (those with projected linear size below 50\\kpc) exhibit spatially resolved absorption in their \\Lya~profiles, with HI column densities of $10^{18}-10^{19.5}$\\psqcm. For the $z=2.92$ HzRG 0943--242, Binette et al.~(2000) discovered associated CIV$\\lambda\\lambda$1548,1551~absorption at the same redshift as the main \\Lya~system. Through photoionisation modelling they argued that the absorbing gas is of lower metallicity ($Z \\sim 0.01Z_{\\rm{\\odot}}$) and is situated beyond the emission line gas, outside the high pressure radio source cocoon. This material, they claimed, is a relic reservoir of low metallicity, low density ($n \\sim 10^{-2.5}$\\pcm) gas. In larger radio sources, the radio source cocoon plausibly engulfs and pressurises the outer halo entirely, causing it to be seen in emission rather than absorption, and thereby accounting for the anti-correlation between radio source size and the presence of HI absorption found by van Ojik et al.~(see also Jarvis et al.~2001b for the HI absorption:radio source size correlation).  With the Ultraviolet-Visible Echelle Spectrograph (UVES) at the VLT, we recently examined the absorbers in two HzRGs with over ten times the spectral resolution of R\\\"{o}ttgering et al. and van Ojik et al. Our targets were the aforementioned 0943--242, and 0200+015 ($z=2.23$) and the results were presented by Jarvis et al.~(2003) and  Wilman et al.~(2003). In 0943--242, the absorbers exhibit no additional structure -- the main absorber is still consistent with an HI column density of $10^{19}$\\psqcm, in line with the Binette et al.~(2000) model. In contrast, a very different view of 0200+015 emerges: the single absorber with HI column density $\\simeq 10^{19}$\\psqcm seen at low resolution now splits into two $\\sim 4 \\times 10^{14}$\\psqcm~systems; these extend by more than 15\\kpc~to obscure addtional \\Lya~emission coincident with a radio lobe and highly fragmented absorption is also seen on the red wing of the emission line at this position. The detection of one of the sub-systems in CIV suggests that the absorbing gas has undergone substantial metal enrichment, to perhaps as high as $10Z_{\\rm{\\odot}}$. Based on the smaller radio source size in 0943--242 (26\\kpc~versus 43\\kpc~for 0200+015), we conjectured that the radio source age (as inferred from its linear size) is the parameter controlling the evolution of: (i) the structure/kinematics of the absorbing halo, through interaction and shredding of the initially quiescent shells; (ii) its metallicity, through enrichment by starburst superwind triggered concurrently with the nuclear radio source.  If the specific effects of the radio source can be understood, high resolution spectroscopy of such \\Lya~haloes provides an opportunity to probe the \\Lya~forest around high-redshift galaxies and hence to test theories for how such systems form out of the intergalactic medium (IGM). For quasars, such attempts are completely hampered by their strong UV continua, which through photoionisation locally reduce the number of \\Lya~forest clouds below the IGM average (the so-called proximity effect). Lacking strong line of sight continua, radio galaxies offer a more direct probe of the distribution of HI in such dense environments. Due to the nature of hierachical clustering, the opacity of the \\Lya~forest is expected to increase in the vicinity of massive objects, and observations of background quasars do indeed show excess HI absorption at radii of 1--5$h^{-1}$\\Mpc~from $z \\sim 3$ Lyman-break galaxies (LBGs) (Adelberger et al.~2003). However, within $0.5h^{-1}$\\Mpc~there appears to be a deficit of HI, which they attribute to the influence of `superwind' outflows from the LBGs. From a theoretical stand-point it is not clear that superwinds alone can account for this observation, and additional effects such as local photoionization (Bruscoli et al.~2003) and the `filling in' of absorption with \\Lya~emission from the LBG itself (Kollmeier et al.~2003) may be needed.  In order to address the origin and evolution of the absorbing shells and their connection with the \\Lya~forest, we need to expand our sample of HzRGs with high resolution spectroscopy. For this reason, we have recently acquired UVES spectra for a further 5 HzRGs, yielding a sample of 7 HzRGs with radio source sizes in the range $\\sim 1-90 $\\kpc~(projected). Here we present the analysis and interpretation of the new dataset. The paper is structured as follows: in section 2, we describe the target selection, observations and data reduction, and in section 3 present the \\Lya~profiles and absorption model fits on an object-by-object basis. In section 4 we compile numerous statistics for the absorption line ensemble and compare them with the properties of HI absorption in the IGM. In section 5, we interpret these results in the light of simulations of: (i) expanding radio sources in dense environments; (ii) the \\Lya~forest around high-redshift galaxies. Section 6 contains our conclusions.  Throughout the paper we assume a cosmology with $H_{\\rm{0}}=70$\\kmpspMpc, $\\Omega_{\\rm{M}}=0.3$ and $\\Omega_{\\rm{\\Lambda}}=0.7$. ",
        "conclusions": "We have built on our successful pilot study and increased the number of high-redshift radio galaxies with UVES echelle spectroscopy of \\Lya~from 2 to 7. With this enlarged ensemble of absorption systems we can begin to make some inferences about their origin. The new spectra demonstrate once again that HI absorption is more complex when seen at this much higher resolution: several absorbers identified in low-resolution spectra with $N_{\\rm{HI}}>10^{18}$\\psqcm~now fragment into a number of much weaker systems with $N_{\\rm{HI}}<10^{15}$\\psqcm. There appears to be a gap in the $N_{\\rm{HI}}$ distribution, as we identify strong absorbers with $N_{\\rm{HI}} \\simeq 10^{18}-10^{20}$\\psqcm~and weaker systems with $N_{\\rm{HI}} \\simeq 10^{13}-10^{15}$\\psqcm, but none at intermediate $N_{\\rm{HI}}$.  We discussed the origin of the strong absorption systems with $N_{\\rm{HI}}>10^{18}$\\psqcm~in the framework put forward by Krause~(2002), based on radio source physics. A highly supersonic jet expanding into a dense proto-galactic medium will be surrounded by a quasi-spherical bow shock. The simulations by Krause~(2002) show that the cooling of post-shock material will lead to neutral column densities comparable to those observed, and that due to instabilities the shells will fragment and offer much less absorption towards larger radio sources. The model provides a very natural explanation for the observation that the shells have high covering factors and surround the entire galaxy. Appealing though it is, it would be premature to enshrine it as the canonical model for the strong absorbers in HzRGs. Several new challenging observations are needed before we can determine whether such strong absorbers really are produced by the radio source bow shock or whether they are instead a by-product of massive galaxy formation. Firstly, in the HzRGs themselves deep, high resolution spectra should be obtained of the faint kinematically-quiescent \\Lya~haloes, identified by Villar-Martin et al.~(2003). Secondly, such spectroscopy should also be obtained of a carefully selected sample of comparably massive high-redshift galaxies without radio-loud AGN, comprising SCUBA galaxies and narrow-line X-ray type II quasars.  We cannot exclude the possibility that some of the lower column density systems with $N_{\\rm{HI}}=10^{13}-10^{15}$\\psqcm~form by fragmentation of larger shells until the simulations of the bow-shock absorbers are extended into this regime. The rate of incidence of such absorption in our spectra and their distribution on the $b$--$N_{\\rm{HI}}$ plane suggests, however, that these systems could be part of the IGM \\Lya~forest population. Simulations predict that the opacity of the forest should increase within a few Mpc~of a massive galaxy at these epochs. Our results show an enhancement (by factors 2--4) in the number of \\Lya~forest systems (relative to the IGM) at $z=2-3$ and a deficit by similar factors in the two $z>3.5$ systems, although on the grounds of small number statistics we do not claim that this implies evolution over this redshift interval. For a proper comparison, independent of any Voigt fitting uncertainties and systematics, new simulations of the \\Lya~forest should be performed. Unlike those in the literature which predict the HI opacity at the galaxy redshift as a function of transverse distance from it, the new simulations should yield the \\Lya~forest along the line of sight to a massive object in redshift space, as a function of galaxy mass and epoch. In lieu of this, we note that the observed column densities are consistent with simple estimates for a multiphase proto-intracluster medium at $z>2$."
    },
    "0403/astro-ph0403343_arXiv.txt": {
        "abstract": "We present \\chandra\\ X-ray observations of the young supernova remnant (SNR) \\e05 in the Large Magellanic Cloud (LMC), believed to be the product of a Type Ia supernova (SN Ia).  The remnant is very round in shape, with a distinct clumpy shell-like structure that extends to an average radius of $14.8\\arcsec$ (3.6 pc) in the X-ray band.  Our \\chandra\\ data reveal the remnant to be rich in silicon, sulfur, and iron.  The yields of our fits to the global spectrum confirm that \\e05 is the remnant of an SN Ia and show a clear preference for delayed detonation explosion models for SNe Ia.  The \\chandra\\ spectra extracted from radial rings are in general quite similar; the most significant variation with radius is a drop in the equivalent widths of the strong emission lines right at the edge of the remnant. We study the spectrum of the single brightest isolated knot in the remnant and find that it is enhanced in iron by a factor of roughly two relative to the global remnant abundances. This feature, along with similar knots seen in Tycho's SNR, argues for the presence of modest small-scale composition inhomogeneities in SNe Ia.  The presence of both Si and Fe, with abundance ratios that vary from knot to knot, indicates that these came from the transition region between the Si- and Fe-rich zones in the exploded star, possibly as a result of energy input to the ejecta at late times due to the radioactive decay of $^{56}$Ni and $^{56}$Co. Two cases for the continuum emission from the global spectrum were modeled: one where the continuum is dominated by hydrogen thermal bremsstrahlung radiation; another where the continuum arises from non-thermal synchrotron radiation. The former case requires a relatively large value for the ambient density ($\\sim$1 cm$^{-3}$).  Another estimate of the ambient density comes from using the shell structure of the remnant in the context of dynamical models.  This requires a much lower value for the density ($<$0.05 cm$^{-3}$) which is more consistent with other evidence known about \\e05.  We therefore conclude that the bulk of the continuum emission from \\e05 has a non-thermal origin. ",
        "introduction": "Type Ia supernovae (SN Ia) are characterized by the absence of Balmer lines and the presence of strong silicon lines in their optical spectra \\citep{hillnie}.  They yield mostly intermediate mass elements from Si to Ca and are the main sources of iron production in galaxies. Each SN Ia produces $\\sim$0.8 ${\\rm M{_\\sun}}$ of Fe \\citep{iwam}, making them important for understanding the chemical evolution of galaxies.  The peak magnitude and shape of their light curves are fairly homogeneous, which makes them useful as distance indicators and tools for cosmology \\citep{rpk,riess,perl}.  For these reasons, SN Ia's are the objects of much study.  Still, there is much that remains uncertain about them.  The progenitors of SN Ia's are believed to be degenerate C+O white dwarfs. Although models for sub-Chandrasekhar mass explosions exist, the favored explanation is that the degenerate object somehow exceeds the Chandrasekhar limit before exploding.  Two possibilities include a single degenerate scenario, in which the white dwarf accretes mass from a companion star, and a double degenerate scenario, in which a pair of white dwarfs coalesce \\citep{branch}.  It is still unclear whether one, both, or none of these offers the correct explanation.  In addition, the mechanism of the explosion, i.e., where the ignition begins and how it propagates outward, continues to elude us \\citep{iwam}.  Two cases studied in detail include a fast deflagration and a delayed detonation.  In the former, a deflagration begins close to the center of the progenitor star and propagates subsonically.  The delayed detonation case involves a deflagration initially, which is transformed at some density to a detonation as the star expands. Models for these two cases are parameterized by the speed of the burning front and the location of the ignition in density space; we compare our results with those of \\citet{iwam}.  By studying the remnants of SN Ia explosions, we can hope to gather clues to some of these puzzles.  The supernova remnant (SNR) \\e05 in the Large Magellanic Cloud (LMC) is believed to be the remnant of an SN Ia \\citep{tuohy, hughes95}.  It was discovered in an X-ray survey of the LMC with the {\\it Einstein X-ray Observatory} \\citep{long}. However, much of the previous study on this remnant was done via optical spectroscopy.  In H$\\alpha$, the remnant appears as a very circular shell \\citep{tuohy}.  \\e05 shows no detectable emission in [\\ion{O}{3}] and [\\ion{S}{2}], indicating it is a Balmer-dominated remnant \\citep{tuohy, smith91,smith94}.  Remnants such as this generally have two components to the H$\\alpha$ emission line: a narrow component produced by neutral hydrogen being excited by the shock, and a broad component produced by charge exchange with fast protons post-shock.  In \\e05, Smith et al. (1994) found the narrow components to be $\\sim$25--31 km/s, higher than expected for neutral hydrogen and a possible indication of the presence of a cosmic ray precursor. This SNR is also intriguing in that its broad component is believed to be so broad that it has escaped detection, pointing to very fast shocks ($\\geq$2000 km/s) and a young age ($\\leq$1000 yr) (Smith et al. 1991; Smith et al. 1994).  \\citet{tuohy} found that \\e05 is expanding into a low density medium, with $n_{\\rm H}\\lesssim0.02$ cm$^{-3}$.  A previous X-ray study of this SNR with {\\it ASCA} revealed it to be strong in Si and Fe L emission, with little O, Ne, and Mg emission \\citep{hughes95}.  This remnant has also been observed in the radio, showing a spectral index of $\\alpha=0.46$ \\citep{mathewson} and a flux density of 0.066 Jy at 1 GHz \\citep{henrey}. ",
        "conclusions": "Our analyses of \\e05 point us toward several conclusions which give us insight into various aspects of SNRs.  We find that the abundances obtained from our model fits, when compared with theoretical models of supernovae, confirm that \\e05 is the remnant of an SN Ia.  In addition, the abundance fits are able to give us information as to which theoretical model is more appropriate.  The global spectrum of \\e05 as well as that of a small knot of material give us important clues as to the properties of iron in the remnant.  Our deprojection and image fits provide us with a shell structure to this SNR, which we use to derive information about the dynamics.  This in turn allows us to draw the conclusion that the bulk of the continuum emission in this remnant must be non-thermal in nature.  \\subsection{SN Ia Yields} In Figure 8 we plot our abundances, normalized to Si = 1, for Cases H (squares) and S (x's).  These yields are also compared to those determined by two models of SN Ia, the W7 model (fast deflagration, stars) and the WDD3 model (delayed detonation, circles) of \\citet{iwam}.  The fast deflagration involves a deflagration wave that propagates outward subsonically, and the delayed detonation scenario begins as a deflagration that then transitions to a detonation at some density.  The other models presented by \\citet{iwam} are effectively represented here.  Their W70 model is similar to W7, and the other DD models are similar to WDD3; therefore we did not plot these.  It is quite clear that the low-Z elements of O, Ne and Mg are 1--2 orders of magnitude less abundant than the higher-Z elements of Si - Ca.  This is a signature of SN Ia's \\citep{thiele} and confirms such an origin for this remnant.  Our yields for O, Ne, and Mg agree better with the WDD3 (delayed detonation) model than the W7 (fast deflagration) model.  Even if all of the Si has not been shocked, the low O/Si ratio would only be exacerbated and the W7 model would still fail to fit the data. \\citet{badenes} have compared various explosion models with the spectrum of Tycho's SNR.  They are able to show that the pure detonation and pure deflagration models do not satisfactorily reproduce important features in Tycho's spectrum.  Their results point to a delayed detonation explosion as the best model.  In their Figure 4, they plot the integrated EM of each species in the remnant as a function of time.  In this context, we looked at another Type Ia remnant, DEM L71.  This SNR is 4400 years old \\citep{ghav} and rich in Si, S, and Fe \\citep{hughl71}.  Shocks in older remnants probe different parts of the SN ejecta.  In comparison to their Figure 4, we find that at the age of DEM L71 the models which show the same dominant elements as seen in this SNR are also delayed detonation models.  This suggests that the delayed detonation scenarios for SN Ia explosions may be more appropriate.  In addition, all three remnants indicate some preference for higher transition densities among the delayed detonation models.  All of these results reveal the power of remnant studies of Ia's to elucidate explosion mechanism physics.  The iron yields we observe are low compared to the models of \\citet{iwam}.  Our model does underestimate the amount of iron present in the remnant.  However, it is extremely unlikely that uncertainties in Fe L-shell atomic physics could account for the difference of a factor of $\\sim$30 between the models and our yield.  We interpret this discrepancy to be due to the fact that the reverse shock has not propagated far enough into the remnant to shock all of the iron.  This is not unprecedented; there is evidence for unshocked iron ejecta in another young Type Ia, SN1006, where \\ion{Fe}{2} absorption lines have been detected in the center of the remnant \\citep{hamilton, wu}. This is also a feature of theoretical models \\citep{badenes}.  One puzzle we encountered was that our fits required iron to be in a different thermodynamic state than the other elements, with a much higher temperature and slightly higher ionization timescale.  We attempted a fit with a component of iron at the same state as the other elements and found we could allow a maximum abundance of 0.85 for both Cases H and S, which is 0.05 and 0.07, respectively, relative to Si.  However, such a fit did not reproduce the Fe K line and quite clearly underestimated the flux from 0.9--1.4 keV.  The fact that iron appears to be in a different thermodynamic state implies that it may have come from a different portion of the progenitor than the other elements.  \\citet{hwang} found a similar situation in the case of Tycho's SNR, where the Fe K emission required a higher temperature, but a lower ionization timescale than the other elements. Intuitively, this might be what one would expect if there were a temperature gradient and the ejecta were stratified, with Fe interior to the other elements.  We may be seeing an indication of the ejecta density profile here.  For example, the constant density ejecta profile of \\citet{dc} predicts an increasing temperature gradient behind the contact discontinuity.  However, it must be emphasized that X-ray studies probe the electron temperature, while the main thermal energy rests in the ions.  The extent of the equilibration between these two components is still under investigation (e.g., Rakowski, Ghavamian, \\& Hughes 2003).  In addition, it is not simple to relate the values fit for in our models to the {\\em actual} temperature and timescales of the ejecta because of astrophysical complications.  For example, our fits assume a single, constant temperature for each species everywhere in the remnant.  \\subsection{Iron Enhanced Knot} We have determined that the knot in the northern interior of \\e05 is enhanced in iron (see Figure 9).  As seen in projection, this knot is contained mostly in the second northeast half-ring.  We used the spectrum of that particular ring as a template for the knot's spectrum.  If we simply renormalize the spectrum, the fit is poor ($\\chi^2/d.o.f. = 59.0/32$).  However, if we allow the iron abundance to be free, we find a good fit at about twice the iron abundance of the half-ring ($\\chi^2/d.o.f. = 32.6/32$).  A fit of pure iron also gave a poor match to the data, indicating that although this knot is enhanced in iron, it is not pure iron ejecta.  We looked at this particular knot simply because it was the brightest and most isolated.  However, given that our image fits (\\S 3.2) show the presence of significant clumpiness in this SNR, it is possible that other knots show variations in the Si/Fe abundance ratio, although we are unable to definitively measure this.  Knots with different X-ray spectra have been detected in Tycho, the so-called ``typical'' Type Ia SNR, on the southeastern edge of the remnant \\citep{van}.  Recent work has confirmed that the spectral differences are a result of differences in Si/S and Fe abundances in these knots \\citep{dec}.  In neither Tycho nor \\e05 do we see knots of {\\em pure} iron or silicon.  Models of SN Ia explosions (i.e., Iwamoto et al. 1999) predict that the ejecta are highly stratified, with an outer O-rich layer, then a Si-rich region, followed by an inner Fe-rich layer.  However, the boundaries between these regions are not sharp; the abundances vary smoothly from one zone to the next.  The differing ratios of Si/Fe observed in the remnants' knots suggest to us that they originated in the transition region between the iron- and silicon-rich zones of the ejecta.  Clumping in SN Ia's has been investigated by \\citet{wc}.  They propose that the nickel bubble effect, in which the radioactive nickel expands and compresses the shell of material around it, may be responsible for producing knots. The shell surrounding the nickel would be the Si/Fe transition zone. Further investigation into possible origins of knots in this region would be worthwhile.  \\subsection{Shell Structure} Our deprojection analysis clearly indicates the presence of a true shell structure for this SNR (see Figure 7).  The deprojection was done for O, Si, S, Fe, and H (the latter only for Case H).  In both cases, and for both hemispheres, we found that the metals all follow this structure, with essentially no material in the interior of the remnant (see below for a discussion of the southwest hemisphere).  The metals extend over the two outermost shells in the deprojection with a peak density in the shell at $R \\sim 3.2$ pc and a slightly lower density in the outer shell at $R \\sim 3.6$ pc.  Hydrogen (Case H) shows a slightly different structure: there is almost no H in the $R \\sim 3.2$ pc shell or further in the interior; rather the continuum comes from the outermost shell.  The continuum emission in Case S, when deprojected, also comes only from the outermost shell.  In each case, therefore, this is clear evidence for a blast-wave component to the remnant.  We note however that the blast wave component appears to be contaminated somewhat by ejecta, presumably by fragments that preceed the main shell.  In the southwest hemisphere, the structure is more complicated.  There is a shell at a radius of $\\sim$ 3.2 pc that is the counterpart to the northeast shell.  In addition, there appears to be another spectral component peaking at a smaller radius, with an apparent gap between the two shells.  We do not believe this is an actual shell, but rather a projection effect due to the clumpy region on the southwest side. Given simply a shell of material, one would expect to see a radial profile like that of the northeast hemisphere.  Now if one adds a clump interior to the shell, the radial profile would include a ``bump'' at some inner radius, as seen in the southwest hemisphere. The radius at which this bump appears gives no indication of the true position of the clump, which could be almost anywhere along the line-of-sight in the remnant.  Our simple deprojection analysis ignores this fact, and assumes that the emission from the clump is spread over an inner shell.  Therefore, the density derived for these inner portions is clearly an overestimate for the interior of the remnant.  On the other hand, it may be an underestimate for the true density of the clump itself.  We believe that the clumpy region in the southwest of \\e05 is a result of enhanced density in or a deeper penetration of the reverse shock into a portion of the ejecta shell. This could be caused by a small region of enhanced ambient density or by intrinsic asymmetry in the explosion process itself.  Because our deprojection analysis revealed the shell-like structure of this remnant, we can be confident that the image fits, which assumed such a structure, produced meaningful results.  The fit results give us a constraint on the evolutionary state of the remnant if we identify the outer edge of the X-ray emitting shell with the location of the blast wave, and the inner edge with the location of the reverse shock.  In the context of self-similar models for the evolution of young SNRs (e.g., Truelove \\& McKee 1999), the observed ratio of these radii (1.185--1.192, including errors), corresponds to a particular point (or range of points) on the scaled radius vs.\\ time curve. These curves depend on the assumed radial density distribution of the ejecta; we consider the cases of constant and exponential profiles. From the constraint on the scaled radius vs.\\ time curve and the physical radius of the remnant, we obtain a relation between the ejected mass and the density of the ambient medium.  Assuming a Chandrasekhar mass for the ejecta, the density is $\\sim$4--$5\\times10^{-2}$ cm$^{-3}$ for the constant density ejecta profile of \\citet{tm}, and $\\sim$0.7--$1\\times10^{-2}$ cm$^{-3}$ for the exponential density ejecta profile of \\citet{dc}. These estimates imply a low density environment and yield swept-up masses for hydrogen of 0.055--$0.068~{\\rm M{_\\sun}}$ or 0.0095--$0.013~{\\rm M{_\\sun}}$, respectively.  The evolutionary models also can provide estimates of the explosion energy and age of the remnant, but they rely on knowing the shock velocity at the current epoch, which is only constrained to be $>$3600 km s$^{-1}$. The explosion energy is $>$$0.15\\times 10^{51}\\,(V_s/3600\\,\\rm km\\, s^{-1})^2$ erg (constant density) and $>$$0.09\\times 10^{51}\\,(V_s/3600\\,\\rm km\\, s^{-1})^2$ erg (exponential). The age is $<$$860\\,(V_s/3600\\,\\rm km\\, s^{-1})^{-1}$ yr (constant density) and $<$$670\\,(V_s/3600\\,\\rm km\\, s^{-1})^{-1}$ yr (exponential).  \\subsection{Non-Thermal Continuum} \\e05 is believed to be situated in a low density environment.  There is the dynamical argument just given, which constrains the density to $n_{\\rm H} \\lesssim0.05$ cm$^{-3}$.  \\citet{tuohy} obtained a value of $n_{\\rm H} \\lesssim0.02$ cm$^{-3}$ through optical studies.  In addition, the remarkable symmetry of the outer shell of this SNR suggests the ambient density is much lower than that surrounding other remnants (given the same level of density fluctuations).  For example, we do not see breaks in the rim or an irregular shape in \\e05 as we do in Tycho \\citep{hwang02}, which has an ambient density of $n_{\\rm H} \\sim0.5$ cm$^{-3}$ \\citep{hughes}.  Our Case H deprojection yields a post-shock density of $3.6~{\\rm cm}^{-3}$ for hydrogen, which in turn yields a swept-up hydrogen mass of 7.4 ${\\rm M{_\\sun}}$.  This density, converted to its preshock value of $\\sim$1 cm$^{-3}$ assuming the strong shock compression factor of 4, is much too high to be consistent with the previous arguments. If we take the ambient density to be $n_{\\rm H}=0.05$ cm$^{-3}$, then the level of hydrogen thermal continuum produced would be approximately $3\\times 10^{-3}$ times lower than the {\\em actual} continuum level observed in the spectrum. Therefore we are led to conclude that there is a significant non-thermal component of continuum emission in this SNR.  In Table 4 we show the masses of H, O, Si, S, and Fe in the two outermost shells, derived from the densities obtained by our deprojection analysis under spherical geometry.  Although the masses derived from the non-thermal case are too high to be consistent with SNe Ia (0.76 M$_\\sun$ for Si), this is a limiting case since it assumes {\\em all} the electrons come from the partially ionized metals.  Therefore, it is an extreme upper estimate to the mass.  If even a small amount of hydrogen were included with the non-thermal emission, the densities, and hence masses, of the metals would decrease.  In addition, the clumpy nature of the SNR inferred from our image fits (\\S 3.2) suggests that the apparently diffuse emission (from the shell component in the image fits) may arise from a smaller volume than expected if it were uniformly spread throughout.  The knots modeled in the image fits produce $\\sim$20\\% of the X-ray events from the shell in the northeast half of the remnant, while occupying only $\\sim$2\\% of the shell volume of the diffuse emission.  If we assume that the diffuse component is made up of similar, though unresolved, knots with the same X-ray events-to-volume ratio as those fit directly, we find that only $\\sim$10\\% of the diffuse emission volume need be occupied by knots to produce the same emission.  This 10\\% filling factor translates to a $\\sim$30\\% reduction in the derived masses, bringing the values for the non-thermal case to a more reasonable level: about 0.2 M$_\\sun$ for Si.  Several other young, shell-like SNRs show evidence for non-thermal emission \\citep{koy2,allen,slane}, the most well-known being SN1006 \\citep{koy}.  The shell of SN1006 is dominated by non-thermal emission, showing a featureless spectrum that can be described by a power-law with a photon index of 2.95 \\citep{koy}.  The synchrotron cut-off model was also applied to SN1006 and a value of $6\\times10^{16}$ Hz for the cut-off frequency was obtained \\citep{reykeo}.  For \\e05 we find a higher power-law photon index ($\\alpha_p\\sim3.3$) and a lower synchrotron cut-off frequency ($1.6\\times10^{16}$ Hz).  Both of these results are consistent with \\e05 showing a steeper non-thermal spectrum in the 0.5--7 keV X-ray band than SN1006.  As to the overall intensity of non-thermal emission, these two remnants are quite similar as well, with \\e05 being about a factor of three more luminous than SN1006.  Other SNRs, such as Tycho, Kepler, and Cas A, are not dominated by non-thermal emission but are well-fit by a power-law at energies $\\gtrsim$10 keV (Allen et al. 1999).  In the case of \\e05, we get as good, if not a better, fit for the global spectrum with a non-thermal continuum.  Thus a non-thermal origin for the continuum emission from \\e05 is plausible.  We also applied the minimum-energy condition for the synchrotron emission (Longair 1994, pp.~292-296) from this remnant, assuming equal energy densities for the protons and electrons and a volume filling factor of one.  By plotting the extrapolated radio power-law against the cut-off power-law determined by the synchrotron cut-off model, we found that the latter begins to deviate from a straight power-law at $\\sim$10$^{14}$ Hz. We used this value as the maximum frequency in our calculation and note that varying this value did not change our results greatly. We obtained a minimum energy of $10^{48}$ ergs and magnetic field of 60 $\\mu$G.  The energy requirements on this component are safely below the available energy of $10^{51}$ ergs, but our magnetic field value is large.  However, the minimum energy condition (effectively equipartition between particles and the magnetic field) is unlikely to be applicable, as shown by Dyer et al.~(2001) for SN1006 from comparing TeV gamma-ray and non-thermal X-ray emission.  If we assume a magnetic field of $\\sim$ 10 $\\mu$G as they find, then the total energy required to explain the non-thermal emission in 0509-67.5 increases to $10^{49}$ ergs, still well below the available energy.  Using $\\sim$ 10 $\\mu$G for the magnetic field along with the fitted value of the roll-off frequency from the synchrotron cut-off spectral model, we find the maximum energy of electrons to be $\\sim$ 20 TeV.  The preceding arguments strongly support the picture that the forward shock of \\e05 is accelerating particles to relativistic energies. \\e05 is thus the first Magellanic Cloud SNR for which such an energetic cosmic ray component has been securely detected.  One task that remains to be done to make our work fully consistent is to explore the effects of modifications to the remnant dynamics due to efficient particle acceleration. As the fraction of the shock front's kinetic energy being diverted to the acceleration of relativistic particles grows, the thicknesses of the forward and reverse shocked regions tend to decrease \\citep{dec2}. Such changes could affect our estimate of the ambient medium density based on the thickness of the X-ray emitting shell.  In the absence of a specific model for \\e05, it is not possible to estimate a correction factor for the ambient density. Although beyond the scope of this work, development of such a model would be a valuable next step toward determining the efficiency of cosmic ray acceleration in \\e05."
    },
    "0403/astro-ph0403669_arXiv.txt": {
        "abstract": "We present {\\it Chandra} and {\\it XMM-Newton} X-ray data of NGC\\,5253, a local starbursting dwarf elliptical galaxy, in the early stages of a starburst episode. Contributions to the X-ray emission come from discrete point sources and extended diffuse emission, in the form of what appear to be multiple superbubbles, and smaller bubbles probably associated with individual star clusters. {\\it Chandra} detects 17 sources within the optical extent of NGC\\,5253 down to a completeness level corresponding to a luminosity of $1.5 \\times 10^{37}\\ergs$. The slope of the point source X-ray luminosity function is $-0.54 \\pm^{0.21}_{0.16}$, similar to that of other nearby dwarf starburst galaxies. Several different types of source are detected within the galaxy, including X-ray binaries and the emission associated with star-clusters.  Comparison of the diffuse X-ray emission with the observed H$\\alpha$ emission shows similarities in their extent. The best spectral fit to the diffuse emission is obtained with an absorbed, two temperature model giving temperatures for the two gas components of $\\sim 0.24\\kev$ and $\\sim 0.75\\kev$. The derived parameters of the diffuse X-ray emitting gas are as follows: a total mass of $\\sim 1.4 \\times 10^{6} f^{1/2}\\msun$, where $f$ is the volume filling factor of the X-ray emitting gas, and a total thermal energy content for the hot X-ray emitting gas of $\\sim 3.4 \\times 10^{54}f^{1/2}$~erg. The pressure in the diffuse gas is $P/k\\sim 10^{6} f^{-1/2}$~K~cm$^{-3}$.  We find that these values are broadly commensurate with the mass and energy injection from the starburst population.  Analysis of the kinematics of the starburst region suggest that the stellar ejecta contained within it can escape the gravitational potential well of the galaxy, and pollute the surrounding IGM. ",
        "introduction": "NGC\\,5253 is a starbursting dwarf elliptical galaxy which lies at a distance of $\\sim 3.15$~Mpc (Freedman \\etal\\ 2001) and with an inclination of 67$^{\\circ}$ (the mean value quoted by the Lyon/Meudon Extra-galactic Database [LEDA]). The signature of Wolf-Rayet stars have been detected in the nucleus of the galaxy (Schaerer \\etal\\ 1997, Walsh \\& Roy 1987) implying that this is a young starburst and as such allows an opportunity for the study of the starburst phenomenon in its earlier stages of development and its effect on galaxy evolution in the local Universe. In fact, Rieke, Lebofsky \\& Walker (1988) went so far as to classify it as one of the youngest starbursts known. Dwarf galaxies, as the basic building blocks in the hierarchical merging cosmology scenario, are likely to have harboured the earliest sites of star-formation in the Universe and so their study in the local Universe can give insight into the evolution of such objects at high redshift. Observations of local edge-on starburst galaxies (Strickland \\etal\\ 2000; Weaver 2001) are presenting a picture of kpc-scale, soft X-ray emitting, bipolar outflows in the form of galactic winds transporting mass, newly synthesised heavy elements and energy into the intergalactic medium (IGM). A similar outflow is seen in the dwarf starburst galaxy NGC\\,1569 (Martin, Kobulnicky \\& Heckman 2002), while the dwarf NGC\\,4449 shows what may be the beginnings of a galactic wind emerging from an extended superbubble (Summers \\etal\\ 2003). The same situation seems to be the case in the NGC\\,3077 dwarf (Ott, Martin \\& Walter 2003). These winds result from the pressure driven outflows along these galaxy's minor axis produced from the efficient thermalization of the mechanical energy from the supernovae (SN) explosions and stellar winds of the massive stars in their OB associations and super star-clusters (SSC).  NGC\\,5253 being both inclined to our line-of-sight and young presents a less clearly observable picture. Its youth means that it is unlikely to have developed a superwind and the diffuse X-ray emission observed is much more likely to be associated with multiple superbubbles around its OB associations and SSC, as suggested from earlier {\\it ROSAT} observations of this galaxy (Strickland \\& Stevens 1999).  Observations of NGC\\,5253 at radio wavelengths have shown it to have a very flat centimetre-wavelength continuum which is indicative of thermal emission from H{\\small II} regions, with only low levels of synchrotron emission from supernova remnants (Beck \\etal\\ 1996; Turner, Ho \\& Beck 1998). Ferrarese \\etal\\ (2000) quote a low metallicity of $0.2 Z_{\\odot}$, based on the work of Webster \\& Smith (1983). This figure suggests that a large amount of metal enrichment has not occurred in the central H{\\small II} regions, and may be further evidence of the youth of the starburst region. The total and H{\\small I} masses for the galaxy are calculated to be $6.4 \\times 10^{8}\\msun$ and $8.3 \\times 10^{7}\\msun$ from 21~cm observations (Reif \\etal\\ 1982), corrected for the difference in assumed distance to NGC\\,5253. Later VLA 21~cm observations show NGC\\,5253 to be peculiar, in so much as its neutral hydrogen appears to rotate about the optical major axis of the galaxy (Kobulnicky \\& Skillman 1995). However, CO observations (Turner, Beck \\& Hurt 1997; Meier, Turner \\& Beck 2002) detect sources which are coincident with the dust lane seen to the SE of the nucleus of the galaxy in optical images, and these CO clouds appear to be infalling into NGC\\,5253. Consequently, the dynamical situation in NGC\\,5253, rotation or infall, is not completely clear.  Infrared emission is detected from what seems to be a highly obscured, massive ($10^{5} -10^{6}$ stars), small ($\\sim$ a few pc diameter) and very young globular cluster in the central starburst region of the galaxy (Gorjian, Turner \\& Beck 2001).  This source contributes $\\sim 50\\%$ of the total observed infrared luminosity of the galaxy.  A summary of some of the properties of NGC\\,5253 is given in Table~1, along with estimates of the current star-formation rate of the galaxy estimated from far-infrared and H$\\alpha$ fluxes.  Caldwell \\& Phillips (1989), using broadband multicolour (UBV) CCD images, narrow-band H$\\alpha$ and long-slit Ca{\\small II} triplet spectra, determined that NGC\\,5253 had undergone an increased rate of star-formation throughout the galaxy about $10^{8} - 10^{9}$~yr ago, as suggested by the presence of over 100 star clusters of a similar age lying outside the nuclear region. They also established that the intense star-formation now occurring in the nuclear region of the galaxy began around 10~Myr ago, giving an upper limit for the age of superbubbles associated with it.  More recently, Tremonti et al. (2001) analysed several clusters in NGC\\,5253, determining ages in the range of 1 to 8~Myr. Calzetti et al. (1997) also noted two other clusters with older ages, determined to be between 10 and 50~Myr.  The two very bright SN explosions seen in this galaxy in 1895 and 1972 are amongst the apparently brightest recorded extra-galactic SNe (Ardeberg \\& de Groot 1973) and may be associated with the large scale star-formation, although neither was particularly centrally located. The Fabry-Perot H$\\alpha$ images of Marlowe \\etal\\ (1995) show that the ionized gas in the galaxy has a complex distribution, consisting of a bright central region extending $\\sim 30{''}$ from the centre in a N-S direction, embedded within a system of loops and filaments, with an overall extent of $\\sim 2{'}$. To the NW and WSW of the centre are 2 large superbubbles with diameters $\\sim 1{'}$, expanding at $\\sim 35\\kms$, while to the east side of the galaxy, a pair of radially oriented filaments are detected. This latter observation is suggestive of the rupture of a superbubble blown by the stellar activity in the central regions.  NGC\\,5253 has previously been observed at X-ray wavelengths with the {\\it Einstein} IPC (Fabbiano, Kim \\& Trinchieri 1992) and both the {\\it ROSAT} PSPC (Martin \\& Kennicutt 1995; Stevens \\& Strickland 1998) and HRI (Strickland \\& Stevens 1999). The {\\it Einstein} IPC and {\\it ROSAT} PSPC results showed soft thermal X-ray emission from what appeared to be an extended source, suggestive of a luminous superbubble. The better spatial resolution of the {\\it ROSAT} HRI instrument detected a complex of at least five X-ray sources, that could be young superbubbles blown by individual young star clusters in the starburst region. The total X-ray luminosity determined from these observations was similar and lies in the range of $(2.4 - 4.1) \\times 10^{38}\\ergs$, depending on the value adopted for the absorbing column density.  In Section 2 we describe the {\\it Chandra} and {\\it XMM-Newton} observations. The X-ray emission from the whole galaxy, the point sources and the diffuse emission are discussed in Section~3. Section~4 contains a more general discussion of the morphology of the X-ray emission and its relationship to emission from other wavebands, along with a determination of the dynamics of the expanding superbubbles of the central region and their effect on NGC\\,5253. Our main conclusions are summarised in Section~5.  \\begin{figure*} \\vspace{13cm} \\special{psfile=n5253_fig1_alt.ps hoffset=-40 voffset=398 hscale=70 vscale=70 angle=-90} \\caption{Low resolution (smoothed using a Gaussian with FWHM of 4 pixels $\\sim 2^{''}$), background subtracted {\\it Chandra} image, in the $0.3 - 8.0\\kev$ energy band, showing the full ACIS-S3 chip field of view, marked with the 31 point sources (Table~2). The image is an inverted log greyscale with the flux density ranging from $7.63 \\times 10^{-14}\\ergs$~cm$^{-2}$~arcmin$^{-2}$ to $9.72 \\times 10^{-11}\\ergs$~cm$^{-2}$~arcmin$^{-2}$. The $D_{25}$ ellipse is also shown. North is to the top and East is to the left.} \\label{sources} \\end{figure*}  \\begin{table} \\caption{Details of NGC\\,5253. The distance is from Freedman \\etal\\ (2001), other data is from the LEDA and NED databases. The infrared luminosity $L_{IR}$ is calculated using the 12, 25, 60 and $100\\mu$m IRAS fluxes (Sanders \\& Mirabel 1996) and the H$\\alpha$ luminosity is from Marlowe \\etal\\ (1997, corrected for distance). The corresponding star-formation rates (SFR) for both the IR and H$\\alpha$ luminosities are calculated using the conversion formulae in Kennicutt (1998).} \\begin{center} \\begin{tabular}{ll} \\hline Parameter & Value \\\\ \\hline Classification\\ \\ \\ \\ \\  & Im pec, H{\\small II} \\\\ Distance & 3.15Mpc \\\\ RA (J2000) & 12 39 56.0 \\\\ Dec (J2000) &$-31$ 38 36\\\\ $D_{25}$ ellipse & $4.8' \\times 1.9'$\\\\ Position angle & $45^\\circ$\\\\ $m_B$ & 10.75\\\\ $B-V$ & 0.43\\\\ $L_{IR}$ & $4.5\\times 10^{42}\\ergs$\\\\ $L_{\\rm {H}\\alpha}$ & $2.6\\times 10^{40}\\ergs$\\\\ SFR rate (IR) & $0.2\\msunyr$\\\\ SFR rate (H$\\alpha$) & $0.2\\msunyr$\\\\ \\hline \\end{tabular} \\end{center} \\end{table}   \\begin{table} \\begin{center} \\caption{Positions and count rates of the 31 sources detected in the NGC\\,5253 {\\it Chandra} S3 chip data. Column 1 is the source number ordered in increasing R.A.. Columns 2 and 3 give the R.A. and Dec. of each source and Column 4 lists their background subtracted count rates.} \\begin{tabular}{|c|c|c|c|} \\hline Source & RA (h m s) & Dec ($^{\\circ} \\; {'}\\; {''}$) & Count Rate \\\\ & & & ($\\times 10^{-3}$ cts s$^{-1}$) \\\\ \\hline 1 & 13 39 37.73 & -31 40 15.8 & $0.77\\pm0.15$ \\\\ 2 & 13 39 39.16 & -31 36 06.9 & $0.32\\pm 0.09$ \\\\ 3 & 13 39 39.17 & -31 36 09.2 & $0.42\\pm 0.10$ \\\\ 4 & 13 39 43.15 & -31 39 20.2 & $1.07\\pm 0.16$\\\\ 5 & 13 39 45.12 & -31 36 58.2 & $0.37\\pm 0.10$\\\\ 6 & 13 39 45.53 & -31 42 06.1 & $0.39\\pm 0.10$\\\\ 7 & 13 39 46.69 & -31 42 28.0 & $0.43\\pm 0.11$\\\\ 8 & 13 39 47.27 & -31 41 40.7 & $1.95\\pm 0.22$\\\\ 9 & 13 39 50.66 & -31 39 20.7 & $1.50\\pm 0.18$\\\\ 10 & 13 39 51.17 & -31 38 40.8 & $0.32\\pm 0.09$\\\\ 11 & 13 39 51.62 & -31 40 16.2 & $0.48\\pm 0.11$\\\\ 12 & 13 39 51.96 & -31 41 43.9 & $4.03\\pm 0.30$\\\\ 13 & 13 39 52.07 & -31 38 59.7 & $0.19\\pm 0.07$\\\\ 14 & 13 39 53.36 & -31 43 43.9 & $1.42\\pm 0.20$\\\\ 15 & 13 39 53.80 & -31 38 32.7 & $0.66\\pm 0.13$\\\\ 16 & 13 39 55.39 & -31 38 32.8 & $0.80\\pm 0.17$\\\\ 17 & 13 39 55.69 & -31 38 31.5 & $1.02\\pm 0.20$\\\\ 18 & 13 39 55.85 & -31 38 22.8 & $1.03\\pm 0.20$\\\\ 19 & 13 39 56.04 & -31 38 24.8 & $1.43\\pm 0.23$\\\\ 20 & 13 39 56.15 & -31 38 35.9 & $0.58\\pm 0.16$\\\\ 21 & 13 39 56.35 & -31 38 20.5 & $8.22\\pm 0.43$\\\\ 22 & 13 39 56.45 & -31 38 25.8 & $1.53\\pm 0.21$\\\\ 23 & 13 39 56.67 & -31 42 45.3 & $1.57\\pm 0.21$\\\\ 24 & 13 39 57.43 & -31 37 32.5 & $1.50\\pm 0.06$\\\\ 25 & 13 39 59.27 & -31 37 50.2 & $0.33\\pm 0.09$\\\\ 26 & 13 40 02.19 & -31 37 19.1 & $0.96\\pm 0.15$\\\\ 27 & 13 40 03.40 & -31 37 14.5 & $0.78\\pm 0.13$\\\\ 28 & 13 40 03.84 & -31 44 03.7 & $1.25\\pm 0.21$\\\\ 29 & 13 40 04.29 & -31 37 30.3 & $0.41\\pm 0.10$\\\\ 30 & 13 40 04.97 & -31 42 03.0 & $0.24\\pm 0.08$\\\\ 31 & 13 40 10.81 & -31 40 33.2 & $0.29\\pm0.09$\\\\ \\hline \\end{tabular} \\end{center} \\end{table}   \\begin{figure*} \\vspace{13cm} \\special{psfile=n5253_fig2.ps hoffset=-20 voffset=375 hscale=70 vscale=70 angle=-90} \\caption{Low resolution (smoothed using a Gaussian with FWHM of 2 pixels $\\sim 8^{''}$), {\\it XMM-Newton}, background subtracted and exposure corrected, mosaiced image of the data from the EPIC MOS1, MOS2 and PN cameras, in the $0.3 - 8.0\\kev$ energy band. The image is an inverted log greyscale with the flux density ranging from $5.2 \\times 10^{-14}\\ergs$~cm$^{-2}$~arcmin$^{-2}$ to $2.1 \\times 10^{-12}\\ergs$~cm$^{-2}$~arcmin$^{-2}$, with North to the top and East to the left. The $D_{25}$ ellipse is also shown. The X-ray emission is extended towards the WSW.} \\label{mosaic} \\end{figure*} ",
        "conclusions": "In summary, we have presented an X-ray analysis of the important dwarf elliptical, starburst galaxy NGC\\,5253 using data from both the {\\it Chandra} and {\\it XMM-Newton} satellites. We detected X-ray emission from 31 discrete point sources in the {\\it Chandra} ACIS-S3 chip data with 17 of these lying within the optical extent of the galaxy, as measured by the $D_{25}$ ellipse. Some of these can be clearly identified as X-ray binaries while others seem to be associated with the emission from super star-clusters rather than individual point sources, as they are extended compared to the {\\it Chandra} PSF. The 5 sources in this category (sources $16 - 20$) all lie in the central region of the galaxy and are associated (particularly in the case of sources 17, 18 and 19) with the most intense regions of emission seen on the H$\\alpha$ image. The unabsorbed luminosities of these sources in the $0.3 - 8.0\\kev$ energy band, range from $(0.25 - 1.50) \\times 10^{38}\\ergs$.  The majority of the X-ray emission is confined to the central region of the galaxy where the intense, current burst of star-formation is occurring. The extent of the emission also correlates well with that of the H$\\alpha$ emission and both are confined to a region of diameter $\\sim 1$~kpc in the centre of the galaxy. Both of these emissions are embedded within a spherical distribution of H{\\small I} that extends out to $\\sim 2$~kpc from the centre.  The diffuse X-ray emission in NGC\\,5253 requires two thermal and one power-law component to best fit its spectrum, suggesting that the hot gas is a multi-phase environment with unresolved point sources, probably lower luminosity X-ray binaries, embedded in it. The fitted gas temperatures are $0.24 \\pm 0.01\\kev$ and $0.75 \\pm 0.05\\kev$ for the soft and medium components respectively and their respective absorption corrected luminosities in the $0.37 - 6.0\\kev$ band are $(2.19 \\pm^{0.16}_{0.15}) \\times 10^{38}\\ergs$ and $(1.75 \\pm^{0.15}_{0.14}) \\times 10^{38}\\ergs$. This emission is seen to be extended in several directions in the {\\it Chandra} data but this is not confirmed by the {\\it XMM-Newton} data. However, the extent is not as great as the $\\sim 1$~kpc shells reported to the NW and WSW by Marlowe \\etal\\ (1995), in the H$\\alpha$ emission. These shells are expanding with velocities of $\\sim 35\\kms$ and hence have required $\\sim$10~Myr to have evolved. NGC\\,5253 is proving to be a rather enigmatic object though, as our analysis of the dynamics of the hot gas are not consistent with the expansion velocity and age of these H$\\alpha$ shells.  The hot X-ray emitting gas has a total thermal energy content of $3.4 \\times 10^{54}$~erg and a total mass of $1.4 \\times 10^{6}\\msun$. From values in the literature, we assume an age for the starburst responsible for this hot gas of 5~Myr and a value for the density of the ambient ISM of $n_{0} = 5$~cm$^{-3}$. The radius of the region in which most of the diffuse X-ray and H$\\alpha$ emission is confined is $0.53$~kpc. Application of standard superbubble models to NGC\\,5253's starburst, assuming the age and density above, result in a predicted radius and expansion velocity for a single superbubble of $\\sim 400$~pc and $\\sim 50\\kms$, figures which are comparable to those observed.  The main problem here would seem to lie in the fact that no clear bipolar outflow is observed in NGC\\,5253 and no allowance has been made for the presence of multiple superbubbles in the central region. If the central starburst is young, the bubbles blown around individual star clusters will overlap and impede each others expansion confining the emission to a smaller region than expected and reducing the observed expansion rate. The H$\\alpha$ shells, reported by Marlowe \\etal\\ (1995), to the West could be the result of the superbubbles closest to the edge of the central region being able to expand more freely. The lower velocity seen could result from less stars contributing to the energy input in these regions and/or the outflow having experienced some recent deceleration on encountering higher density material such as the dark clouds reported in the West of the galaxy (Hunter 1982). However, our radial surface brightness profile does not show any detectable X-ray emission beyond $\\sim 0.5$~kpc.  NGC\\,5253 is a very complex object, which shows that standard superbubble models are difficult to apply to dwarf starbursts galaxies. However, the energy injection rate into the galaxy would seem sufficient to allow the expanding hot gas to escape the gravitational potential well of NGC\\,5253 and its relatively small ($\\sim 1.8$~kpc along the minor axis) H{\\small I} halo would also offer little resistance to this. Thus, it seems likely then that NGC\\,5253 will lose metal-enriched material, mass and energy as a result of its current bout of star-formation."
    },
    "0403/astro-ph0403496_arXiv.txt": {
        "abstract": "We report the strength of seed magnetic flux of accretion disk surrounding the PopIII stars. The magnetic field in accretion disk might play an important role in the transport of angular momentum because of the turbulence induced by Magneto-Rotational Instability (MRI). On the other hand, since the primordial star-forming clouds contain no heavy elements and grains, they experience much different thermal history and different dissipation history of the magnetic field in the course of their gravitational contraction, from those in the present-day star-forming molecular clouds. In order to assess the magnetic field strength in the accretion disk of PopIII stars, we calculate the thermal history of the primordial collapsing clouds, and investigate the coupling of magnetic field with primordial gas. As a result, we find that the magnetic field strongly couple with primordial gas cloud throughout the collapse, i.e. the magnetic field are frozen to the gas as far as initial field strength satisfies $ B\\lesssim 10^{-5}(n_{\\rm H}/10^3~{\\rm cm^{-3}})^{0.55}~{\\rm G}$. ",
        "introduction": "PopIII stars are considered to have very important impacts on the thermal history and the chemical evolution of the universe. Recent observations  of the polarization of Cosmic Microwave Background (CMB) photons by WMAP(Wilkinson Microwave Anisotropy Probe) \\citep{Kogut03} revealed that the optical depth of the universe by Thomson scattering is fairly large ($\\tau = 0.17 \\pm 0.04$). This result tells that reionization epoch is earlier ($z = 17 \\pm 5$) than expected from the observations on Gunn-Peterson trough \\citep{Becker01,Fan02,White03}. On the other hand, the results of theoretical calculations of structure formation \\citep{Sokasian03} indicate that such early reionization seems to be difficult solely by the Pop II stars, but it might be possible if the ionizing photons from PopIII stars with top heavy initial mass function (IMF) is taken into account. Thus, the mass (or IMF) of PopIII star have great significance on the reionization epoch of the universe. In addition, some amount of metals are found at high redshift intergalactic matter by the observation of high redshift QSO absorption line systems\\citep{songaila01,ma02,vladilo02}. This means significant fraction of baryons are already processed in the stars by $z=5$ \\citep{songaila01}. PopIII stars also should be responsible for such metal pollution of the early universe, and the abundance pattern of heavy elements depends on the mass of PopIII stars. Thus, typical mass (or IMF) of the PopIII stars is a key quantity also for the chemical evolution of the universe.  In order to assess the typical mass of PopIII stars, typical scale of the prestellar core (or fragments of primordial gas) should be evaluated as a first step. Several authors have studied on this issue by simple one-zone approach before middle of '90s \\citep{MST69,Hut,Carl,PSS83,SUN96,USNY96,Puy}. Recently, multi-dimensional numerical simulations of fragmentation of primordial gas have been performed intensively by several authors\\citep{nakamura99,bromm99,abel00,nakamura01,bromm02}, and they find that the mass of PopIII prestellar core is quite massive ($\\sim 10^{3-4}M_\\odot$).  Further evolution of prestellar core was investigated by \\cite{omukai98}, and it is found that the collapse proceeds in a run-away fashion and converges to Larson-Penston similarity solution \\citep{Larson69,Penston69} with $\\gamma\\simeq 1.1$. They also find the mass accretion rate is also very large compared to the present-day forming stars, although spherical symmetry is assumed in their numerical calculation. In reality, however, prestellar cores have some amount of angular momentum, which prevent the mass from accreting onto the protostars. Consequently, accretion disks surrounding the protostars are expected, and some mechanisms that transport their angular momentum are required in order to enable the mass accretion onto protostars.  There are a few possibilities of the angular momentum transport such as gravitational torque by the nonaxisymmetric structures in the accretion disk, the interaction among the fragments of the disk \\citep{stone00,bodenheimer00}, and the turbulent viscosity triggered by Magneto-Rotational Instability (MRI) \\citep{hawley92,sano98,sano01}. All of these mechanisms are regarded to be important for present-day star formation. Thus, it is worth to investigate these processes for primordial case. Among these possibilities, we concentrate on the last one, the MRI induced turbulence.  In order to activate MRI, the initial magnetic field strength in the accretion disk should be larger than a critical value. Otherwise the magnetic field is dissipated before the field is amplified by MRI \\citep{TM04,TB04}. Thus, it is important to assess the ``initial'' magnetic field strength in the accretion disk. As was discussed by \\citet{nakano86} for preset-day case, magnetic field could be dissipated while prestellar core collapses. However, the dissipation processes strongly depend on the components of the gas as well as the the temperature evolution in the course of the collapse. On the other hand, the temperature of the collapsing primordial gas is much higher than that of the present-day case \\citep{omukai00}, which might bring about different dissipation history of the magnetic field.  In this paper, we investigate the dissipation of magnetic field in collapsing primordial gas cloud. Consequently, we obtain the initial magnetic field strength in the accretion disk of PopIII stars. In the next section, formulations are described. In \\S\\ref{results}, results of our calculations are shown, and the importance of MRI in PopIII star formation is discussed in \\S\\ref{discussion}. Final section is devoted to summary. ",
        "conclusions": "\\label{discussion} The condition for which the MRI can grow at the PopIII accretion disk is investigated by \\citet{TB04}. This condition requires that the MRI growth timescale is shorter than the diffusion timescale. Thus, there is the minimum field strength in the disk for the MRI to be driven. According to \\citet{TB04}, this minimum field strength in the disk becomes  \\begin{eqnarray} B&=&1.1\\times 10^{-4}{\\rm G} \\left(\\frac{m_*}{10\\ M_\\odot}\\right)^{1/4} \\left(\\frac{T}{10^4\\ {\\rm K}}\\right)^{-3/4} \\left(\\frac{\\ln\\Lambda}{10}\\right)^{1/2}\\nonumber\\\\ &&\\times\\left(\\frac{\\rho_{\\rm disk}} {5\\times 10^{-10}\\ {\\rm g\\;cm^{-3}}}\\right)^{1/2} \\left(\\frac{r}{600\\ R_\\odot}\\right)^{-3/4}, \\label{eq:discon} \\end{eqnarray} where $m_*$ is the stellar mass, $r$ is the radius from the cloud center, $\\rho_{\\rm disk}$ denotes the density of the accretion disk and $\\ln \\Lambda$ is the Coulomb logarithm. Characteristic values of these parameters are taken from Figure 3 in \\citet{TB04}.  Hence, it is concluded that if the initial field strength in prestellar core is at least $\\gtrsim 10^{-10}\\ {\\rm G}$ at $n_{\\rm H}=10^3\\ {\\rm cm^{-3}}$, the transport of angular momentum could be driven by the turbulence due to the MRI in the accretion disk.  In fact, there are the models for the generation of initial seed magnetic field in the universe \\citep{pudritz89,kulsrud97,widrow02,langer03}. Among these models, \\citet{langer03} propose the generation mechanism of magnetic field based on the radiation force around very luminous objects such as QSOs. According to their study, the magnetic field $\\sim 10^{-11}-10^{-12}\\ {\\rm G}$ is generated in the inter galactic matter. Consequently, field strength is amplified to $\\sim 10^{-7}-10^{-8}\\ {\\rm G}$ when the clouds collapses to $n_{\\rm H}=10^3\\ {\\rm cm^{-3}}$, which is the initial condition of our analysis. Thus, in this case, the MRI can be driven and it could be the possible mechanism of angular momentum transport. On the other hand, most of the other seed field generation mechanism predict $B \\lesssim 10^{-19}\\ {\\rm G}$, which is too small to drive MRI. Thus, MRI may not be important for the formation of very first stars, since the number of sources which provide anisotropic radiation field should be very small at the epoch of very first stars."
    },
    "0403/astro-ph0403482_arXiv.txt": {
        "abstract": "We report the discovery of only the fourth massive WO star to be found in the Milky Way, and only the seventh identified within the Local Group.  This has resulted from the first observations made in a programme of follow-up spectroscopy of candidate emission line stars from the AAO/UK Schmidt Southern Galactic Plane H$\\alpha$ Survey.  The optical spectrum of this star, to become WR~93b in the Catalogue of Galactic Wolf-Rayet stars, is presented and described.  WR~93b is classified as WO3 and is shown to be highly reddened ($E_{B-V} = 2.1\\pm0.1$). A recombination line analysis of the emission lines yields the abundance ratios C/He = 0.95 and O/He = 0.13 (by number). Comparisons at near infrared wavelengths of reddening corrected photometry between WR~93b and both of Sand~2 (WO3, $D = 49$~kpc) and Sand~5 (WO2, $D = 1.75$~kpc) yields a consistent distance to WR~93b of 3.4~kpc.  Positioned at Galactic co-ordinates $\\ell = 353.27^{\\rm o}$, $b = -0.85^{\\rm o}$, the star is most likely located in the Scutum-Crux Arm of the inner Milky Way. We note that none of the four Galactic WO stars lies significantly beyond the Solar Circle (with two well inside).  Estimation of the wind terminal velocity in WR~93b at 5750~km~s$^{-1}$ makes this star the current wind speed record holder among all non-degenerate stars. ",
        "introduction": "In this paper we present the discovery of a particularly rare type of emission line star made in the initial phase of a long-term spectroscopy programme which aims to confirm and provide preliminary classification of spatially-unresolved line-excess objects contained within the AAO/UK Schmidt Southern Galactic Plane H$\\alpha$ Survey (Parker et al 2003).  The star in question has been revealed as only the fourth massive WO star to be found in the Galaxy -- bringing the total identified within the Local Group to just seven.  When it is assimilated into the next revision of the catalogue of massive Wolf-Rayet stars in the Galaxy, compiled and maintained by van der Hucht (2001), its position in the sky will earn it the designation WR 93b.  Hereafter we refer to this newly discovered star by this name.  The WO stars are the most chemically extreme Wolf-Rayet (WR) stars, whose spectra are dominated by high-excitation oxygen and carbon lines.  Objects in this elite group of stars are viewed as plausible progenitors for extreme, chemically-peculiar Type Ib/c supernovae (Woosley, Heger \\& Weaver 2002) and some GRBs (e.g. Schaefer et al 2003).  The history of this presently rare type of object begins with the last generation of galactic emission line surveys: Sanduleak (1971) presented a list of five WR stars, two in the Magellanic Clouds, which had strong O~{\\sc vi} 3811,34~\\AA\\ emission features that had previously been found only amongst planetary nebula nuclei.  Barlow \\& Hummer (1982) proposed that one of these five stars was indeed a PN central star but that the other four corresponded to an advanced stage of evolution of massive stars, beyond the WC phase, and classified them as WO Wolf-Rayet stars. Since this time, two more have been added to this grouping (MS 4, by Smith, Shara \\& Moffat 1990 and DR~1 in the dwarf irregular galaxy IC 1613 by Kingsburgh \\& Barlow 1995). In general terms, discoveries of Galactic WR stars have been assisted greatly by optical objective-prism surveys (Stephenson \\& Sanduleak 1971, MacConnell \\& Sanduleak 1970) and narrow-band imaging surveys (most recently, Shara et al. 1999). The result we present here uses the latter technique in the red part of the spectrum, and serves as an encouragement that much remains to be discovered.  In the next section we introduce the AAO/UK Schmidt Southern Galactic Plane H$\\alpha$ Survey and describe the data it contains relevant to the newly-discovered WO star and its locale.  We then describe how the WO star came to be included in our initial programme of follow-up spectroscopy using the AAO/UK Schmidt multi-fibre spectroscopy facility, 6dF.  This is followed, in Section 3, by presentation of both the 6dF data obtained in May/June 2003 (\\S 3.1), and flux-calibrated long-slit spectra obtained at the William Herschel Telescope (WHT) in August 2003 (\\S 3.2).  We are then in a position to classify the WO star and determine its reddening (\\S 4).  We also present estimates of C/He and O/He abundance ratios deduced from the emission lines (\\S 5).  Finally, by comparison with other WO stars at known distances, we derive the distance to WR~93b (\\S 6): at 3.4~kpc, and along a line of sight passing within a few degrees of the Galactic Centre, it seems likely the new WO star is associated with the Scutum-Crux spiral arm (Russeil 2003) -- without a doubt it is well inside the Solar Circle. We close in \\S 7 with comment on the 4-strong group of galactic WO stars and on the prospects for future discoveries. ",
        "conclusions": "\\begin{tabular}{lll} \\hline Parameter & value & comment\\\\ \\hline Position (RA,Dec) & 17 32 03.30 $-$35 04 32.5 & J2000, from SHS database \\\\ $R$, $I$ magnitudes & 14.7, 12.5 & from SHS database\\\\ $J$, $H$, $K$ magnitudes & 11.33$\\pm$0.03, 10.56$\\pm$0.04, 10.17$\\pm$0.04 & from 2MASS database \\\\ E($B-V$) & 2.1$\\pm$0.1 & $R=3.1$ gives A$_V \\simeq 6.5$  \\\\ spectral type & WO3 & criteria place WR 93b closer \\\\ &     & to WO2/3 boundary than Sand~2 \\\\ wind terminal velocity & 5750 km s$^{-1}$ & \\\\ abundance ratios: & C/He $=$ 0.95 & by number \\\\ & O/He $=$ 0.13 & \\ `` \\ \\ ``     \\\\ & (C+O)/He $=$ 1.08 & \\ `` \\ \\ ``     \\\\ distance      &  3.4$\\pm$0.3 kpc & \\\\ \\hline \\end{tabular} \\end{table*}  Adopting a distance of 49~kpc to the LMC (see e.g.Gibson 2000), the difference of 5.7 magnitudes with respect to Sand~2 yields a distance of 3.5~kpc to WR~93b. Two recent estimates of the distance to Be~87, the cluster hosting Sand~5, are respectively 1.6~kpc (Massey, DeGioia-Eastwood \\& Waterhouse 2001) and 1.9~kpc (Kn\\\"odlseder et al 2002): we adopt the intermediate value of 1.75~kpc.  The NIR magnitude difference of $-$1.43, combined with this distance then yields a distance to WR~93b of 3.4~kpc.  This is very satisfying agreement. But, given the systematic uncertainty contained within the assumption of shared stellar parameters, we should anticipate an error on 3.4~kpc of up to 10 \\% (equivalent to $\\sim$0.2 offset in $M_V$).  The Galactic co-ordinates of WR~93b are $\\ell = 353.27\\deg$, $b = -0.85\\deg$, placing it less than 7\\deg from the Galactic Centre line of sight, and outside the strips surveyed for Wolf-Rayet stars by Shara et al (1999).  At a distance of 3.4~kpc the new WO star would be located very close to the structure Russeil (2003) identifies as the Scutum-Crux arm, well inside the long-established Sagittarius-Carina arm at $\\sim$2~kpc in the same direction.  For a  necessarily young object like a Wolf-Rayet star this is a perfectly acceptable place to be. Projected onto the plane of the sky, the nearest H~{\\sc ii} region (a signature of continuing star formation) listed by Russeil (2003, Table 3) is about half a degree away from WR~93b at $\\ell = 353.43\\deg$, $b = -0.368\\deg$.  At 3.4~kpc this angular separation converts to a length of about 30~pc.  The kinematic distance to this H{\\sc ii} region, which is `compact' and without a known optical counterpart, has been estimated from Caswell \\& Haynes' (1987) measurement of its radio recombination line radial velocity.  In the framework adopted by Russeil (2003), wherein the Galactic Centre is 8.5~kpc away, the distance to this H~{\\sc ii} region is required to be $3.5^{+0.6}_{-0.7}$ kpc.  This is consistent with our estimate of the distance to the WO star. If the Galactic Centre is taken to be 7.94$\\pm$0.42~kpc as determined recently by Eisenhauer et al (2003), this revises the kinematic distance downwards a little to 3.3~kpc -- leaving the quality of agreement unchanged."
    },
    "0403/astro-ph0403161_arXiv.txt": {
        "abstract": "We present a redshift $z=6.535$ galaxy discovered by its \\lya\\ emission in a 9180\\AA\\ narrowband image from the Large Area Lyman Alpha (LALA) survey.  The \\lya\\ line luminosity ($1.1\\times 10^{43} \\ergsec$) is among the largest known for star forming galaxies at $z\\approx 6.5$.  The line shows the distinct asymmetry that is characteristic of high-redshift \\lya. The $2\\sigma$ lower bound on the observer-frame equivalent width is $>530$\\AA.  This is hard to reconcile with a neutral intergalactic medium unless the \\lya\\ line is intrinsically strong {\\it and\\/} is emitted from its host galaxy with an intrinsic Doppler shift of several hundred $\\kms$.  If the IGM is ionized, it corresponds to a rest frame equivalent width $> 40$\\AA\\ after correcting for \\lya\\ forest absorption. We also present complete spectroscopic followup of the remaining candidates with line flux  $> 2\\times 10^{-17} \\ergcm2s$ in our $1200\\sqamin$ narrowband image.  These include another galaxy with a strong emission line at $9136$\\AA\\ and no detected continuum flux, which however is most likely an \\oiii\\ source at $z=0.824$ based on a weak detection of the \\oiiib\\ line. ",
        "introduction": "Observational study of the redshift range $6 \\la z \\la 7$ is crucial for understanding the reionization of intergalactic hydrogen, which was the most recent major phase transition for most of the baryonic matter in the universe.  While polarization of the microwave background measured by the Wilkinson Microwave Anisotropy Probe (WMAP) satellite indicates that substantial ionization had begun as early as $z \\sim 15$ (Spergel et al. 2003), both the opaque Gunn-Peterson troughs observed in the spectra of $z \\ga 6.3$ quasars (Becker et al. 2001; Fan et al. 2002) and the high temperature of the intergalactic medium at $z \\sim 4$ (Theuns et al. 2002; Hui \\& Haiman 2003) imply that a large part of the reionization was relatively recent, with substantial neutral gas lasting up to $z \\sim 6$.  \\lya\\ emission has proven the most effective tool so far for identifying star forming galaxies at redshifts $z>6$; indeed, all but two of the $\\sim 6$ galaxies that have been spectroscopically confirmed at $z>6$ were found through their \\lya\\ line emission in either narrowband or spectroscopic searches (Hu et al 2002 [H02]; Kodaira et al 2003 [K03]; Cuby et al 2003; this work; Kneib et al 2004), and at least one of the continuum-selected sources is also a \\lya\\ emitter (Cuby et al 2003).  We present the extension of the Large Area Lyman Alpha (LALA) survey to the $z\\approx 6.5$ window.  We have obtained spectra for each of the three $z\\approx 6.5$ candidates that pass all photometric selection criteria.  Two of these show strong emission lines in the narrowband filter bandpass. In one case, the line is identified as \\lya\\ at $z\\approx 6.535$ based on its asymmetric profile and the absence of other detected lines down to faint flux levels.  The second object is identified as an \\oiii\\ source at $z\\approx 0.824$ based on a probable ($4\\sigma$) detection of the \\oiiib\\ line. ",
        "conclusions": "We have performed a narrowband search for \\lya\\ emitting galaxies at $z\\approx 6.5$ in the Bo\\\"{o}tes field of the Large Area Lyman Alpha survey, and have obtained spectra for all of our viable candidates.  One source, \\gemwin, is confirmed as \\lya\\ at $z = 6.535$, based on an isolated, asymmetric emission line. The \\lya\\ luminosity of this source is large, $\\approx 1.1\\times 10^{43} \\ergsec$.  Objects this bright in the line are rare, occurring at a rate $\\sim 1$ per $2\\times 10^5 \\hbox{comoving} \\Mpc^3$ (this work and K03). It is therefore likely to correspond to a high peak in the density distribution at redshift $z=6.535$.  The equivalent width is also larger than most other $z\\approx 6.5$ galaxies, with a $2\\sigma$ lower bound of $>40$\\AA\\ (rest frame, corrected for IGM absorption).  This is difficult to reconcile with a neutral intergalactic medium unless the \\lya\\ line is  intrinsically strong {\\it and\\/} is emitted from its host galaxy with an intrinsic Doppler shift of several hundred $\\kms$."
    },
    "0403/astro-ph0403027_arXiv.txt": {
        "abstract": "{ It is shown that if the self--gravitating shearing sheet, a model of a patch of a galactic disk, is embedded in a live dark halo, this has a strong effect on the dynamics of density waves in the sheet. I describe how the density waves and the halo interact via halo particles either on orbits in resonance with the wave or on non-resonant orbits. Contrary to expectation the presence of the halo leads to a very considerable enhancement of the amplitudes of the density waves in the shearing sheet. This effect appears to be the equivalent of the recently reported enhanced growth of bars in numerically simulated stellar disks embedded in live dark halos. Finally I discuss the transfer of linear momentum from a density wave in the sheet to the halo and show that it is mediated only by halo particles on resonant orbits. ",
        "introduction": "The shearing sheet (Goldreich \\& Lynden--Bell 1965, Julian \\& Toomre 1966) model has been developed as a tool to study the dynamics of galactic disks and is particularly well suited to describe theoretically the dynamical mechanisms responsible for the formation of spiral arms. For the sake of simplicity the model describes only the dynamics of a patch of a galactic disk. It is assumed to be infinitesimally thin and its radial size is assumed to be much smaller than the disk. Polar coordinates can be therefore rectified to pseudo-Cartesian coordinates and the velocity field of the differential rotation of the disk can be approximated by a linear shear flow. These simplifications allow an analytical treatment of the problem, which helps to clarify the underlying physical processes operating in the disk.  In the present paper of this series I discuss the dynamical effects if the shearing sheet is immersed in a live dark halo. Dark halos are usually thought to stabilize galactic disks against non-axisymmetric instabilities. This was first proposed by Ostriker \\& Peebles (1973) on the basis of -- low--resolution -- numerical simulations. Their physical argument was that the presence of a dark halo reduces the destabilizing self--gravity of the disks. Doubts about an entirely passive role of dark halos were raised by Toomre (1977), but he (Toomre 1981) also pointed out that a dense core of a dark halo may cut the feed--back loop of the corotation amplifier of bars or spiral density waves and suppress thus their growth. Recent high-resolution numerical simulations by Debattista \\& Sellwood (2000) and Athanassoula (2002, 2003) have shown that quite the reverse, a {\\em destabilization} of disks immersed in dark halos, might be actually true. Athanassoula (2002) demonstrated clearly that much stronger bars grow in the simulations if the disk is embedded in a live dark halo instead of a static halo potential. This is attributed to angular momentum transfer from the bar to the halo via halo particles on resonant orbits. Angular momentum exchange between disk and halo has been addressed since the pioneering work of Weinberg (1985) in many studies theoretically or by numerical simulations and I refer to Athanassoula (2003) for an overview of the literature. Toomre (1981) has shown how the bar instability can be understood as an interference of spiral density waves in a resonance cavity between the corotation amplifier and an inner reflector (cf.~also Fuchs 2004). Thus it is to be expected that a live dark halo will be also responsive to spiral density waves and develop wakes, which I investigate here using the shearing sheet model. ",
        "conclusions": "If a self--gravitating shearing sheet is embedded in a live dark halo, the halo particles respond unexpectedly strong to density waves in the sheet. The interaction between the density waves and the halo particles is mediated both by halo particles on orbits in resonance with the waves and on non--resonant orbits. If the embedded shearing sheet is initially perturbed by a sinusoidal wave, a swing amplified density wave develops in the disk, which is of the same type as in an isolated sheet or a sheet in a static halo potential, but with an amplitude enhanced by a surprisingly large amount. This appears to be the equivalent of the enhanced bar growth in stellar disks embedded in live dark halos instead of static halo potentials seen in numerical simulations. There is transfer of linear momentum of density waves in the shearing sheet to the halo particles. This is mediated, however, entirely by halo particles on orbits in resonance with the waves similar to the torque exerted by bars on the surrounding halo."
    },
    "0403/astro-ph0403211_arXiv.txt": {
        "abstract": "We present a multi-wavelength catalog (15 $\\micron$, R, K-band, 1.4 GHz flux) plus spectroscopic identifications for 406 15 $\\micron$ sources detected in the European Large Area {\\it ISO} Survey (ELAIS) region S1, over the flux density range 0.5$<$$S_{15\\micron}$$<$150 mJy. 332 ($\\sim$82\\%) sources are optically identified down to R$\\sim$23.0. Spectra or bona fide stellar identifications are obtained for 290 objects ($\\sim$88\\% of the optically identified sources). The areal coverage, mid-infrared (MIR) and optical completeness of the sample are discussed in order to allow statistical and evolutionary analyses.  Two main spectroscopic classes have been found to dominate the MIR extragalactic population: $z<0.5$ star-forming galaxies (from absorbed to extreme starbursts: $\\nu L_{\\nu}(15 \\micron)$$\\approx 10^{8}-10^{11}$~L$_{\\odot}$), which account for $\\sim$75\\% of the sources, and Active Galactic Nuclei (AGN; both type 1 and 2), which account for $\\sim$25\\% of the sources. About 20\\% of the extragalactic sources are dust-enshrouded starburst galaxies [e(a) spectra], and all the starburst galaxies appear more dust extincted in the optical than nearby normal galaxies. We also identified 91 stellar objects ($\\sim$22\\% of the MIR sources). The counts for starburst galaxies and AGN down to 0.6 mJy have been derived. A general trend is found in the optical-MIR spectral energy distribution (SED) of the galaxies, where the MIR-luminous objects have larger MIR to optical luminosity ratios. Based on a variety of analyses, we suggest that the ELAIS sources fainter than R$\\sim$23 are luminous and ultra-luminous MIR galaxies (LIG-ULIGs; $\\nu L_{\\nu}(15 \\micron)$=10$^{11}$-10$^{12}$ L$_\\odot$) at intermediate redshifts ($z$=0.5--1.5), and that consequently the present sample is virtually 100\\% spectroscopically complete up to $z$=0.5. ",
        "introduction": "The ISOCAM instrument (Cesarsky et al. 1996) on board of the {\\it Infrared Space Observatory} ({\\it ISO}; Kessler et al. 1996) has been able to unveil at 15 $\\micron$ a population of MIR galaxies with two orders of magnitude fainter flux densities than observed by the {\\it Infrared Astronomical Satellite} ({\\it IRAS}).  These galaxies result to be ten times more numerous than expected if there were no evolution from $z = 0$ up to $z= 1 - 1.5$ (see Elbaz et al. 1999). These results are confirmed by the detection of a substantial diffuse cosmic Infrared Background in the 140 $\\mu$m - 1 mm range (Puget et al. 1996; Hauser et al.  1998; Fixsen et al. 1998; Lagache et al. 1999), which suggests a strong evolution of the galaxies in the IR (stronger than observed at any other wavelengths). The suggested evolution can either imply a larger fraction of galaxies experiencing an IR luminous phase in the past, or that MIR galaxies were more luminous in the past (see e.g. Chary \\& Elbaz 2001).  So far, the nature of the sources responsible for the strong evolution observed in the MIR band has only been studied in small fields, containing sources well known at other wavelengths: the Hubble Deep Field North (HDF-N; Aussel et al. 1999a,b) and South (HDF-S; Oliver et al. 2002; Mann et al. 2002; Franceschini et al. 2003), the Canada-France Redshift Survey (CFRS) 1452+52 field (Flores et al. 1999) and the ELAIS-S2 field (Pozzi et al. 2003). This has yielded a small but meaningful sample of sources sufficient for most multiwavelength studies. Most of these sources can easily be identified with relatively bright optical counterparts ($I < 22.5$) with a median redshift of $z\\simeq0.7-0.8$ imposed by the MIR {\\it k}-correction (Aussel et al. 1999a,b; Flores et al. 1999).  However, to obtain reliable source counts and statistically significant information about the nature and evolution of the MIR extragalactic source population, large-area surveys covering a wide range of flux densities are required. In particular, it is very important to bridge the gap between the {\\it IRAS} surveys ($S_{15\\micron}>200$ mJy) and the Deep ISOCAM surveys ($0.1< S_{15\\micron}<2-3$ mJy). ELAIS (Oliver et~al. 2000), with its broad flux range (0.5$<S_{15\\micron}<$150 mJy) and several thousand sources detected over a $\\sim$12 deg$^2$ area (Rowan-Robinson et~al. 2004), is the crucial survey to investigate the nature and evolution of the MIR sources in the region where the source counts start diverging from the non-evolution predictions (a few mJy at 15 $\\mu$m; Elbaz et al. 1999; Gruppioni et al. 2002) and with data not affected by cosmic variance (which might be a strong effect in the narrow area surveys).  In order to better understand the nature of the MIR extragalactic population on the basis of a larger identification statistics than the ones obtained so far, we performed photometric and spectroscopic identifications for a source sample selected at 15 $\\mu$m in the ELAIS southern field S1 ($\\sim$4 deg$^2$). Here we present the results of the identification of 406 objects, which form a complete and highly reliable sub-sample drawn from the 15 $\\micron$ catalog of Lari et~al. (2001). The areal coverage, MIR and optical completeness are discussed in detail to allow statistical and evolutionary analyses with this sample. The counts down to 0.6~mJy for each class of source have been derived. The multiwavelength (radio, MIR, near-infrared (NIR), optical) properties of the 15 $\\mu$m sources are investigated in order to study the nature and evolution of the infrared extragalactic population. The infrared luminosity function and the evolution of 15 $\\mu$m galaxies and AGN are discussed in two companion papers (Pozzi et al. 2004; Matute et al. 2004, in preparation). A preliminary analysis of the AGN evolution has been presented in Matute et~al. (2002) \\footnote{A collection of papers and scientific results related to the ELAIS southern fields is available at \\url{http://www.fis.uniroma3.it/$\\sim$ELAIS\\_S}. }    This paper is structured as follows: in section 2 we describe the 15 $\\micron$ sample; in section 3 we present the optical CCD observations, the source identification and the cross-correlation with catalogs in other bands; in section 4 we describe the spectroscopic observations, the classification and the spectroscopic catalog; in section 5 we present the counts; in section 6 we discuss the nature of the MIR galaxies.  Throughout this paper we will assume $H_0 = 75$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m = 0.3$ and $\\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": "We have presented multiband (R, K, 1.4 GHz) and optical spectroscopic identifications of a reliable subsample of 406 15-$\\micron$ sources from the ELAIS-S1 survey over the flux density range 0.5$<S_{15\\micron}<$150 mJy. These observations substantially contribute to the understanding of the evolution of star-forming galaxies and AGN in the MIR, which is presented in two companion papers by Pozzi et al. (2004) and Matute et al. (in preparation) respectively. This paper has described the details of the ground-based follow-up observations and data reduction, and has presented optical identifications for $\\sim$82$\\%$ of the 15 $\\micron$ sources down to R$\\sim$23, and spectroscopic identification for $\\sim$88$\\%$ of the optically identified sample. The areal coverage and completeness of the survey has been described in detail, in order to allow statistical analyses. The source counts for star-forming galaxies and AGN have been derived: the ELAIS extragalactic sources down to S$_{15\\micron}$=0.6 mJy have a density of 370$\\pm$14/deg$^2$, and produce a 15 $\\micron$ background of $\\nu$I($\\nu$) = 0.34 nW m$^{-2}$sr$^{-1}$. This corresponds to $\\sim$12$\\%$ of the total extragalactic background estimated by Metcalfe et al. (2003).  Taking into account the optical based classification, we estimate an AGN contribution of at least $\\sim$6\\% to the MIR sources brighter than S$_{15\\micron}$=0.6 mJy. A larger fraction is expected if the classification will be based on the X-ray luminosity.  The line ratios of the optical spectra of the MIR galaxies suggest that they are affected by substantial reddening, whith the more MIR-luminous sources having larger dust extinction. A general trend is found in the MIR-optical SEDs of {\\it IRAS} and {\\it ISO} galaxies whith the more MIR-luminous sources having larger $L_{15}/L_{R}$ ratios. According to our analyses we conclude that most of the ELAIS galaxies fainter than R$\\sim$23 should have redshift greater than 0.5 and are likely to be LIG-ULIG galaxies. As a consequence, the present sample of star-forming galaxies has to be considered virtually complete up to $z=0.5$."
    },
    "0403/astro-ph0403094_arXiv.txt": {
        "abstract": "We present numerical results on the properties of young binary and multiple stellar systems. Our analysis is based on a series of SPH $+$ $N$-body simulations of the fragmentation of small molecular clouds, that fully resolve the opacity limit for fragmentation. These simulations demonstrate that multiple star formation is a major channel for star formation in turbulent flows. We have produced a statistically significant number of stable multiple systems, with components separations in the range $\\sim 1-10^3$ AU. At the end of the hydrodynamic stage (0.5 Myr) we find that $\\approx 60\\%$ of stars and brown dwarfs are members of multiples systems, with about a third of these being low mass, weakly bound outliers in wide eccentric orbits. Our results imply that in the stellar regime most stars are in multiples ($\\approx 80\\%$) and that this fraction is an increasing function of primary mass. After $N$-body integration to 10.5 Myr, the percentage of bound objects has dropped to about 40\\%, this decrease arising mostly from very low mass stars and brown dwarfs that have been released into the field. Brown dwarfs are never found to be very close companions to stars (the {\\it brown dwarf desert at very small separations}), but one case exists of a brown dwarf companion at intermediate separations (10 AU). Our simulations can accommodate the existence of brown dwarf companions at large separations, but only if the primaries of these systems are themselves multiples.  We have compared the outcome of our simulations with the properties of real stellar systems as deduced from the IR colour-magnitude diagram of the Praesepe cluster and from spectroscopic and high-resolution imaging surveys of young clusters and the field. We find that the spread of the observed main sequence of Praesepe in the $0.4-1$M$_\\odot$ range appears to require that stars are indeed commonly assembled into high-order multiple systems. Similarly, observational results from Taurus and $\\rho$ Ophiuchus, or moving groups such a TW~Hydrae and MBM~12, suggest that companion frequencies in young systems can be indeed as high as we predict. The comparison with observational data also illustrates two problems with the simulation results. Firstly, low mass ratio ($q < 0.2$) binaries are not produced by our models, in conflict with both the Praesepe colour magnitude diagram and independent evidence from field binary surveys. Secondly, very low mass stars and brown dwarf binaries appear to be considerably under-produced by our simulations. ",
        "introduction": "The formation of binary and multiple systems provides one of the most exacting areas in which star formation theory can be compared with observational data. Stars are known to have a binary frequency in excess of 50\\%, both in the field (Duquennoy \\& Mayor 1991, henceforth DM91; Fisher \\& Marcy 1992; Halbwachs et al. 2003) and in clusters (K\\\"ohler \\& Leinert 1998; Patience et al. 1998; Mathieu et al. 2000; Bouvier et al. 2001).  For pre-main-sequence stars this frequency seems to be even higher (Reipurth \\& Zinnecker 1993; Simon et al. 1995; Duch\\^ene 1999; Reipurth 2000). Thus, understanding the formation of multiple stars becomes necessary if we are to understand star formation in general. Recently, Bate, Bonnell \\& Bromm (2003; henceforth BBB) have shown that multiple stars are a natural byproduct of the collapse and fragmentation of turbulent molecular clouds. In particular, close binary stars can be indirectly formed out of the fragmentation of a molecular cloud core by means of dynamical interactions with other stellar objects in the vicinity, as well as interactions with circumbinary discs (Bate, Bonnell \\& Bromm 2002b). However, this sort of calculation (that fully resolve fragmentation in a cloud of e.g. 50 M$_\\odot$ of gas in BBB's case) is very demanding computationally, and a wide range of initial conditions cannot easily be explored. In addition, such calculations remain inconclusive about the properties of multiple stars (in particular the existence of long lived multiples and the production of wide binaries) since the evolution of the resulting stellar systems is not followed until decay to a stable hierarchical configuration is achieved. For example, out of the 50 objects formed in the BBB simulation, 18 are contained in non-hierarchical groupings containing 11 and 7 members. Since, in this simulation, only the more compact groupings have had time to decay to a stable configuration, BBB understandably focused their attention on the implications of their work for {\\it close} binary formation.  In this paper we have taken an alternative approach which is a natural complement to the BBB simulations. Through the modelling of less massive clouds over longer timescales, we have been able to follow wide binary formation and the production of long lived multiple systems.  We find that by modelling an ensemble of isolated cores of mass 5~M$_\\odot$, we can improve the number of stars formed per CPU hour by a factor 7 compared with the BBB calculation, which comprised a total mass of 50~M$_\\odot$.  This economy stems from the fact that by focusing on individual dense cores, we dispense with the computational expense of following the diffuse gas in the BBB simulation. We remark, however, as a caveat, that we find that the formation of multiples proceeds {\\it hierarchically}, with structures on any scale being progressively modified by interactions on larger scales. Thus in these models with scale free turbulence, {\\it any} upper mass cut-off (i.e. the finite mass of the system modelled) may have some influence on the multiple systems produced. By comparing our results with those of BBB, we are able to assess how sensitive our results are to the total mass of the region simulated.  The computational economies that we gain through the simulation of less massive cores, allows us to instead concentrate our efforts on the longer timescale integration of the multiple systems produced.  We wish to understand how our results are affected if we {\\it throw away} the gas component at a certain evolutionary stage (when 60\\% of the initial gas has been accreted) and thereafter evolve the system as an $N$-body ensemble. We follow the stellar dynamical $N$-body problem for 10~Myr (this typically corresponding to $10^4$ orbits of multiple systems at the median separation), using the {\\sc nbody1} code, by Aarseth (Aarseth 1963). Another advantage that we reap from this less computationally expensive approach is that we gain improved statistics (i.e. we form $\\approx 150$ stars and brown dwarfs, compared with the $\\approx 50$ of BBB), and from this dataset are able to extract reasonable statistics, in a manner similar to Sterzik \\& Durisen (1998; 2003) for the purely $N$-body case. We stress that our simulations share the property of the BBB simulations that they resolve the opacity limit for fragmentation (Low \\& Lynden-Bell 1976; Rees 1976) and that, assuming that fragmentation does not occur at densities greater than those at which the gas becomes opaque to infrared radiation, these calculations are able to model the formation of all the stars and brown dwarfs that, under the initial conditions imposed, can be produced. Our spatial resolution limit for binaries allows us to study a wide range of separations, and the particle numbers we employ allow us to model accretion discs around the protostars which are as long lived as those modelled by BBB.  Our study has four major findings: \\begin{itemize} \\item The incidence of multiple systems (with N $>$ 2) is {\\it high} (8 out of 18 N $>$ 1 systems at $t=10.5$ Myr) and we therefore test whether our results are compatible with the photometric width of the main sequence in young clusters. The companion frequency varies significantly during the pure $N$-body evolution of the systems. \\item Although we produce a variety of triples, quadruples and higher order multiples, they tend to follow a characteristic pattern of {\\it internal mass distribution}. This involves each binary pair having a mass ratio that does not deviate strongly from unity, and likewise, for quadruples, the mass ratio of the total mass in each binary is not greatly different from unity (i.e. mass ratios in the range $0.5-1$ are favoured in each case). The exception to this is that multiple systems are commonly orbited by {\\it low mass outliers} on eccentric orbits, which are the result of incomplete ejections of low mass objects from multiple systems. We discuss how deep images of the environs of multiple systems, and spectroscopic studies of the primaries of extreme mass ratio binaries, may be used to test this prediction. \\item The multiplicity fraction is an increasing function of primary mass. The brown dwarf desert at {\\it very small separations} is reproduced by our models. \\item We confirm the result of previous simulations that brown dwarf binaries indeed appear to be under-produced by turbulent fragmentation calculations. \\end{itemize}  The structure of this paper is as follows. In Section 2 the computational method and initial conditions applied to our models are described. The results on multiple stars are given in Section 3. In Section 4 we perform a detailed comparison of our results with available observational data, and suggest future experiments. Our conclusions are given in Section 5. ",
        "conclusions": "We have undertaken the first hydrodynamical $+$ $N$-body simulations of multiple star formation that have produced a statistically significant number of stable hierarchical systems, with component separations in the range $\\sim 1-1000$ AU. These simulations have demonstrated that multiple star formation is a major channel for star formation in turbulent flows. The hydrodynamical simulations are followed for $\\approx 0.5$ Myr; subsequently, the remaining gas is removed and the stellar systems followed as $N$-body ensembles for an additional 10 Myr. At this point, all but one of the surviving multiple systems are stable, according to the criteria of Eggleton \\& Kiseleva (1995).  We find that the properties of the resulting multiple systems are not significantly sensitive to the large scale geometry of the cloud -- determined by the turbulence -- but rather to the dynamical and competitive accretion processes taking place within the mini-clusters formed out of the collapse and fragmentation of the cloud.   At an age of $0.5$ Myr, we find that about $60 \\%$ of stars and brown dwarfs are in multiple systems, with about a third of these being low mass, weakly bound outliers. Excluding these outliers and unbound objects, $7 \\%$ of the remaining objects are in pure binaries ($2$ systems), $14 \\%$ are in quadruples ($2$ systems), $35 \\%$ are in quintuples ($4$ systems), $32 \\%$ are in sextuples ($3$ systems) and $12 \\%$ are in multiples with seven components ($1$ system). The companion frequency is therefore very high, $\\approx 1$. We find that our multiples consist of hierarchies of binaries and triples and that {\\it planetary multiples} (in which companions are not members of binary/triple systems other than the multiple itself) are comparatively rare (occurring $\\sim 25 \\%$ of the time). There is a distinctive pattern of mass distribution within these multiples, with the mass ratio within binaries, and the mass ratios between binaries, rarely deviating far from unity (values of $0.5-1$ are typical). On the other hand, such systems are typically orbited by several low mass outliers (typically at separations of $\\sim 10^4$ AU) on eccentric orbits. About 90\\% of these objects are unstable in timescales of a few $\\times 10^6$ yr (i.e. a few $\\times$ their typical orbital timescale).  We find that the $40 \\%$ of objects that are unbound are overwhelmingly of low mass (median mass $\\approx 0.02$ M$_\\odot$). Thus our results imply that in the stellar regime, most stars are in multiples ($\\approx 80\\%$) and that this multiplicity fraction $f_{\\rm m}$ is an increasing function of mass. In this latter respect, these results are qualitatively consistent with a large body of previous works on the decay of small-$N$ systems, both with and without gas (van Albada 1968; McDonald \\& Clarke 1993, 1995; Sterzik \\& Durisen 1998, 2003; DCB03). The high $f_{\\rm m}$ values for GK stars are consistent with adaptive optics measurements of nearby young associations such as MBM~12 and TW~Hydrae (e.g. Brandeker, Jayawardhana \\& Najita 2003), where multiplicity fractions as high 0.64 are found, and radial velocity surveys of visual binaries (e.g. Tokovinin \\& Smekhov 2002) which raise the percentage of spectroscopic sub-systems to at least 40\\%. Low-mass SFR such as Taurus or $\\rho$ Ophiuchus also show companion frequencies in the range $0.3-0.5$, comparable to those predicted by our models at later times. It must be pointed out that the values of the multiplicity fraction $f_{\\rm m}$ for each mass range do not change significantly during the $N$-body evolution of the systems.  At an age of 10.5 Myr the fraction of bound and unbound objects has reversed: 40\\% remain in multiples and 60\\% are singles. The companion frequency has dropped to $\\approx 0.3$ due to the ejection of bound outliers to the field. This transference of objects from bound to unbound orbits results in an increase of the number of free floating brown dwarfs by $\\approx 60 \\%$. In this 10 Myr time-span, many multiple systems also experience internal decay: excluding the remaining 3 outliers, 42\\% of the remaining bound objects are in pure binaries (11 systems), 12\\% are in triples (2 systems), 15\\% are in quadruples (2 systems), 19\\% are in quintuples (2 systems) and 12\\% are in sextuples (1 system).  We pointed out that low mass stars (and, especially, brown dwarfs) are {\\it locked up} in multiple systems at early times and subsequently released into the field. This remark needs some qualification however. Multiple star formation is hierarchical in our simulations, with structures forming on a particular scale being modified as a result of subsequent merging with structures on a larger scale.  It is notable that we find low mass outliers at a separation that is similar to the initial size of the core, and we speculate that if we had modelled a larger volume of cloud, rather than isolated cores, we might have found that these outliers would already have been stripped off by interactions with structures on a larger scale before some of them could settle into stable orbits. This suggests that such outliers should be sought in imaging of relatively isolated pre-main sequence groups and may explain why no such outliers have been detected through deep imaging of multiple systems in Taurus (G. Duch\\^ene, private communication).  If such outliers do indeed survive the formation process, then about 10\\% of them are in stable hierarchical orbits at 10.5 Myr.  We would thus expect some brown dwarfs to remain in the outer reaches of multiple systems even in the field.  Our simulations can thus accommodate the existence of systems with brown dwarfs as wide companions (Gizis et al. 2001) but {\\it only if the primaries of these systems are themselves multiple systems}. (We find that three out of the four bound brown dwarfs present at $t=10.5$ Myr are orbiting a multiple system).  We therefore predict that the primaries of binaries containing a brown dwarf in wide orbit should themselves be multiple systems.  We have examined how well the products of our simulations compare with the properties of real stellar systems as deduced from the colour magnitude diagram of young clusters (specifically the infrared colour magnitude diagram for Praesepe). Our simulation data compares very favourably with the width of the main sequence in the mass range 0.4-1 M$_\\odot$; indeed the spread of the main sequence in this mass range appears to {\\it require} that stars are commonly assembled into high order multiple systems, although the number of outliers from a pure equal mass binary sequence is not large.  The comparison with the Praesepe colour magnitude diagram however illustrates two problems with the simulation results.  Firstly, the simulation produces no single star sequence at masses greater than 0.35 M$_\\odot$ (colour bluer than $I-K=2$; there are only two blue circles above 0.35 M$_\\odot$, at 0.581 and 0.584 M$_\\odot$), whereas the observational data shows such a sequence, indicating that single stars and/or low mass ratio binaries are produced in this colour range. We could probably alleviate this problem by modelling a more massive cloud. This would increase the maximum mass of stars produced and enable some stars in the colour range considered to be ejected as singles by encounters with more massive stars.  Although this would bring closer agreement with the observed colour magnitude diagram, the lack of low mass ratio binaries in our simulations is in conflict with independent evidence from field binary surveys, such as that of DM91 for G dwarfs.  The DM91 data (containing binaries with median separation $30$ AU) showed a distribution with mass ratio ($q$) that rose with decreasing $q$ towards their completeness limit of $q\\sim 0.2$, whereas our mass ratio distribution (with the exception of the very low mass outliers) is strongly concentrated between $0.5$ and unity.  This inability to reproduce enough extreme mass ratio systems is a feature of all hydrodynamic modelling of multiple systems to date. It has not however been widely discussed previously, since BBB focused on the close systems, where the observed mass ratio distribution is in any case much more concentrated towards unit mass ratio than for the binary population as a whole (Mazeh et al. 1992; Halbwachs et al. 2003) and where the simulations very naturally reproduce the observed brown dwarf desert at {\\it very small} separations (Marcy \\& Butler 2000; Halbwachs et al. 2000).  The second problem revealed by the colour magnitude diagram relates to the fact that the simulations produce almost no binaries with primaries redder than $I-K=2.5$ whereas the data exhibits a pronounced scatter (consistent with large numbers of binaries) in this mass range ($0.15-0.4 M_\\odot$).  This scarcity of binaries with low mass primaries in our simulations also conflicts somewhat with the results of binary surveys among M dwarfs (Fischer \\& Marcy 1992; see Figure 4, this paper) and brown dwarfs (Close et al. 2003; Gizis et al. 2003; Mart\\'{\\i}n et al. 2003; Bouy et al. 2003).  In summary, we have found that our simulations produce large numbers of hierarchical multiple systems and that relatively isolated young multiples may harbour weakly bound brown dwarf outliers, as a relic of the hierarchical formation process in turbulent flows. We predict that where brown dwarfs are found in wide orbits, the primary should itself turn out to be a multiple.  The simulations are consistent with a number of observational constraints: the high (but presently poorly constrained) incidence of hierarchical multiples among field stars and pre-main sequence stars, the absence of brown dwarfs as close companions to normal stars (the brown dwarf desert) and, at a qualitative level at any rate, the positive dependence of the binary fraction on primary mass. There are two areas in which the simulations are not able to replicate observations properly: the simulations under-produce the binary fraction at low masses (M dwarfs and brown dwarfs) and also do not generate enough wide stellar pairs with low mass ratios. ?"
    },
    "0403/astro-ph0403577_arXiv.txt": {
        "abstract": "{We report on the first simultaneous near-infrared/X-ray detection of the SgrA* counterpart which is associated with the massive 3--4$\\times$10$^6$\\solm ~black hole at the center of the Milky Way. The observations have been carried out using the NACO adaptive optics (AO) instrument at the European Southern Observatory's Very Large Telescope\\footnote{Based on observations at the Very Large Telescope (VLT) of the European Southern Observatory (ESO) on Paranal in Chile; Program: 271.B-5019(A).} and the ACIS-I instrument aboard the \\emph{Chandra X-ray Observatory}. We also report on quasi-simultaneous observations at a wavelength of 3.4~mm using the Berkeley-Illinois-Maryland Association (BIMA) array. A flare was detected in the X-domain with an excess 2 - 8 keV luminosity of about 6$\\times$10$^{33}$~erg/s.  A fading flare of Sgr~A* with $>$2 times the interim-quiescent flux was also detected at the beginning of the NIR observations, that overlapped with the fading part of the X-ray flare. Compared to 8-9 hours before the NIR/X-ray flare we detected a marginally significant increase in the millimeter flux density of Sgr A* during measurements about 7-9 hours afterwards. We find that the flaring state can be conveniently explained with a synchrotron self-Compton model involving up-scattered sub-millimeter photons from a compact source component, possibly with modest bulk relativistic motion. The size of that component is assumed to be of the order of a few times the Schwarzschild radius.  The overall spectral indices $\\alpha_{NIR/X-ray}$ ($S_{\\nu}$$\\propto$$\\nu^{-\\alpha}$) of both states are quite comparable with a value of $\\sim$1.3.  Since the interim-quiescent X-ray emission is spatially extended, the spectral index for the interim-quiescent state is probably only a lower limit for the compact source Sgr~A*.  A conservative estimate of the upper limit of the time lag between the ends of the NIR and X-ray flare is of the order of 15~minutes. ",
        "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 dynamics of stars (Eckart \\& Genzel 1996, Genzel et al. 1997, 2000, Ghez et al. 1998, 2000, 2003a, 2003b, Eckart et al. 2002, Sch\\\"odel et al. 2002, 2003).  Located at a distance of only 8 kpc from the solar system (Reid 1993, Eisenhauer et al. 2003), it allows detailed observations of stars at distances much less than 1~pc from the central black hole candidate, the compact radio source Sgr~A*.  Additional strong evidence for a massive black hole at the position of Sgr~A* came from the observation of interim-quiescent and flare activity from that position both in the X-ray and recently in the near-infrared wavelength domain (Baganoff et al. 2001, 2003, Eckart et al. 2003, Porquet et al. 2003, Goldwurm et al. 2003, Genzel et al. 2003, Ghez et al. 2004).  Throughout the paper we will use the term 'interim-quiescent' (or IQ) for the apparently constant, low-level flux density states at any given observational epoch since current data cannot exclude flux density variations of that state on longer time scales (days to years). This is especially true for the more compact NIR source (Genzel et al. 2003, Ghez 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 observations of SgrA* covering an X-ray flare simultaneously in the near-infrared using seeing limited exposures revealed only upper limits to the NIR flux density (Eckart et al. 2003). In section 2 of the present paper we report on the first successful simultaneous NIR/X-ray observations using adaptive optics.  These observations were for the first time successful in detecting radiation from the SgrA* counterpart both in the NIR and the X-ray wavelength domain.  A detailed statistical analysis that supports the simultaneous detection of a SgrA* flare event in the NIR and X-ray domain is given in section 3 of this paper. In section 2 we also describe the quasi-simultaneous mm-observations that were taken just before and after the NIR/X-ray observations. In section 4 we briefly discuss the flux densities and spectral indices we derived from the available data.  Section 5 gives a first physical interpretation of the simultaneous detection of SgrA*. A short summary and discussion of the results is given in section 6. ",
        "conclusions": "We have presented the first successful simultaneous X-ray and NIR detection of SgrA* in a flaring and the IQ-state.  The X-ray flare lasted for about 42 minutes and began on 19 June 2003 at about 23:10 UT, more than 4 hours into the X-ray observation.  The peak of the flare occurred less than 10 minutes prior to the time of the first VLT/NACO DDT image. The emission during the flare can successfully be described by a SSC model in which the NIR and X-ray flux density excess is produced by up-scattering sub-mm-wavelength photons into the NIR and X-ray domain.   With respect to its FWZP duration of 55 - 115 minutes the flare reported here compares well with other flares measured so far. Baganoff et al. (2001), Eckart et al. (2003), and Porquet et al. (2003) report on X-ray events of 45 to 170 minutes.  Ghez et al. (2004) and Genzel et al. (2003) report on NIR flare events that last 50 to 80 minutes, respectively.  Our newly detected flare event is weaker than most others which have been reported and not necessarily representative of the characteristics of the stronger flares (factor of 50: Baganoff et al. 2001; factor of $>$100: Porquet et al. 2003; Goldwurm et al. 2003).  However, during our 2002 \\emph{Chandra} monitoring session we found that flares that are a factor of $>$10 stronger than the quiescent emission occur at a rate of $0.53\\pm0.27$ per day.  Weaker flares are more frequent and our newly detected flare event is probably more representative for those weaker flares.   Although we have given preference to a simple SSC model in explaining the observed simultaneous flare emission, flux density contributions via other emission mechanisms may be of relevance.  Along with enhanced electron heating leading to SSC flares, Markoff et al. (2001) also suggested the possibility that the flares may result from acceleration.  This could result from electrons which are energetic enough to account for both the NIR and X-ray flares via direct synchrotron radiation.  In fact even SSC models presented by Markoff et al. (2001) and Yuan, Quataert, \\& Narayan (2003) result in a significant amount of direct synchrotron emission in the infrared (see also synchrotron models in Yuan, Quataert, \\& Narayan 2004).  The different possible emission mechanisms may also be coupled in a more complicated way. Ambient thermal electrons may be heated during a flare and may produce excess sub-millimeter and infrared flux density. This process could lead to correlated radio/NIR/X-ray variability quite similar to to what is expected in SSC models (Yuan, Quataert, \\& Narayan 2003, 2004, Genzel et al. 2003, Ghez et al. 2004).  These synchrotron/ inverse Compton models may also result in small (or no) time lag between the NIR and X-ray emission - compatible with our limit on the time lag for the decaying flank of the flare.  In such a scenario it will be difficult to determine the relative importance and flux density contributions of the different emission mechanisms by variability.   Since the X-ray source responsible for most of the X-ray IQ-state flux density of SgrA* is extended over a radius of about 0.6~arcsec (Baganoff et al. 2003) its emission can most likely be ascribed to bremsstrahlung from a thermal particle distribution.  While the interim-quiescent NIR flux density probably can be attributed to a fairly compact component (Genzel et al. 2003), it is currently not clear how much of the quiescent X-ray flux density originates from that compact region.  The extended X-ray component probably originates from hot gas within the $\\sim$1\" Bondi radius that is associated with the accretion flow (see Baganoff et al. 2003 and Quataert 2003 for detailed discussions).   If thermal bremsstrahlung is an important mechanism to explain the simultaneous NIR and the X-ray emission, then geometries are conceivable in which the cooler $\\ga$10$^3$~K plasma that gives rise to the NIR emission is spatially offset from the hotter $\\ga$10$^8$~K plasma. One may also speculate about the possibility that the amount of cooler plasma might be larger than the amount of hot plasma. These facts could result in a situation in which the X-ray variations lead the NIR-variations. On the other hand the cooling timescales of the hotter plasma may be larger which would result in the opposite behavior.  There are of course also models conceivable that involve small source sizes or apparently cospatial distributions that would result in no measurable time lag or even NIR and X-ray flare events that need not always be correlated (see discussion by Yuan, Quataert, \\& Narayan 2004).   If we assume that the (marginally) higher mm-flux density after the flare is related to the flare activity we observed in the X-ray and NIR-domain, then there seems to be a correlation between flare activity in the mm/sub-mm and the X-ray domain in three cases now: the original X-ray flare reported by Baganoff et al. (2001; see also Zhao et al. 2004), the activity around the largest X-ray flare reported by Porquet et al. (2003) and Zhao et al. (2004), and now the small X-ray flare we report on in this publication.  However, during an 8-day simultaneous observing campaign using the Owens Valley mm-array at a wavelength of 3-mm and Chandra in the X-ray domain (Mauerhan et al. 2004) did not detect an obvious correspondence of significant fluctuations in both wavelength domains. The currently available data also indicate that the mm-activity may be a function of the X-ray flare magnitude. In the event reported by Porquet et al. (2003) and Zhao et al. (2004) the mm-flux density flared by 100\\% of the continuum following the factor $160$ X-ray flare. Here we report a factor 2 to 3 X-ray flare for which we observed a no more than 10\\% mm-flare following the NIR/X-ray event.   Future observations will reveal possible relations between different wavelength regimes. An important question is whether individual mm-flare events are related to events in the NIR or X-ray regime,  or whether there are in general more or stronger flares per unit time, when the average mm- or sub-mm flux density is higher. 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-state and flare phenomena associated with SgrA* at the position of the massive black hole at the center of the Milky Way."
    },
    "0403/astro-ph0403354.txt": {
        "abstract": "For a detailed comparison of the appearance of  cluster of galaxies in X-rays and  in the optical,  we have compiled a comprehensive database of X-ray  and optical properties of a  sample of clusters based on the largest available X-ray and optical  surveys: the ROSAT All Sky Survey (RASS) and  the Sloan Digital  Sky Survey  (SDSS).  The  X-ray  galaxy clusters of this RASS-SDSS catalog cover a  wide range of masses, from groups of $10^{12.5}$ $M_{\\odot}$  to  massive clusters of  $10^{15}$ $M_{\\odot}$ in the redshift range  from 0.002 to 0.45.  The RASS-SDSS sample comprises  all the X-ray selected   objects already observed by the Sloan Digital Sky Survey (114 clusters).  For each system we have uniformly determined the X-ray    (luminosity  in the  ROSAT    band, bolometric luminosity,   center  coordinates)  and  optical properties (Schechter luminosity function parameters, luminosity, central galaxy density, core, total  and half-light radii).  For  a  subsample of 53 clusters we have also compiled the  temperature and the iron abundance from the literature.  The  total optical luminosity can be  determined with a  typical uncertainty of 20\\%  with a  result independent of the choice  of local  or  global background subtraction.   We searched for parameters which  provide   the  best correlation between   the  X-ray luminosity and  the optical properties   and found  that  the z   band luminosity  determined within  a   cluster aperture  of    0.5 Mpc $\\rm{h}_{70}^{-1}$   provides the best  correlation with a scatter of about 60-70\\%.  The scatter  decreases to less  than 40\\%  if the correlation  is limited to the  bright X-ray clusters.  The resulting correlation of $L_X$  and $L_{op}$   in  the z and   i  bands shows  a logarithmic slope of 0.38, a value  not consistent with the assumption of a constant $M/L$.  Consistency is found, however, for an increasing M/L  with  luminosity as  suggested   by other observations.  We  also investigated     the  correlation between  $L_{op}$     and the  X-ray temperature obtaining the same result. ",
        "introduction": "Cluster  of galaxies are the largest  well  defined building blocks of our Universe. They form via  gravitational collapse of cosmic matter over a region of several megaparsecs. Cosmic baryons, which represent approximately  10-15\\%  of the  mass content of  the Universe, follow dynamically the dominant dark matter during the collapse. As a result of  adiabatic  compression and   of  shocks generated  by  supersonic motions,  a    thin hot  gas   permeates   the  cluster gravitational potential.  For a typical cluster  mass  of $10^{14}$ $M_{\\odot}$ the intracluster gas reaches a temperature of the order of $10^7$ keV and, thus, radiates   optically  thin    thermal bremsstrahlung and    line radiation in  the X-ray band.  In 1978, the launch  of the first X-ray imaging telescope, the \\emph{Einstein} observatory, began a new era of cluster discovery, as clusters proved to be luminous ($\\ge 10^{42-45}$ ergs $\\rm{s}^{-1}$),  extended  ($\\rm{r}\\sim 1-5$ Mpc) X-ray  sources, readily identified in the  X-ray sky. Therefore, X-ray observations of galaxy clusters  provide an efficient  and physically motivated method of identification  of  these structures.  The X-ray  selection is more robust against contamination along the  line of sight than traditional optical methods since  the richest  clusters  are relatively  rare and since    X-ray  emission, which is    proportional  to the gas density squared, is  far more sensitive to physical  overdensities than in the projected  number   density of galaxies   in   the sky.  In  fact  the existence  of diffuse,   very hot  X-ray  emitting  gas   implies  the existence  of  a   massive confining   dark  matter  halo.  Moreover, selection  according to X-ray luminosity is  also an efficient way to find the highest mass  concentrations due to well defined correlation between the X-ray luminosity and  the total cluster mass (Reiprich \\& Bh\\\"oringher.  2002). In addition  to allowing the identification  of galaxy clusters, X-ray observations provide  a wealth of  information on the intracluster medium  itself, e.g. its metal abundance,  radial density    distribution and   temperature   profile. These   latter quantities, in turn, can be used to reliably estimate the \\emph{total mass} of the system.    In addition  to  the   hot,  diffuse  component,   baryons   are also concentrated in the individual galaxies within the cluster. These are best studied through photometric  and spectroscopic  optical surveys, which provide  essential   information about luminosity,   morphology, stellar population and  age. Solid observational evidences indicate a strong interaction between  the two baryonic  components, as galaxies pollute the intracluster medium  expelling metals via  galactic winds producing the  observed  metal abundances  in clusters  (De Grandi  et al. 2002, Finoguenov et al. 2001). On the other hand, the evolution of galaxies in clusters is  influenced by processes  due  to the hot  gas (e.g. gas stripping  by  ram pressure, etc.)   as  it is by  internal processes like star formation,  galactic winds, supernovae  explosions etc., operating  inside galaxies themselves   (Dressler et al.   1997, Fasano et al. 2000, van Dokkum et al.  2000, Lubin et al. 2002, Kelson et  al.  1997,2000, Ziegler\\&Bender   1997,  Gomez et  al.  2003).  In conclusion, understanding the complex  physics at play in clusters of galaxies   requires combined   X-ray  and  optical   observation  of a statistically significant sample of these objects.   On  the    basis of  these considerations,  we    have created a large database of clusters of galaxies based on the largest available X-ray and optical surveys: the ROSAT All Sky Survey  (RASS), the only X-ray all sky survey ever  realized using an  imaging X-ray  telescope, and the Sloan  Digital Sky Survey  (SDSS), which  is observing the  whole Northern Galactic  Cap and part of  the Southern Galactic Cap in five wide  optical bands covering the entire optical range.  By carefully combining  the data of the two  surveys we have created the RASS-SDSS galaxy cluster catalog.  Although two galaxy cluster catalogs from the SDSS already exist, the Cut and Enhance Galaxy Cluster Catalog of Goto et al.  (2002)  and  the  Merged Cluster   Catalog of  Bahcall et  al. (2003, see also Kim et al. 2002), we prefered to compile a new cluster catalog by selecting the systems in the  X-ray band, for which we have reliable X-ray  characteristics and  for the reasons  explained above. The X-ray-selected galaxy clusters cover a wide range of masses, from groups of $10^{12.5}$ $M_{\\odot}$  to  massive clusters of  $10^{15}$ $M_{\\odot}$ in a redshift range  from 0.002  to 0.45.  The RASS-SDSS sample comprises  all the X-ray  detected objects already  observed in the sky region covered by the Sloan Digital Sky Survey.  One of the first goals is the comparison  of the X-ray and the optical appearance  of   the clusters.   We want   in  particular find optical parameters that  provide the    closest   correlation to the     X-ray parameters, such that we can predict  within narrow uncertainty limits the X-ray luminosity from these  optical parameters and vice versa.  So far  optical     and   X-ray cluster surveys    have    been conducted independently without much  intercomparison.  Therefore, the empirical relationship between the   X-ray luminosity and optical  luminosity of clusters is   not so  well defined,  in   large part  because  of  the difficulties inherent in  measuring the cluster optical luminosity and in getting a homogeneous set of total optical luminosities for a large number of clusters. Edge and Stewart  (1991) found that the bolometric X-ray  luminosity  of  a    local sample of    X-ray-selected clusters correlated very roughly with Abell number and somewhat better with the Bahcall galaxy   density (number    of  bright galaxies    within  0.5 $\\rm{h}^{-1}$ Mpc; Bahcall 1977,1981) for the small subsample that had Bahcall galaxy densities.  Arnaud et al.  (1992) made an heroic effort in computing   cluster  optical luminosities at   low redshift  from a heterogenous literature.  The    first joint X-ray/optical survey   of galaxy clusters was  the ROSAT Optical X-ray  Survey (ROXS, Donahue et al.  2002).  They observed 23 ROSAT pointings for a  total of 5 square degrees in the I band and partially in the V band. The X-ray selection and    the optical   selection    of  cluster   candidates were   done independently, with the wavelets algorithm in the former case and with a matched  filter algorithm in the  latter one.   They found X-ray and optical coincident detections for 26  galaxy clusters.  Donahue et al. (2001)  studied  the relation  between  the  X-ray luminosity  and the matched  filter  parameter  $\\Lambda   _{cl}$, which  is approximately equivalent to the  number of $L^*$ galaxies  in the system (Postman et al.  1996).  They  found a marginally significant correlation  between the  two quantities with  a  significant scatter.   Yee and  Ellingson (2003) defined  a  new  richness parameter  as  the  number of cluster galaxies  within some fixed  aperture, scaled by a luminosity function and  a  spatial distribution function.  They   analysed a sample of 15 clusters from  CNOC1  Cluster Redshift  Survey, and  found a very poor correlation between this  parameter and other cluster  properties such as the X-ray luminosity, temperature and the velocity dispersion.  In the present  paper we describe  the properties  and the information contained in the RASS-SDSS   catalog  and the  resulting  correlations between the X-ray and  optical properties in  the sample. In section 2 we explain how the cluster sample  has been created by X-ray selecting the systems from  the available X-ray  cluster and group  catalogs. In section  3 we describe the   method for calculating  the X-ray cluster properties. We describe in section 4 the optical data and in section 5 the  data  reduction method. We  analyse and  discuss the correlations between the optical luminosity and the X-ray  properties in section 6. We summarize and   discuss the catalog properties  and  the results in section 7.  For    all     derived   quantities,   we   have    used $\\rm{H}_0=70$ $\\rm{km}$$\\rm{s}^{-1}$$\\rm{Mpc}^{-1}$, $\\Omega_ {m}=0.3$  and $\\Omega_ {\\lambda}=0.7$. ",
        "conclusions": "We created a  database of clusters  of galaxies  based on  the largest available X-ray  and optical surveys: the ROSAT  All Sky Survey (RASS) and the Sloan Digital Sky Survey (SDSS).  The RASS-SDSS galaxy cluster catalog is the first catalog which combines X-ray and optical data for a  large   number (114) of  galaxy clusters.    The systems  are X-ray selected, ranging from  groups of $10^{12.5}$  $M_{\\odot}$ to massive clusters of  $10^{15}$ $M_{\\odot}$ in  a redshift range from 0.002 to 0.45.    The X-ray  (luminosity     in the   ROSAT band, bolometric luminosity,  redshift, center  coordinates)  and   optical  properties (Schechter luminosity function parameters, luminosity, central galaxy density, core, total and half-light radii) are  computed in a uniform and accurate  way.  The catalog   contains also temperature  and  iron abundance for a subsample of 53 clusters from the Asca Cluster Catalog and the Group Sample. The resulting RASS-SDSS galaxy cluster catalog, then,  constitutes an  important database to  study  the  properties of galaxy clusters   and  in   particular the  relation   of  the  galaxy population seen in the optical to the properties of the X-ray luminous ICM.  The  first  investigations reported  have  shown  a tight  correlation between the X-ray  and optical   properties, when  the choice  of  the measurement   aperture  for the optical     luminosity  and the  optical wavelength band  are optimized.  We  found that the optical luminosity calculated in the i and in the z band correlates better with the X-ray luminosity and the  ICM temperature, in  comparison to the other Sloan photometric   bands.  Thus the    red optical  bands,  which are  more sensitive to the light of the  old stellar population and therefore to the stellar mass of cluster galaxies, have tight correlations with the X-ray properties of the systems.  Moreover,  we   found  that the   scatter  in   the  $L_{op}-L_X$  and $L_{op}-T_X$ relations can be minimized  if the optical luminosity  is measured   within   a    cluster    aperture  between    0.2-0.8   Mpc $\\rm{h}_{70}^{-1}$, with an absolute minimum of the scatter at 0.5 Mpc $\\rm{h}_{70}^{-1}$. The best  aperture   in the  central  part of  the cluster for the  measurement of the optical luminosity  is due  to the fact that it is  a good compromise to  simultaneously assess the total richness and the compactness of the cluster.    Finally   by  using the   relations obtained   from   the  z band,  we demonstrated that, given  the optical properties  of a cluster, we can predict the X-ray luminosity and temperature with  an accuracy of 60\\% and  vice versa.  By    restricting  the correlation  analysis  to  the subsample of X-ray detected REFLEX-NORAS clusters, the minimum scatter decreases to less than 40\\%  for the $L_{op}-L_X$ relation.  Since the observational uncertainties in the optical and in the X-ray luminosity are   about 20\\%, the observed scatters   in both  relations should be intrinsic.  The resulting logarithmic slope for the $L_{op}-L_X$ relation with the minimum  scatter    is   $0.38\\pm 0.02$,  while   the   value  for the $L_{op}-T_X$  relation  is $1.12  \\pm  0.08$.   Both results  are  not consistent  with   the assumption of  hydrostatic  equilibrium   and a constant  M/L.  If we  assume that M/L depends  on the luminosity with the power law $M/L \\propto L^{0.3}$ (Girardi et al. 2002), our results are  in very   good   agreement with  the  expected values   under the assumption of hydrostatic equilibrium.  The analysis carried out in this paper on the correlation between X-ray     and optical  apparence  of  galaxy    clusters is completely empirical. In principle, the best way to proceed in this kind of study is  to measure the optical luminosity  within the physical size of the cluster, like the virial radius. Without optical spectroscopic data or accurate  temperature measurements, the  cluster virial  radius can be calculated by assuming  a    theoretical model relating   the  optical luminosity to the  cluster   mass.  At this stage   of  the  work,  we preferred  to tackle the  cluster   X-ray-optical connection with  the empirical method   explained    in the    paper,  in  order   to  have model-independent results. On the other hand,  not taking into account the different cluster sizes could have affected both the slope and the scatter of the given relations (eq.  \\ref{eq1}, \\ref{eq2} and \\ref{eq3}).  Therefore,  for a better   understanding of the important connection between the X-ray and optical apparence of galaxy clusters, the optical luminosity has to  be calculated within the physical  size of the cluster.  This work is in progress and will be published in the second  paper   of this  series about   the  RASS-SDSS  galaxy cluster sample. The next step will be the  study of the fundamental plane of galaxy clusters. Through this kind of analysis we will find out if the observed scatter  in the correlations   between the optical  and X-ray properties  depends  on  another  parameter   related to the   cluster compactness.  Moreover, because of the link between the galaxy cluster fundamental plane and the M/L parameter,  we will connect directly the slope of two relations to the behavior of M/L.  \\vspace{2cm}  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,  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, Princeton University, the United States Naval Observatory, and the University of Washington."
    },
    "0403/astro-ph0403431_arXiv.txt": {
        "abstract": "Data mining is an important and challenging problem for the efficient analysis of large astronomical databases and will become even more important with the development of the Global Virtual Observatory. In this study, learning vector quantization (LVQ), single-layer perceptron (SLP) and support vector machines (SVM) were put forward for multi-wavelength data classification. A feature selection technique was used to evaluate the significance of the considered features to the results of classification. From the results, we conclude that in the situation of less features, LVQ and SLP show better performance. In contrast, SVM shows better performance when considering more features. The focus of the automatic classification is on the development of efficient feature-based classifier. The classifiers trained by these methods can be used for preselecting AGN candidates. ",
        "introduction": "Today, there are many impressive archives painstakingly constructed from observations associated with an instrument. The Hubble Space Telescope (HST), the Chandra X-Ray Observatory, the Sloan Digital Sky Survey (SDSS), the Two Micron All Sky Survey (2MASS), and the Digitized Palomar Observatory Sky Survey (DPOSS) are examples of this. Furthermore yearly advances in electronics bring new instruments, doubling the amount of data we collect each year. For example, approximately a gigapixels is deployed on all telescopes today, and new gigapixel instruments are under construction. This trend is bound to continue. Just like what Szalay says, the astronomy is facing \"data avalanche\" (See e.g., Szalay \\& Gray 2001). How to organize, use, and make sense of the enormous amounts of data generated by today's instruments and experiments? It is very time consuming and demands high quality human resources. Therefore, better features and better classifiers are required. In addition, expert systems are also useful to get quantitative information.  It is possible to solve the above questions with Neural Networks (NNs), because they permit application of expert knowledge and experience through network training. Furthermore, astronomical object classification based on neural networks requires no priori assumptions or knowledge of the data to be classified as some conventional methods need. Neural networks , over the years, have proven to be a powerful tool capable to extract reliable information and patterns from large amounts of data even in the absence of models describing the data (cf. Bishop 1995) and are finding a wide range of applications also in the astronomical community: catalogue extraction (Andreon et al. 2000), star/galaxy classification (Odewahn et al. 1992; Naim et al. 1995; Miller \\& Coe 1996; M$\\ddot {a}$h$\\ddot {o}$nen \\& Hakala 1995; Bertin \\& Arnout 1996; Bazell \\& Peng 1998), galaxy morphology (Storrie-Lombardi et al. 1992; Lahav et al. 1996), classification of stellar spectra (Bailer-Jones et al. 1998; Allende Prieto et al. 2000; Weaver  2000). Just to name a few, the rising importance of artificial neural networks is confirmed in this kind of task. There is also a very important and promising recent contribution by Andreon et al. (2000) covering a large number of neural algorithms.  In this work, a class of supervised neural networks called learning vector quantization (LVQ) was proposed. LVQ shares the same network architecture as the Kohonen self-organizing map (SOM), although it uses a supervised learning algorithm. Bazell \\& Peng (1998) pioneered the use of it in astronomical applications. Another class of supervised neural networks named multi-layer perceptrons (MLP) was presented. Goderya \\& McGuire (2000) summarized progress made in the development of automated galaxy classifiers using neural networks including MLP. Qu et al. (2003) experimented and compared multi-layer perceptrons (MLP), radial basis function (RBF), and support vector machines (SVM) classifiers for solar-flare detection. Meanwhile, an automated algorithm called support vector machines (SVM) for classification was introduced. The approach was originally developed by Vapnik (1995). Wozniak et al. (2001) and Humphreys et al. (2001) have pioneered the use of SVM in astronomy. Wozniak et al. (2001) evaluated SVM, K-means and Autoclass for automated classification of variable stars and compared their effectiveness. Their results suggested a very high efficiency of SVM in isolating a few best defined classes against the rest of the sample, and good accuracy for all classes considered simultaneously. Humphreys et al. (2001) used different classification algorithms including decision trees, K-nearest neighbor and support vector machines for classifying the morphological type of the galaxy. Furthermore, they got the very promising results of their first experiments with different algorithms.  Celestial objects radiate energy over an extremely wide range of wavelengths from radio waves to infrared, optical to ultraviolet, X-ray and even gamma rays. Each of these observations carries important information about the nature of objects. Different physical processes show different properties in different bands. Based on these, we apply learning vector quantization (LVQ), single-layer perceptron (SLP) and support vector machines (SVM) to classify AGNs, stars and normal galaxies with data from optical, X-ray, infrared bands. In this paper we present the principles of LVQ, SLP and SVM in section 2. In section 3, we discuss the sample selection and analysis the distribution of parameters. In section 4 the computed results and discussion are given. Finally, in section 5 we conclude this paper with a discussion of general technique and its applicability. ",
        "conclusions": "Sources classification depends on the quality and amount of real-time data and on the algorithm used to extract generalized mappings. Availability of the high-resolution multi-wavelength data constantly increases. The best possible use of this observational information requires efficient processing and generalization of high-dimensional input data. Moreover, good feature selection techniques, as well as good data mining methods, are in great demand. A very promising algorithm that combines the power of the best nonlinear techniques and tolerance to very high-dimensional data is support vector machines (SVM). In this work we have used histogram as the feature selection technique and applied LVQ, SLP and SVM to multi-wavelength astronomy to classify AGNs from stars and normal galaxies. We conclude that the features selected by histogram are applicable and the performance of SVM models can be comparable to or be superior to that of the NN-based models in the high dimensional space. The advantages of the SVM-based techniques are expected to be much more pronounced in future large multi-wavelength survey, which will incorporate many types of high-dimensional, multi-wavelength input data once real-time availability of this information becomes technologically feasible. All these methods can be used for astronomical object classification, data mining and preselecting AGN candidates for large survey, such as the Large Sky Area Multi-Object Fiber Spectroscopic Telescope (LAMOST). Various data, incuding morphology, photometry, spectral data and so on, can be applied to train the methods and obtain classifiers to classify astronomical objects or preselect intresting objects. When lacking training sets, we may explore some unsupervised methods or outlier finding algorithms to find unusual, rare, or even new types of objects and phenomena. In addition, with the development of the Virtual Observatory, these methods will be part of the toolkits of the International Virtual Observatory."
    },
    "0403/hep-th0403140_arXiv.txt": {
        "abstract": " ",
        "introduction": "In AdS/CFT, the $O(2,4)$ isometry group of anti-de Sitter space is reinterpreted as conformal symmetry of the dual theory. For de Sitter space, we do not yet possess the holographically dual theory; nevertheless, we would expect the $O(1,d)$ isometry group of de Sitter space to be a symmetry of that theory. But, unlike AdS, de Sitter space also has horizons, with an associated entropy. There are differing interpretations of the entropy but one view, which we shall explore here, is that it indicates that the dual theory has only a finite number of states\\cite{tombanks,fischler}.  However, there is a well-known theorem that says that there are no nontrivial finite-dimensional unitary representations of a noncompact group. Consequently, the Fock space of the dual theory can be either finite-dimensional or a unitary representation of the de Sitter group, but not both. There have thus been claims that de Sitter space has no holographic dual\\cite{gks}, or that the symmetry group is not the de Sitter group\\cite{Witten,alberto}.  Here we will consider the possibility that the dual to de Sitter space is indeed defined by a finite-dimensional Fock space, so that the entropy is finite. The finiteness of the Fock space requires that particles in the dual theory obey Fermi-Dirac statistics. Hence we shall construct a spinor dual to de Sitter space. The finitely many Fock states describing a static patch of de Sitter space do not form representations of the de Sitter group but only of the (compact) rotation subgroup. Global de Sitter space is described by two copies of the Fock space. After antipodal identification, these become the space of initial and final states. Some of the de Sitter generators then mix the in- and out-states i.e. bras with kets. The main idea is that there are certain pairings of bras and kets that are nevertheless de Sitter-invariant. These can be used to construct a de Sitter-invariant S-matrix so, in that sense, the isometry group of de Sitter space indeed appears as a symmetry of the dual theory. In this paper, we will illustrate the reconciliation of symmetry and entropy with a toy model based on Dirac spinors; a more refined construction involving fuzzy spheres will appear in a companion paper\\cite{finitereps}. ",
        "conclusions": ""
    },
    "0403/astro-ph0403695_arXiv.txt": {
        "abstract": "A primordial cosmological magnetic field induces and supports vorticity or Alfv\\'en waves, which in turn generate cosmic microwave background (CMB) anisotropies. A homogeneous primordial magnetic field with fixed direction induces correlations between the $a_{l-1,m}$ and $a_{l+1,m}$ multipole coefficients of the CMB temperature anisotropy field. We discuss the constraints that can be placed on the strength of such a primordial magnetic field using CMB anisotropy data from the WMAP experiment. We place 3 $\\sigma$ upper limits on the strength of the magnetic field of $B < 15$ nG for vector perturbation spectral index $n=-5$ and $B<1.7$ nG for $n=-7$. ",
        "introduction": "The origin of the large-scale part of observed galactic magnetic fields, of $\\sim \\mu$G (microGauss) strength and apparently coherent over $\\sim 10$ kpc scales, is unknown. They could be the consequence of nonlinear amplification of a tiny seed field by galactic dynamo processes. An alternate possibility is amplification of a weak seed field through anisotropic protogalactic collapse and subsequent further amplification via galactic differential rotation. In both cases a primordial seed field of strength exceeding $10^{-13}$ to $10^{-12}$ G, coherent over $\\sim $ Mpc scales, is apparently needed,  and it is often suggested that upto a $\\sim $ nG strength seed field might be required. See Kulsrud (1999), Widrow (2002), and Giovannini (2003) for reviews of the state of the art in this area. A primordial magnetic field of present strength $\\sim $ nG can leave observable signatures in the cosmic microwave background (CMB) anisotropy.  In standard cosmologies vorticity perturbations decay and so do not contribute to CMB temperature or polarization anisotropies. The presence of a cosmological magnetic field generated during an earlier epoch\\footnote{ Quantum fluctuations during an early epoch of inflation can generate a primordial nG magnetic field, coherent over very large scales (see, e.g., Ratra 1992; Bamba \\& Yokoyama 2004)} changes this situation: a primordial magnetic field induces and supports vorticity or Alfv\\'en waves (Adams et al.~1996; Durrer, Kahniashvili, \\& Yates 1998, hereafter DKY). These vector perturbations generate CMB anisotropies.\\footnote{ In the future one can hope to constrain vector modes through their effect on CMB anisotropy polarization anisotropy spectra. CMB polarization spectra that result from vector perturbations due to a primordial magnetic field have been discussed by Seshadri \\& Subramanian (2001), Pogosian, Vachaspati, \\& Winitzki (2002), Mack, Kahniashvili, \\& Kosowsky (2002), and Subramanian, Seshadri, \\& Barrow (2003), while Lewis (2004) considers the case of vector modes supported by free-streaming neutrinos.} The presence of a preferred direction due to a homogeneous magnetic field background leads to an $m$ dependence of $\\langle a_{lm}a_{lm}^*\\rangle$, and induces correlations between the $a_{l+1,m}$ and $a_{l-1,m}$ multipole coefficients of the CMB temperature anisotropy field. Since the CMB anisotropies are observed to be random Gaussian\\footnote{ Colley, Gott, \\& Park (1996), Mukherjee, Hobson, \\& Lasenby (2000), and Park et al.~(2001) are some early discussions of the Gaussianity of the CMB anisotropy. More recent discussions of the Gaussianity of the WMAP CMB anisotropy data are in Komatsu et al.~(2003), Colley \\& Gott (2003), Chiang et al.~(2003), Park (2004), Eriksen et al.~(2004a, 2004b), Coles et al.~(2004), Vielva et al.~(2004), Copi, Huterer, \\& Starkman (2003), Hansen et al.~(2004), Gurzadyan et al.~(2004), and Mukherjee \\& Wang (2004). The simplest inflation models predict Gaussian fluctuations (see, e.g., Fischler, Ratra, \\& Susskind 1985; Ratra 1985), and this is consistent with most observational indications (see, e.g., Peebles \\& Ratra 2003). While there are indications of mild peculiarities in some subsets of the WMAP data, for example the apparent paucity of large-scale power (Spergel et al.~2003; see G\\'orski et al.~1998a for a similar indication from COBE data) and the differences between data from different parts of the sky (see papers cited above), foreground contamination (see, e.g., Park, Park, \\& Ratra 2002; Mukherjee et al.~2002, 2003; de Oliveira-Costa et al.~2003; Bennett et al.~2003b; Tegmark, de Oliveira-Costa, \\& Hamilton 2003) and other systematics might be responsible for part of this.}, it is known that such a contribution can only be subdominant.  We use the observed $\\overline{\\langle a_{l-1,m}a_{l+1,m}^* \\rangle}$ correlations measured by WMAP to place constraints on the strength of a homogeneous primordial magnetic field. The angular brackets here denote an ensemble average, and the overbar indicates an average over $m$ for each $l$. Limited by cosmic variance uncertainties, this would be the useful measure to characterize the signature of a homogeneous primordial magnetic field.  The model on which we base our analysis is introduced in \\S 2. Our analysis of the WMAP data and our results are discussed in \\S 3. We conclude in \\S 4. ",
        "conclusions": "We study off-diagonal correlations of the form $D_l=\\overline{\\langle a_{l-1,m}a_{l+1,m}^* \\rangle}$ in the first year WMAP CMB anisotropy data. Such correlations can result from a homogeneous primordial magnetic field. We do not find significant off-diagonal correlations in the data, which appear to be satisfactorily fit by a zero primordial magnetic field hypothesis. We place 3 $\\sigma$ upper limits on the strength of the magnetic field of $B < 15$ nG for spectral index $n=-5$ and $B<1.7$ nG for $n=-7$. These two cases are interesting as they correspond to a Harrison-Peebles-Yu-Zel'dovich scale-invariant spectrum result for the $C_l$'s and $D_l$'s, and to a possible inflation model primordial magnetic field perturbation spectrum, respectively. These two cases also span the range of constraints that can be placed on $B$ using this method. Future CMB anisotropy data should allow for tighter constraints on a primordial cosmological magnetic field.  \\bigskip  We acknowledge useful discussions with R.~Durrer and A.~Kosowsky. GC, TK, and BR acknowledge support from NSF CAREER grant AST-9875031 and DOE EPSCoR grant DE-FG02-00ER45824. TK also acknowledges CRDF-GRDF grant 3316. PM and YW acknowledge support from NSF CAREER grant AST-0094335."
    },
    "0403/astro-ph0403140_arXiv.txt": {
        "abstract": "We have followed the evolution of multi-mass star clusters containing massive central black holes by N-body simulations on the GRAPE6 computers of Tokyo University. We find a strong cluster expansion and significant structural changes of the clusters. Clusters with IMBHs have power-law density profiles $\\rho \\sim r^{-\\alpha}$ with slopes $\\alpha=1.55$ inside the influence sphere of the central black hole. This leads to a constant density profile of bright stars in projection, which rules out the presence of intermediate mass black holes in core collapse clusters. If the star clusters are surrounded by a tidal field, a central IMBH speeds up the destruction of the cluster until a remnant of a few hundred stars remains, which stays bound to the IMBH for a long time.  We also discuss the efficiency of different detection mechanisms for finding IMBHs in star clusters. ",
        "introduction": "X-ray observations of starburst and interacting galaxies have revealed a class of ultra-luminous X-ray sources (ULX), with luminosities of order $L\\approx 10^{39}$ to $10^{41}$ \\citep{Makishimaetal2000}. If the flux is radiated isotropically, this exceeds the Eddington luminosities of stellar mass black holes by orders of magnitude, making ULX good candidates for IMBHs. Many ULX appear to be associated with star clusters \\citep{Fabbianoetal1997}, the irregular galaxy M82 for example hosts an ULX with luminosity $L > 10^{40}$ erg/sec near its center \\citep{Matsumotoetal2001, Kaaretetal2001} whose position coincides with that of the young ($T \\approx 10$ Myrs) star cluster MGG-11. \\citet{PortegiesZwartetal2004a} and \\citet{McMillanetal2004} have performed $N$-body simulations of several star clusters in M82 and showed that runaway merging of massive stars could have led to the formation of an IMBH with a few hundred to a few thousand solar masses in MGG-11, thereby explaining the presence of the ultraluminous X-ray source.  The fact that a considerable fraction of star cluster might have formed intermediate mass black holes (IMBHs) has interesting consequences. For example, IMBHs of a few 100 to a few 1000 $M_\\odot$ would explain why the mass-to-light ratios in several globular clusters increase towards the center \\citep{Gerssenetal2002, Colpietal2003}, although the data presented so far is also compatible with an unseen concentration of neutron stars and heavy mass white dwarfs \\citep{Baumgardtetal2003}. IMBHs in star clusters would also be prime targets of the forthcoming generation of ground and space-based gravitational wave detectors and could provide the missing link between the stellar mass black holes formed as the end product of stellar evolution and the $10^6$ to $10^9 M_\\odot$ sized black holes found in galactic centers \\citet{Ebisuzakietal2001}.  In this paper we explore the dynamical evolution of star clusters containing massive black holes. We study how star clusters evolve during a Hubble time and compare the outcome of our simulations with galactic globular clusters in order to determine which clusters are likely to contain IMBHs.  We also study what is left bound to an IMBH after the parent cluster is dissolved and discuss ways how to detect an IMBH in a globular cluster. ",
        "conclusions": "We have performed two sets of $N$-body simulations of multi-mass star clusters containing intermediate mass black holes. We found that the 3-dimensional mass-density follows a $\\rho \\sim r^{-1.55}$ density profile around the central black hole. When viewed in projection, the luminosity profile of clusters with massive black holes has a constant density core. The presence of intermediate mass black holes in core collapse globular clusters like M15 is therefore ruled out by our simulations. As was shown in \\citet{Baumgardtetal2003}, a more natural explanation for mass-to-light ratios that increase towards the center in such clusters is a dense concentration of neutron stars, white dwarfs and stellar mass black holes. The detection of a central black hole through proper motion or radial velocity measurements of stars in the central cusp around the black hole is possible with HST for the nearest globular clusters. It might also be possible to find black holes in globular clusters by their gravitational wave emission. The detection through the X-ray emission arising from the IMBH is possible only after the tidal disruption of a star or when a star captured through tidal heating is in a close enough orbit to the IMBH. Intermediate mass black holes also speed up the dissolution of star clusters if the clusters are surrounded by a tidal field."
    },
    "0403/astro-ph0403230_arXiv.txt": {
        "abstract": "An Artificial Neural Network (ANN) scheme has been employed that uses a supervised back-propagation algorithm to classify 2000 bright sources from the Calgary database of IRAS (Infrared Astronomical Satellite) spectra in the region $8\\mu$m to $23\\mu$m. The database has been classified into 17 predefined classes based on the spectral morphology. We have been able to classify over 80 percent of the sources correctly in the first instance. The speed and robustness of the scheme will allow us to classify the whole of the LRS database, containing more that 50,000 sources, in the near future. ",
        "introduction": "Infrared Astronomical Satellite Low Resolution Spectrometer (LRS) recorded spectra of some 50,000 sources in blue ($8 - 15\\mu$m) with $\\lambda/\\Delta\\,\\lambda \\sim 40$ and in red ($13 - 23\\mu$m) with a resolution $\\sim 20$. A total of 5425 objects with better quality spectra were included in the Atlas of low-resolution IRAS spectra (1986, hereafter the Atlas). Volk \\& Cohen (1989a) published spectra of 356 IRAS point sources with $\\rm F_{\\nu} (12 \\mu m) > 40 Jy$ that were not included in the Atlas. These brighter sources were classified into nine classes based upon the spectral morphology. Sixty percent of the sources have silicate emission and red-continuum spectra associated with H II region sources. No emission-line sources formed part of the set of 356 spectra. This sample was also used to test the classification scheme of IRAS sources based on broad-band colors. Classifiable spectra were found for 338 of the sources in the sample. The remaining 18 sources had either extremely noisy or incomplete spectra. It was found that some class of sources overlapped on the color-color diagrams and, therefore, the nature of some of the IRAS sources could not be determined from the IRAS photometry.  Volk et al. (1991) published an additional 486 spectra belonging to sources with $\\rm 12 \\mu m$ fluxes between $20$ and $\\rm 40 Jy$ that were also not in the Atlas. Classifiable spectra were found for 424 sources. The spectra were classified into nine groups as in Volk \\& Cohen (1989a) that describe the astrophysical nature of these sources. Kwok, Volk \\& Bidelman (1997) processed 11,224 spectra (including sources in the Atlas), corresponding to a flux limit of $\\rm 7 Jy$ at $\\rm 12 \\mu m$. These spectra were also classified by-eye and put into nine classes based on the presence of emission and absorption features and on the shape of the continuum. They identified optical counterparts of these IRAS sources in the existing optical and infrared catalogs and listed the optical spectral types if they were known.  It is evident that large databases like the one referred to above require automated schemes for any analysis. Artificial Neural Networks (ANN) have been employed extensively in several branches of Astronomy for automated data analysis (Lahav \\& Storrie-Lombardi, 1994). ANN have been used previously by the IUCAA group in three distinct areas of stellar Astronomy. They have been applied to classify digitized optical and ultraviolet spectra (Gulati et. al. 1994a,b; Singh et al. 1998); to compare a set of observed spectra of F \\& G dwarfs with a library of synthetic spectra (Gulati et al. 1997a, Gupta et al. 2001); and to determine the reddening properties of hot stars from the low-dispersion ultraviolet spectra (Gulati et al. 1997b).  We have attempted to classify 2000 brightest sources from the Atlas into 17 classes by means of Artificial Neural Networks. In the next section we describe features of the 17 new classes. In Section 3, we present details of the ANN scheme. Results are discussed in Section 4 and important conclusions of the study are presented in the last section. Further, an appendix has been added just before the reference section, to explain the general ANN architecture for the benefit of the readers. ",
        "conclusions": "We have demonstrated in this paper the application of ANN scheme to a large database of 2000 IRAS spectra and have been able to correctly classify more than 80\\% of the data-set. The misclassified spectra were looked in detail and most of them could have been wrongly catalogued or had features which would have been confusing for even a human classifier.   We stress here that the speed and robustness of this scheme can be very useful for classifying the whole of LRS database containing over 50,000 sources."
    },
    "0403/gr-qc0403088_arXiv.txt": {
        "abstract": "Japanese laser interferometric gravitational wave detectors, TAMA300 and LISM, performed a coincident observation during 2001. We perform a coincidence analysis to search for inspiraling compact binaries. The length of data used for the coincidence analysis is 275 hours when both TAMA300 and LISM detectors are operated simultaneously. TAMA300 and LISM data are analyzed by matched filtering, and candidates for gravitational wave events are obtained. If there is a true gravitational wave signal, it should appear in both data of detectors with consistent waveforms characterized by masses of stars, amplitude of the signal, the coalescence time and so on. We introduce a set of coincidence conditions of the parameters, and search for coincident events. This procedure reduces the number of fake events considerably, by a factor $\\sim 10^{-4}$ compared with the number of fake events in single detector analysis. We find that the number of events after imposing the coincidence conditions is consistent with the number of accidental coincidences produced purely by noise. We thus find no evidence of gravitational wave signals. We obtain an upper limit of 0.046 /hours (CL $= 90 \\%$) to the Galactic event rate within 1kpc from the Earth. The method used in this paper can be applied straightforwardly to the case of coincidence observations with more than two detectors with arbitrary arm directions. ",
        "introduction": "In the past several years, there has been substantial progress in gravitational wave detection experiments by the ground-based laser interferometers, LIGO\\cite{ref:LIGOdescription}, VIRGO\\cite{ref:virgo}, GEO600\\cite{ref:geo}, and TAMA300\\cite{ref:tama, ref:ando}. The observation of gravitational waves will not only be a powerful tool to test general relativity, but also be a new tool to investigate various unsolved astronomical problems and to find new objects which were not seen by other observational methods.  The Japanese two laser interferometers, TAMA300 and LISM, performed a coincident observation during August 1 and September 20, 2001 (JST). Both detectors showed sufficient stability that was acceptable for an analysis to search for gravitational wave signals. Given the sufficient amount of data, it was a very good opportunity to perform a coincidence analysis with real interferometers' data.  There were several works to search for gravitational waves using interferometeric data. A coincidence analysis searching for generic gravitational wave bursts in a pair of laser interferometers has been reported in \\cite{ref:Nicholson}. Allen et al.\\cite{ref:40m} analyzed LIGO 40m data and obtained an upper limit of 0.5/hour (CL = 90$\\%$) on the Galactic event rate of the coalescence of neutron star binaries with mass between 1$M_\\odot$ and 3$M_\\odot$. Tagoshi et al.\\cite{ref:DT2} analyzed TAMA300 data taken during 1999 and obtained an upper limit of 0.59/hour (CL = 90$\\%$) on the event rate of inspirals of compact binaries with mass between 0.3$M_\\odot$ and 10$M_\\odot$ and with signal-to-noise ratio greater than 7.2. Very recently, an analysis using the first scientific data of the three LIGO detectors was reported~\\cite{ref:LIGOinspiral}, and an upper limit of $1.7\\times 10^2$ per year per Milky Way Equivalent Galaxy is reported. Recently, International Gravitational Event Collaboration (IGEC) of bar detectors reported their analysis using four years of data to search for gravitational wave bursts \\cite{ref:IGEC03}. They found that the event rate they obtained was consistent with the background of the detectors' noise.   In the matched filtering analysis using real data of single laser interferometer (e.g. \\cite{ref:40m}, \\cite{ref:DT2}), many fake events were produced by non-Gaussian and non-stationary noise. In order to remove such fake events, it is useful to perform coincidence analysis between two or more independent detectors. Furthermore, coincidence analysis is indispensable to confirm the detection of gravitational waves when candidates for real gravitational wave signals are obtained. The purpose of this paper is to perform coincidence analysis using the real data of TAMA300 and LISM.  We consider gravitational waves from inspiraling compact binaries, comprized of neutron stars or black holes. They are consider to be one of the most promising sources for ground based laser interferometers. Since the waveforms of the inspiraling compact binaries are known accurately, we employ the matched filtering by using the theoretical waveforms as templates. Matched filtering is the optimal detection strategy in the case of stationary and Gaussian noise of detector. However, since the detectors' noise is not stationary and Gaussian in the real laser interferomters, we introduce $\\chi^2$ selection method to the matched filtering.  We analyze the data from each detector by matched filtering which produces event lists. Each event is characterized by the time of coalescence, masses of the two stars, and the amplitude of the signal. If there is a real gravitational wave event, there must be an event in each of the event list with consistent values of parameters. We define a set of coincidence conditions to search for coincident events in the two detectors. We find that we can reduce the number of events to about $10^{-4}$ times the original number. The coincidence conditions are tested by injecting the simulated inspiraling waves into the data and by checking the detection efficiency. We find that the detection efficiency is not affected significantly by imposing the coincidence conditions.  We estimate the number of coincident events produced accidentally by the instrumental noise. By using a technique of shifting the time series of data artificially, we find that the number of events survived after imposing the coincidence conditions is consistent with the number of accidental coincidences produced purely by noise.  We propose a method to set an upper limit to the real event rate using results of the coincidence analysis. In the case of TAMA300 and LISM, we obtain an upper limit of the event rate as 0.046/hour (CL = 90$\\%$) for inspiraling compact binaries with mass between 1$M_\\odot$ and 2$M_\\odot$ which are located within 1kpc from the Earth. In this case, since TAMA300 is much more sensitive than LISM, the upper limit obtained from the coincidence analysis is less stringent than that obtained from the TAMA300 single detector data analysis. This is because the detection efficiency in the coincidence analysis is determined by the sensitivity of LISM. Thus, the upper limit obtained here is not the optimal one which we could obtain using the TAMA300 data taken during 2001.  The method to set an upper limit to the event rate proposed here can be extended straightforwardly to the case of a coincidence analysis for a network of interferometric gravitational wave detectors.  This paper is organized as follows. In Section~\\ref{sec:detector}, we briefly describe the TAMA300 and LISM detectors. In Section~\\ref{sec:method}, we discuss a method of matched filtering search used for TAMA300 and LISM data. In Section~\\ref{sec:matchedfilterresults}, the results of the matched filtering search for each detector are shown. In section~\\ref{sec:coincidence}, we discuss a method of the coincidence analysis using the results of single-detector searches, and the result of the coincidence analysis is shown. We also derive the upper limit to the event rate in Section~\\ref{sec:limit}. Section~\\ref{sec:summary} is devoted to summary. In Appendix A, we discuss a $\\chi^2$ veto method to distinguish between real events and fake events produced by non-Gaussian noise. In Appendix B, we examine a different choice of $\\Delta t$ ( the length of duration to find local maximum of matched filtering output ) for comparison. In Appendix C, we discuss a sidereal time distribution of coincidence events. In Appendix D, we review a method to estimate the errors in the parameters due to noise using the Fisher matrix.  Throughout this paper, the Fourier transform of a function $h(t)$ is denoted by $\\tilde{h} (f)$, which is defined by \\begin{equation} \\tilde{h} (f) = \\int_{-\\infty}^{\\infty} dt\\ e^{2 \\pi i f t} h(t) . \\end{equation} ",
        "conclusions": "\\end{table}   \\begin{table}[htpd] \\begin{center} \\renewcommand{\\arraystretch}{1.5} \\begin{tabular}{c c c c c} \\hline \\hline & Year  & period & obsevation time [hours] &Topics   \\\\ \\hline DT1&1999& 6-7 Aug. & 11 &Total detector system check \\\\ &&&&and  Calibration test\\\\ DT2&1999& 17-20 Sept.&31& First event search \\\\ DT3&2000& 20-23 April&13& Sensitivity improved\\\\ DT4&2000& 21 Aug.-4 Sept.&167 &100 hours observation\\\\ DT5&2001& 2-10 Mar. &111& Full time observation\\\\ DT6&2001& 1 Aug.-20 Sept.&1038& 1000 hours observation \\\\ &&&&and coincident observation with LISM\\\\ DT7&2002& 31 Aug.-2 Sept.&25& Power recycling installed (Full configuration)\\\\ DT8&2003& 14 Feb.-14 April & 1158 & Coincident observation with LIGO \\\\ DT9&2003 - 2004&  28 Nov.- 10 Jan. &  557 & Full automatic operation\\\\ &  &  & &and Partial coincident observation \\\\ &&&&with LIGO and GEO600 \\\\ \\hline \\hline \\end{tabular} \\end{center} \\caption{Observation history of TAMA300} \\label{tab:history} \\end{table}  \\subsection{LISM}\\label{sec:detectorl}  LISM is a laser interferometer gravitational wave antenna with arm length of 20m, located in the Kamioka mine ($36.25^\\circ$N, $137.18^\\circ$E), 219.02km west of Tokyo. The detector's arm orientation is $165^\\circ$ measured counter clockwise from East. The LISM antenna was originally developed as a prototype detector from 1991 to 1998 at the National Astronomical Observatory of Japan, in Mitaka, Tokyo, to demonstrate advanced technologies \\cite{ref:sato}. In 1999, it was moved to the Kamioka mine in order to perform long-term, stable observations. Details of the LISM detector is found in \\cite{ref:satoLISM}.  The laboratory site is 1000m underground in the Kamioka mine. The primary benefit of this location is extremely low seismic noise level except artificial seismic excitations. Furthermore, much smaller environmental variations at this underground site are beneficial to stable operation of a high-sensitivity laser interferometer. The optical configuration is the Locked Fabry-Perot interferometer. The finesse of each arm cavity was about 25000 to have a cavity pole frequency of 150Hz. The main interferometer was illuminated by a Nd:YAG laser yielding 700mW of output power, and the detector sensitivity spectrum was shot-noise limited at frequencies above about 1kHz.  The operation of LISM was started in early 2000, and has repeatedly been tested and improved since. The data used in this analysis were taken in the observations between August 1st and 23th and between September 3rd and 17th, 2001. The total length of data is 780 hours. The first half of the period was in a test-run and some improvements were made after that. The data from the second half were of good quality to be suitable for a gravitational wave event search, so 323 hours of data for the latter half was dedicated for this analysis. The best sensitivity during this period was about $h\\sim 6.5 \\times 10^{-20}/\\sqrt{\\rm{Hz}}$ around 800Hz.   \\label{sec:summary}  In this paper, we performed a coincidence analysis using the data of TAMA300 and LISM taken during DT6 observation in 2001.  We analyzed the data from each detector by matched filtering and obtained event lists. Each event in the lists was characterized by the time of coalescence, masses of the two stars, and the amplitude of events. If any of the events are true gravitational wave events, they should have the consistent values of these parameters in the both event lists. We proposed a method to set coincidence conditions for the source parameters such like the time of coalescence, chirp mass, reduced mass, and the amplitude of events. We took account of the time delay due to the distance between the two detectors, the finite mesh size of the mass parameter space, the difference in the signal amplitudes due to the different sensitivities and antenna patterns of the detectors, and errors in the estimated parameters due to the instrumental noise. Our Monte Carlo studies showed that we would not lose events significantly by imposing the coincidence conditions.  By applying the above method of the coincidence analysis to the event lists of TAMA300 and LISM, we can reduce the number of fake events by a factor $10^{-4}$ compared with the number of fake events before the coincidence analysis. In order to estimate the number of accidental coincidences produced by noise, we used the time shift procedure. We found that the number of events survived after imposing the coincidence conditions is consistent with the expected number of accidental coincidences within the statistical fluctuations. Thus we found no evidence of gravitational wave signals. As discussed in Appendix~\\ref{sec:sidereal}, the sidereal time distribution of the survived events were also consistent with the distribution of accidentals.  Finally, we proposed a simple method to set an upper limit to the event rate and applied it to the above results of the coincidence analysis. We obtained an upper limit to the Galactic event rate within 1kpc from the Earth to be 0.046 [1/hour] (90\\%\\,CL). In our case, since  LISM has a much lower sensitivity than TAMA300, we were unable to obtain a more stringent upper limit to the event rate than the one obtained by the single-detector analysis of TAMA300. This is because the detection efficiency in the coincidence analysis is determined by the detector with a lower sensitivity.  However, if we have two detectors that have comparable sensitivities, it is possible to obtain an improved upper limit compared to a single-detector analysis. As an example, let us imagine the case when the sensitivity of LISM is the same as that of TAMA300. The result of Galactic event simulations suggests that the detection efficiency in the case of a single-detector analysis is 0.35, while it improves to 0.48 in the case of a coincidence analysis. These values are translated to upper limits on the Galactic event rate of 0.026 [1/hour] ($90\\%$\\,CL) for the single-detector case and 0.019 [1/hour] ($90\\%$\\,CL) for the two-detector case.  The method of a coincidence analysis and the method to set an upper limit to the event rate proposed here can be readily applied to the case when there are more than two detectors with arbitrary arm directions. Hence these methods will be useful for data analysis for a network of interferometeric gravitational wave detectors in the near future."
    },
    "0403/astro-ph0403006_arXiv.txt": {
        "abstract": "Gauge invariant quintessence perturbations in the quintessence and cold dark matter ($QCDM$) model are investigated. For three cases of constant equation-of-state (EOS) parameter, linear scalar field potential, and supergravity scalar field potential, their perturbation evolutions have a similar dependence on EOS parameter and scale, but they have different sensitivity to the initial conditions due to the different shapes of the quintessence potential. They have a minor effect on primary CMB anisotropies, but change the secondary CMB effect in different ways. The first case is insensitive to initial quintessence perturbations and only modifies the Integrated Sachs-Wolfe (ISW) effect within a factor of 2. The other two cases are sensitive to initial conditions at large scales and could affect the secondary CMB anisotropies drastically depending on how smooth the initial perturbations are. This makes it possible for future cosmological probes to provide constraints on quintessence properties. ",
        "introduction": "Recent observations of an accelerating universe \\cite{Riess1998, Perlmutter1999, Riess2001} imply the existence of dark energy characterized by a negative pressure to density ratio, also known as the EOS parameter $w$. There have been discussions of the cosmological constant, the dark energy evolving according to a specific EOS parameter \\cite{Silveira1997, Chiba1997}, and the dark energy consisting of a dynamical cosmic scalar field, the quintessence \\cite{Ratra1988, Ferreira1997}. In contrast to the cosmological constant, the quintessence component in $QCDM$ model has kinematic behavior and can develop perturbations. It has important features of time-evoluting EOS parameter and scale-dependent effective sound speed \\cite{Caldwell1998a, Caldwell1998b}. One significant issue is the sensitivity to initial perturbation conditions, which attracts much attention.  For some $QCDM$ models, the quintessence potentials can be approximated with constant EOS parameter, and the perturbations evolutions are insensitive to the initial quintessence perturbations, with minor effect on observations \\cite{Dave2002, DaveThesis, Weller2003}. For exponential potential $V(Q)=\\hat{V} e^{-(c/M)Q}$ \\cite{Wetterich1985, Wetterich1988, Ratra1988}, although quintessence perturbations may stay non-zero at some time for different background attractor solutions, they die out in matter-dominant time for physically reasonable models \\cite{Tassilo2001}, which shows the insensitivity again. For the tracking quintessence models \\cite{Brax2000}, large-scale non-adiabatic perturbations can grow or stay constant before entering tracker regime, but have to get suppressed afterwards \\cite{MalquartiLiddle2002}. However, the first-order matrix formulation developed latter \\cite{Bartolo2003} applies to general quintessence potential cases on all linear scales and indicates the possibility of non-vanishing entropy perturbations. Recently some other $QCDM$ models with special quintessence potentials, such as linear $V(Q)$ \\cite{Dimopoulos2003} and supergravity $V(Q)$ \\cite{Kallosh2002}, lead to special background evolutions. So it will be interesting to see the perturbations behavior and their sensitivity to initial conditions in these models, as well as their effects on cosmic microwave background (CMB) \\cite{Bennett2003} and baryon mass power spectrum \\cite{Scranton2003} in the universe.  In this paper, we describe the universe by a simple $QCDM$ model with the conformal-Newtonian metric in section 2. The cosmological perturbations are solved for constant $w$ cases in section 3, for linear-$V(Q)$ cases in section 4, and for supergravity-$V(Q)$ cases in section 5. The results are analyzed by the matrix formulation and their effects on cosmological observations are discussed in section 6. ",
        "conclusions": "In another viewpoint, we can apply the first-order matrix formulation in \\cite{Bartolo2003} to analyze the different quintessence perturbation behaviors in above $QCDM$ models. At large scales ($k \\ll 1$), the relative entropy perturbation $S$ and intrinsic entropy perturbation $\\Gamma$ \\cite{Wands2000, Malik2003, Kodama1984} have coupled evolution equation $$ \\frac{\\partial}{\\partial \\ln(a/a_0)} \\left( \\begin{array}{c} S \\\\ \\Gamma \\end{array} \\right) = 3 \\left( \\begin{array}{cc} w_Q - w_f + \\gamma_f \\Omega_f (w_f - c_{sQ}^2) / \\gamma & \\gamma_f \\Omega_f (1 - c_{sQ}^2) / \\gamma \\\\ - \\gamma /2 & w_Q - \\gamma /2 \\end{array} \\right) \\times \\left( \\begin{array}{c} S \\\\ \\Gamma \\end{array} \\right) $$ where $_f$ means the perfect fluid component, $\\gamma_i \\equiv 1+w_i$, and $c_{sQ}^2 \\equiv \\dot{p} / \\dot{\\rho}$. For most of the time in our $QCDM$ models, $CDM$ dominates and have $w_f \\simeq 0$, $\\Omega_f \\simeq \\gamma_f \\simeq \\gamma \\simeq 1$, and $-1<w_Q<-1/3$. Then the eigenvalues are $n_{\\pm} = \\frac{3}{2} \\left[ 2w_Q - c_{sQ}^2 - \\frac{1}{2} \\pm \\sqrt{(c_{sQ}^2 + \\frac{1}{2})^2 - 2} \\right]$. For the constant $w$ case, $w_Q = c_{sQ}^2$ leads to $Re(n_{\\pm}) = \\frac{3}{2}(w_Q - 1/2) < 0$, corresponding to decaying entropy perturbations. While for the linear $V(Q)$ case, $c_{sQ}^2 \\simeq -2$ leads to $n_{\\pm} = (3w_Q+3,3w_Q+\\frac{3}{2})$, corresponding to at least one growing entropy perturbation mode. This fits our previous analysis well.  In all $QCDM$ models that we studied above, the quintessence perturbations have the common characteristic feature of scale dependence. At small scales the perturbations do not grow and the gravitational potential decays more only due to background effects; at large scales the perturbations grow and sustain the gravitational potential. On the other hand, because of the different quintessence potentials, the three $QCDM$ models have different sensitivity to the initial conditions on quintessence perturbations. The constant-$w$ model has a large quintessence mass at early times and damps perturbations fast so that it is insensitive to initial conditions at all scales. In this model the quintessence has no effect before the last-scattering time. In contrast, the linear-$V(Q)$ model and supergravity-$V(Q)$ model have zero and tachyon quintessence mass respectively, and their perturbations remain near their initial value ($\\delta\\tilde{Q}(0) / \\Phi(0) \\leq 10^4$) at large scales. Since the EOS parameter $w$ in these two models approaches -1 at early times, their dark energy dominates only very recently, like the $\\Lambda CDM$ model. So the total back-reaction of dark energy perturbations (background times relative perturbations) on other perturbations remains negligible before the last-scattering time, when the universe evolves like $a \\propto \\eta^2$, $Q_0 \\simeq Q_0(0)$, $Q_0^{\\prime} \\simeq 0$, $\\delta Q \\simeq \\delta Q(0)$, and $\\Phi \\simeq \\Phi(0)$. So they do not change the primary CMB anisotropies either. But later when the dark energy background become dominant, this back-reaction of dark energy perturbations will affect other perturbations drastically, depending on how smooth the initial perturbations are.  For smooth initial conditions in all models the dark energy has no primary CMB effect except that its background contributes to the distance to the last-scattering surface. But later on both the dark energy background and perturbations will contribute to the change of the gravitational potential, which may lead to a considerable secondary CMB effect such as the ISW effect \\cite{SW1967}. For the scale-invariant primordial spectrum $P_{\\Phi} \\equiv |\\Phi(0,k)|^2 = A^2/k^3$, the large-scale CMB power spectrum due to the SW effect is: \\begin{eqnarray} & & C_l^{SW} = \\frac{A^2}{36\\pi^2 l(l+1)} K_l^2 \\; , \\nonumber \\\\ & & K_l^2 = 2l(l+1) \\int_0^{\\infty} \\frac{dk}{k} \\left[ j_l(k\\eta_0) + 6\\int_{\\eta_{LS}}^{\\eta_0} d\\eta f^{\\prime}(\\eta,k) j_l(k(\\eta_0-\\eta)) \\right]^2 \\; , \\label{ISW_K} \\end{eqnarray} where $j_l(z)$ is the spherical Bessel function and $f(\\eta, k)$ is the evolution factor of the gravitational potential defined as $\\Phi(\\eta, k) \\equiv \\Phi(0, k) f(\\eta, k)$. The integral over the $j_l^2(k\\eta_0)$ term corresponds to the primary SW effect on the last-scattering surface, and the integral over the remaining terms corresponds to the ISW effect. The flat $CDM$ model gives $f^{\\prime}=0$ and only the primary SW effect, with $K_l^2=1$. The flat $\\Lambda CDM$ model has no dark energy perturbations and gives a scale-independent evolution factor $f(\\eta)$ which leads to a positive ISW effect in the low-$l$ CMB power spectrum \\cite{Kofman1985}, rising at the lower-$l$ ($l=2,3$) end. For the flat $QCDM$ model, the bigger EOS parameter $w$ $(\\geq -1)$ brings out the dark energy earlier and decreases the gravitational potential more, corresponding to a bigger $f^{\\prime}$ value for all scales and affecting the magnitude of the CMB power spectrum. The quintessence perturbations grow differently at different scales and give a $k$-dependent $|f^{\\prime}(\\eta,k)|$, which is bigger at small scales than at large scales with the transition around the horizon size. This means that the lower-$l$ ($l=2,3$) multipoles will get suppressed relative to the higher-$l$ ones such that the whole low-$l$ spectrum will get flattened. These two effects can be seen by numerical calculation for the  standard $\\Lambda CDM$ model, constant-$w$ $QCDM$ model with $w=-0.5$, and linear-$V(Q)$ $QCDM$ model with $w(\\eta_0)=-0.5$ in Figure~\\ref{ISW_Figure}. This is done with eq.(\\ref{ISW_K}) approximately for illustrative purpose. Exact results should refer to more accurate numerical work like the CMB-FAST code. For non-smooth initial conditions, the perturbation evolution can be seen in Table~\\ref{LinearV_Initial_Value_Dependence_Table} and Table~\\ref{SugraV_Initial_Value_Dependence_Table}, and corresponds to very unusual cosmological scenarios which we will not discuss further.  \\begin{figure}[ht] \\includegraphics[scale=0.75]{QuintessencePerturbations_ISW_Fig.eps} \\caption{The low-$l$ CMB power spectrum due to the Sachs-Wolfe effect in various $QCDM$ models with the smooth initial condition. The solid line is for the $\\Lambda CDM$ model, the dash line for the constant-$w$ $QCDM$ model with $w=-0.6$, and the dot line for the linear-$V(Q)$ $QCDM$ model with $w(\\eta_0)=-0.5$.} \\label{ISW_Figure} \\end{figure}  The scale dependence and sensitivity to initial perturbations can also be tested in other ways. The CMB results give the primordial $\\Phi_k$ spectrum, while the gravitational lensing and other large-scale observations yield $\\Phi_k$ information at low redshift. The ratio of these two, the evolution of gravitational potential, can be directly compared with the theoretical plot on different scales shown above for various $QCDM$ models. Furthermore, since the matter density perturbation is related to the gravitational potential by eq.(\\ref{MatterEMC.4}), the evolution of gravitational potential can also be inferred from the measurement of the visible matter density spectrum (by SDSS, etc.), and checked with $QCDM$ predictions.  By considering the dark energy perturbations in these $QCDM$ models, we see that they all have a scale dependence which is different from the $\\Lambda CMD$ model, and they have different sensitivity to initial conditions. These properties make the dark energy perturbations an effective way to test these models. More precise cosmological observations in the future will give better constraints on these $QCDM$ models."
    },
    "0403/hep-ph0403019_arXiv.txt": {
        "abstract": "We show that a gravitational interaction between the derivative of the Ricci scalar curvature and the baryon-number current dynamically breaks $CPT$ in an expanding universe and, combined with baryon-number-violating interactions, can drive the universe towards an equilibrium baryon asymmetry that is observationally acceptable. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403283_arXiv.txt": {
        "abstract": "Saturn's largest satellite, Titan, transited the Crab Nebula on 5 January 2003. We observed this astronomical event with the {\\it Chandra} X-ray Observatory. An ``occultation shadow'' has clearly been detected and is found to be larger than the diameter of Titan's solid surface. The difference gives a thickness for Titan's atmosphere of 880 $\\pm$ 60 km. This is the first measurement of Titan's atmospheric extent at X-ray wavelengths. The value measured is consistent with or slightly larger than those estimated from earlier Voyager observations at other wavelengths. We discuss the possibility of temporal variations in the thickness of Titan's atmosphere. ",
        "introduction": " ",
        "conclusions": "DISCUSSION}  Taken with the distance of 1.214 $\\times$ 10$^{9}$ km to Titan at the time of this observation, $R_{shadow} = 0.588 \\pm 0.011$ arcseconds gives a thickness for Titan's atmosphere of $880 \\pm 60$ km. This value can be compared with estimates from the models for Titan's atmosphere which have been constructed based on Voyager~1 observations at radio, IR, and UV wavelengths. Figure~\\ref{fig:comparison}a shows profiles of the tangential column density (cm$^{-2}$) along the line-of-sight as a function of altitude (distance above Titan's surface), which are compiled from density (cm$^{-3}$) profiles provided by Yelle et al.\\ (1997) (red) and Vervack, Sandel, \\& Strobel (manuscript in preparation) (blue). Nitrogen dominates, making up more than 95\\% of the atmospheric constituents, and methane is the second major component, although proportions slightly differ in the two models. In those models, the profiles in the altitude range of 500--1000 km, where no data were available, were interpolated between the data of higher and lower altitudes assuming hydrostatic equilibrium. The step function we used to determine the atmospheric thickness is a simplification of the transmission curve of X-rays through the atmosphere. In reality, X-rays suffer photoelectric absorption as they pass through the atmosphere, with a transmissivity determined by the atmospheric composition, the tangential column density for a given altitude, and photon energy.  We have calculated transmission curves by using the atmospheric compositions of the above two models, the tangential column densities in Figure~\\ref{fig:comparison}a, and integrating over the observed Crab spectrum.  The curves are shown in Figure~\\ref{fig:comparison}b. The step function with a threshold of $R_{shadow}$ is also shown. Figure~\\ref{fig:comparison}c shows a $\\chi^{2}$ curve which is identical with the black curve shown in Figure~\\ref{fig:chi2}, but the definitions of the x-axis are different. We compared the simulation treating the atmosphere as a solid disk (disk model) with the simulations based on the transmission curves shown in Figure~\\ref{fig:comparison}b (atmosphere model) as follows. We calculated $\\chi^{2}$ for the atmosphere model simulations and defined the ``characteristic altitude'' $h$ for those models by  \\begin{equation} h = \\int_{0}^{\\infty}\\, \\left( 1-T(h') \\right) \\,dh', \\end{equation}  \\noindent where $T(h')$ is transmissivity as a function of altitude, $h'$. We find that $h$ = 750 and 780 km for the Yelle et al.\\ and Vervack et al.\\ models, respectively. The points ($h$, $\\chi^{2}$) for those two atmosphere models fall very close to the $\\chi^{2}$ curve for the disk model, as shown in Figure~\\ref{fig:comparison}c.  In other words, our step-function transmissivity is indistinguishable from a more realistic transmissivity with $h$ at the threshold of the step in the simulation.  This fact strongly justifies the approximation of transmission curve with a step function in our simulation and suggests that $R_{shadow}$ and $h$ can be directly compared with each other. Then, our result for the thickness of Titan's atmosphere is consistent with or slightly larger than those estimated from Voyager~1 observations (1.9 and 1.4 $\\sigma$ separation).  Although the statistical significance is not so large, it is still worth discussing a temporal variation of Titan's upper atmosphere. Distances between the Sun and Saturn were 1.42 and 1.35 $\\times$ 10$^{9}$ km at the time of Voyager's encounter (November 1980) and our observation (January 2003), respectively. The solar luminosity is known to be fairly stable with an uncertainty of 0.3\\% regardless of its 11 years active cycle (Fr\\\"{o}ehlich \\& Lean 1998; Quinn \\& Fr\\\"{o}ehlich 1999). Accordingly, the closer distance to the Sun in our observation might have resulted in the higher temperature and the larger extent of the atmosphere.  Although the size of this effect cannot be ascertained without detailed modeling, the increase in solar flux incident on Titan may account for some part of the slightly larger value of Titan's atmosphere measured here. More detailed information will be obtained by the Cassini/Huygens mission in 2005.  Finally, we note that the spectrum taken from the shadow has no absorption nor emission line feature and is statistically indistinguishable from the spectrum taken from the surrounding region. Those facts indicate that photons penetrating the atmosphere hardly contributed to the spectrum taken from the shadow; if they did, excess absorption due to Titan's atmosphere would be seen in the shadow spectrum. The absence of excess absorption is reasonable considering that the PSF is much broader (the HPD corresponds to about 5000 km) than the corresponding atmospheric thickness interval of about 500 km within which the tangential optical depth changes from unity to zero (see Fig.~\\ref{fig:comparison}b)."
    },
    "0403/astro-ph0403180.txt": {
        "abstract": " ",
        "introduction": "MAGIC was designed back in 1998 with the very clear goal to lower the energy threshold at which gamma rays can be observed \\cite{MagicDR}. As the first MAGIC telescope has started operation and is now, in early 2004, on the way to reach its design performance, it is likely that, beyond the physics that has been predicted, unexpected avenues will open. Many unanswered physics questions loom in the low energy range, and we are confident that several of them will already be answered by MAGIC.  The instruments of ECO, in particular the eventual larger telescope(s), will subsequently be able to extend the energy limit to even lower values, close to the lowest energies accessible to the Cherenkov technique, and to fully explore this range of wavelengths.  The energy threshold is primarily a question of two components: photon detection efficiency and mirror surface. Fast electronics resulting in signal integration over a short time is also an important handle to suppress background, and contributes to observing showers at lower energies. A natural evolution of the MAGIC principles, therefore, will be to take advantage of technological progress in light detection and the building of large structures, viz. the construction of a larger and more performant telescope. The mirror surface must be pushed to the technical limit without taking undue risks. To be constructed with a large mirror surface, maintaining a fast slewing possibility, the telescope structure must be made light-weight and easy to assemble. For high sensitivity, photodetectors must be brought to the highest possible quantum efficiency, and the light collection along the optical path must be optimal.  These are goals that have already guided the construction of MAGIC. The addition to MAGIC of further telescopes, on the Roque de los Muchachos site (La Palma), constitutes the initial step in building up a gamma ray observatory, the European Cherenkov Observatory or ECO. The experience gained during the construction of MAGIC, and progress in industrially available components and production methods do allow an improvement program, for a future second telescope of the MAGIC type. A so far rather unchanged copy of MAGIC is already under way.  With some additional support, this telescope can incorporate key innovations in technology, viz. the latest commercially available photodetectors, a new light-weight structure supporting highly reflective aspherical mirrors, and much improved mirror control that does not interfere with observations. Prototypes developed in that framework, and an accompanying feasibility and design study will pave the way for a substantial future leap in telescope size, i.e. towards ECO-1000.  First results from the MAGIC observatory are being obtained now (early 2004), and will likely turn out to be the start of a bonanza of sources understood or discovered, beyond the ones which we are guided to by the EGRET (in the future: GLAST) catalogues. Many sources will require detailed study, for which small-angle telescopes like MAGIC or ECO-1000 are ideally suited.  This report is structured as follows: The physics arguments are presented in the next section \\ref{sec2}. Sections  \\ref{sec_sigbkg} and \\ref{sec3} discuss the characteristics of the gamma signal and the expected sources of backgrounds, and the conceptual choices for a low-energy telescope, based on simulations of gamma ray induced showers and background. We concentrate on a comparison between the present MAGIC and both a MAGIC with a high-QE camera and an ECO-1000 telescope, also with a high-QE camera. In section \\ref{sec:analysis}, we discuss specific problems associated with opening the low energy sector. We then follow, in section \\ref{sec4}, with a discussion of the practical implementation, viz. a crude work plan. ",
        "conclusions": "Multiple strong arguments can be made for the extraordinary physics potential that can be tapped by gamma ray telescopes with the lowest possible energy threshold. A key instrument to aim for is a telescope with a large mirror surface and a high quantum efficiency camera. We call this telescope ECO-1000, and have presented above the necessary steps to work towards such a device over the next three years. One important step is a design and feasibility study, including some prototyping; we propose here to equip part of a MAGIC-size telescope, now under construction, with novel photodetectors, with optimized mirrors, and with a permanently active mirror control system. Use in a working telescope will not only allow long-term field testing of these components, but also improve the physics performance of this telescope. We also propose detailed feasibility studies towards an ultra-light telescope structure, based on computer modeling.  The successful completion of the proposed studies will give us confidence in the new technological features, and will allow a riskless extrapolation from MAGIC towards ECO-1000, as a final target. A proposal for ECO-1000 could be completed as early as 2007, with the goal of having an operating instrument in 2009 or 2010.  We further argue strongly in favor of using the existing infrastructure at the La Palma site for a gradual development towards a multi-telescope European Cherenkov Observatory. We intend to open this observatory to guest observers, and to make the data available to the entire international community of astronomers and astrophysicists."
    },
    "0403/astro-ph0403556_arXiv.txt": {
        "abstract": "{We examine the environmental effects on the photometric properties of galaxies for the rich galaxy cluster ABCG 209 at $z=0.209$. We use archive CFHT optical imaging of a \\mbox{$42\\times28\\,{\\rm arcmin}^{2}$} field centred on the cluster to produce a galaxy sample complete to $B=25.0$ and $R=24.5$. Both the composite and red sequence galaxy luminosity functions are found to be dependent on the local galaxy surface density, their faint-end slopes becoming shallower with increasing density. We explain this as a combination of the morphology-density relation, and dwarf galaxies being cannibalised and/or disrupted by the cD galaxy and the ICM in the cluster core. The $B-R$ colour of the red sequence itself appears 0.02\\,mag redder for the highest-density regions, indicative of their stellar populations being marginally ($<5$\\%) older or ($<20$\\%) more metal-rich. This may be due to the galaxies themselves forming earliest in the rarest overdensities marked by rich clusters, or their star-formation being suppressed earliest by the ICM. ",
        "introduction": "The effect of local environment on the formation and evolution of galaxies is one of the most pressing issues in cosmology today, with evidence that environmental processes affect the mass distribution of galaxies, as well as their star-formation and morphological characteristics. Rich clusters provide a unique opportunity to study these environmental effects, providing large numbers of galaxies at the same redshift which have been exposed to a wide variety of environments. Numerous authors have studied the environmental effects of rich clusters on the luminosity functions, star-formation rates and morphologies of their constituent galaxies (e.g. L\\'{o}pez-Cruz \\etal \\cite{lopezcruz}; Balogh, Navarro \\& Morris \\cite{balogh}; Hogg \\etal \\cite{hogg}; Treu \\etal \\cite{treu}).  The galaxy luminosity function (LF), which describes the number of galaxies per unit volume as a function of luminosity, is a powerful tool for examining galaxy formation and evolution, since it can be directly related to the galaxy mass function. The Press-Schechter prescription for the hierarchical assembly of galaxies predicts a simple analytical formula for the mass distribution of the form \\mbox{${\\rm n(M)dM}\\propto {\\rm M}^{\\alpha}\\exp(-{\\rm M/M}^{*})$} (Press \\& Schechter \\cite{press}) which, despite not considering non-linear dynamical effects, reproduces well the results of N-body simulations of the growth of dark matter halos. This prescription was confirmed by the observation of Schechter (\\cite{schechter}) that the galaxy luminosity function can be described by the Schechter function: \\begin{equation} \\psi{\\rm (L)\\,dL}=\\psi^{*}\\left(\\frac{{\\rm L}}{{\\rm L}^{*}}\\right)^{\\alpha}\\exp({\\rm -L/L}^{*})\\,{\\rm d}\\left(\\frac{{\\rm L}}{{\\rm L}^{*}}\\right), \\end{equation} where the shape of the LF is described by a characteristic cut-off luminosity \\mbox{${\\rm L}^{*}$}, and $\\alpha$, the faint-end slope of the distribution. He also suggested that the shape of the galaxy LF is universal, and only the multiplicative constant $\\psi^{*}$ differs between clusters.  Numerous studies have since been made to determine the galaxy LF: some confirming the validity of the ``{\\em universal luminosity function hypothesis}'' (e.g. Lugger \\cite{lugger}; Colless \\cite{colless}; Trentham \\cite{trentham}); whereas others indicate that the LF is instead dependent on the environment of the galaxies, resulting in significant differences between the LFs from cluster to cluster and between cluster and field (e.g. L\\'{o}pez-Cruz \\etal \\cite{lopezcruz}; Valotto \\etal \\cite{valotto}). More recent studies based on large surveys empirically suggest that the galaxy LF is indeed dependent on environment, with dynamically-evolved, rich clusters and clusters with central dominant galaxies having brighter characteristic luminosities and shallower faint-end slopes, than poorer clusters, those with substructure (L\\'{o}pez-Cruz \\etal \\cite{lopezcruz}; De Propris \\etal \\cite{depropris}). In particular these differences can be explained through considering the composite LF as the sum of type-specific luminosity functions (TSLFs), each with its universal shape for a specific type of galaxies (e.g. elliptical, spiral, dwarf). The shape of the composite LF then can vary from cluster to cluster, and cluster to field, according to the mixture of different galaxy types in each environment resulting from the morphology-density relation (Dressler \\etal \\cite{dressler87}; Binggeli \\etal \\cite{binggeli}).  The galaxies most synonymous with the cluster environment are the bulge-dominated, passively-evolving galaxies which make up the cluster red sequence. Studies of low-redshift clusters  (e.g. Bower, Lucey \\& Ellis \\cite{bower}) indicate that, irrespective of the richness or morphology of the cluster, all clusters have red sequences, whose k-corrected slopes, scatters and colours are indistinguishable. Other studies find similar precise regularities between other physical properties of the early-type galaxies (surface-brightness, mass-to-light ratio, radius, velocity distribution, metallicity) resulting in the Faber-Jackson (Faber \\& Jackson \\cite{faber}) and Kormendy (\\cite{kormendy}) relations and the fundamental plane (Djorgovski \\& Davis \\cite{djorgovski}; Dressler \\etal \\cite{dressler87}). These all indicate that the early-type galaxies which make up the red sequences form a homogeneous population, not only within each cluster, but from cluster to cluster, and also that red sequences are universal and homogeneous features of galaxy clusters, at least at \\mbox{$z\\lsim0.2$}.  The photometric evolution of the cluster red sequence with redshift has been studied by a number of authors (e.g. Arag\\'{o}n-Salamanca \\etal \\cite{aragon}; Ellis \\etal \\cite{ellis}; Stanford \\etal \\cite{stanford98}; Kodama \\etal \\cite{kodama98}) for clusters out to \\mbox{$z\\simeq1.2$} indicating: that red sequence remains a universal feature of clusters; that the stellar populations of its constituent early-type galaxies are formed in a single, short burst at an early epoch \\mbox{$(z_{f}\\gsim2)$}; and that the galaxy colours have evolved passively ever since. This is confirmed by the spectra of red-sequence galaxies which in nearby clusters are best fit by simple stellar populations of ages \\mbox{9--12\\,Gyr}, resulting in their having characteristically strong spectral breaks at 4000\\,\\AA, and correspondingly red \\mbox{$U-V$} colours.  Cosmological models of structure formation indicate that the densest regions of the universe corresponding to the rarest overdensities in the primordial density field will have collapsed earliest, and will contain the most massive objects, i.e. rich galaxy clusters. Detailed cosmological simulations which follow the formation and evolution of galaxies (e.g. Kauffmann \\cite{kauffmann}; Blanton \\etal \\cite{blanton00}) predict that galaxies in these high-density regions are older and more luminous than those in typical density (i.e. field) regions. Having initially (\\mbox{$z\\sim5$}) been the most likely place for the formation of stars and galaxies, the high-density environment of rich cluster cores is later (\\mbox{$z\\sim$1--2}) filled with shock-heated virialised gas that does not easily cool and collapse (Blanton \\etal \\cite{blanton99}), inhibiting both the formation of stars and galaxies (Blanton \\etal \\cite{blanton00}), which instead occurs most efficiently in increasingly less massive dark matter halos. These models are able to successfully predict the observed cluster-centric star-formation and colour gradients, and the morphology-density relation of $z<0.5$ clusters (Balogh \\etal \\cite{balogh}; Diaferio \\etal \\cite{diaferio}; Springel \\etal \\cite{springel}; Okamoto \\& Nagashima \\cite{okamoto}), although it is not clear whether they predict trends with density of colour or star-formation for galaxies of a fixed morphology and luminosity. They show that these gradients in galaxy properties naturally arise in hierarchical models, because mixing is incomplete during cluster assembly, and the positions of galaxies within the clusters are correlated with the epoch at which they were accreted.  To examine the effect of environment on galaxy properties, in particular the galaxy luminosity function, and the cluster red sequence, we have performed a photometric study of the galaxy cluster A\\,209 at \\mbox{$z=0.21$} (Kristian, Sandage \\& Westphal 1978; Wilkinson \\& Oke \\cite{wilkinson}; Fetisova \\cite{fetisova}; Mercurio \\etal \\cite{paper1}) using archive wide-field $B$ and $R$-band imaging, which allows the photometric properties of the cluster galaxies to be followed out to radii of \\mbox{3--4$h_{70}^{-1}$\\,Mpc}. A\\,209 is a rich (richness class \\mbox{${\\rm R}=3$}; Abell \\etal \\cite{abell}) , X-ray luminous (\\mbox{${\\rm L}_{{\\rm X}}$(0.1--2.4\\,keV$)\\sim2.7\\times10^{45}\\,h^{-2}_{70}\\,$erg\\,s$^{-1}$}, Ebeling \\etal \\cite{ebeling}; \\mbox{${\\rm T}_{{\\rm X}}\\sim10$\\,keV}, Rizza \\etal \\cite{rizza}), and massive cluster (Mercurio \\etal \\cite{paper1}, but see also Dahle \\etal \\cite{dahle}). The cluster was initially chosen for its richness, allowing its internal velocity field and dynamical properties to be studied in great detail, and also for its known substructure, allowing the effect of cluster dynamics and evolution on the properties of its member galaxies to be examined. This substructure is manifested by an elongation and asymmetry in the X-ray emission with two main clumps (Rizza \\etal \\cite{rizza}), but no strong cooling flow is detected. Moreover, the young dynamical state is indicated by the possible presence of a radio halo (Giovannini, Tordi \\& Feretti \\cite{giovannini}), which has been suggested to be the result of a recent cluster merger, through the acceleration of relativistic particles by the merger shocks (Feretti \\etal \\cite{feretti}).  The internal dynamics of the cluster were studied by Mercurio \\etal (\\cite{paper1}) through a spectroscopic survey of 112 cluster members. A high value of the line-of-sight velocity dispersion was found, with \\mbox{$\\sigma_{v}=1394^{+88}_{-99}$\\,km\\,s$^{-1}$}. Assuming dynamic equilibrium, this value of $\\sigma_{v}$ leads to a virial radius of \\mbox{${\\rm R}_{{\\rm vir}}\\sim2.5h_{70}^{-1}$\\,Mpc}, and a virial mass of \\mbox{${\\rm M(R}<{\\rm R}_{{\\rm vir}})=2.$3--3.1$\\times10^{15}h^{-1}_{70}{\\rm M}_{\\odot}$}.  Evidence in favour of the cluster undergoing a dynamical evolution is found in the form of a velocity gradient acting along a SE-NW axis, which is the same preferential direction found from the elongation in the spatial distribution of galaxies and X-ray flux, as well as that of the cD galaxy. There is also significant deviation of the velocity distribution from a Gaussian, with evidence for two secondary clumps at \\mbox{$z=0.199$} and \\mbox{$z=0.215$}, which appear spatially segregated from the main cluster. These all indicate that A\\,209 is undergoing strong dynamic evolution with the merging of two or more sub-clumps along the SE-NW direction. This dynamical situation is confirmed by the weak lensing analysis of Dahle \\etal (\\cite{dahle}) which shows the cluster to be highly elongated in the N-S direction, and two significant peaks, the largest coincident with the cluster centre and cD galaxy, and the second 6\\,arcmin to the north.  The cluster was studied previously using optical data covering a \\mbox{78\\,arcmin$^{2}$} (\\mbox{160\\,arcmin$^{2}$} in $V$) region, reaching \\mbox{$B=22.8$}, \\mbox{$V=22.5$} and \\mbox{$R=22.0$}, and galaxy luminosity functions determined for each passband (Mercurio \\etal \\cite{paper2}).  The photometric data are presented in Section 2. As this cluster has been found to be significantly elongated along the SE-NW direction, we examine the effect of the cluster environment by measuring galaxy properties as a function of local surface density rather than cluster-centric radius. We describe the method for determining the cluster environment in Sect. 3. The ability to successfully subtract the background field population is vital to study the global properties of cluster galaxies from photometric data alone, and we describe the approach to this problem in Sect. 4. The composite galaxy LF is presented in Sect. 5, and the properties of the red sequence galaxies described in Sect. 6. The effects of environment on the blue galaxy fraction and mean galaxy colours are presented in Sects. 7 and 8. Sect. 9 is dedicated to the summary and discussion of the results. In this work we assume \\mbox{H$_{0}$=70\\,km\\,s$^{-1}$\\,Mpc$^{-1}$}, \\mbox{$\\Omega_{m}=0.3$}, and \\mbox{$\\Omega_{\\Lambda}=0.7$}. In this cosmology, the lookback time to the cluster A\\,209 at \\mbox{$z=0.209$} is 2.5\\,Gyr. ",
        "conclusions": "We have examined the effect of cluster environment, as measured in terms of the local surface density of \\mbox{$R<23.0$} galaxies, on the global properties of the cluster galaxies, through their luminosity functions, colour-magnitude relations, and average colours. For this study we have considered three cluster environments, a high-density region sampling the cluster core, and intermediate- and low-density regions which sample the cluster periphery.  \\subsection{The Galaxy Luminosity Function}  The LFs for galaxies within the virialised region are found to be well described by single Schechter functions to M$_{R}$, M$_{B}\\sim-16$, although there is an indication of a dip at $R=2$0--20.5 (M$_{R}=-20$). Dips in the LF at such absolute magnitudes appear common for low-redshift clusters (see Mercurio \\etal \\cite{paper2}), and appear stronger for richer, early-type dominated clusters (Yagi \\etal \\cite{yagi}). %  The faint-end slope, $\\alpha$, shows a strong dependence on environment, becoming steeper at $>3\\sigma$ significance level from high- to low-density environments. We explain this trend as the combination of two related effects:  a manifestation of the morphology-density relation whereby the fraction of early-type galaxies which have shallow faint-end slopes increases with density, at the expense of late-type galaxies which have steep faint-end slopes; and a luminosity-segregation due to dwarf galaxies being cannibalised and disrupted by the cD galaxy and the ICM in the cluster core \\cite{lopezcruz}.  To separate the two effects, we consider the TSLF of galaxies belonging to the cluster red sequence, as these are predominately early-type galaxies, and so any trends with environment should be independent of the morphology-density relation. Some luminosity-segregation is observed, with a reduction in the fraction of dwarf galaxies in the high-density regions, as manifested by their shallower TSLF faint-end slopes. This effect is smaller than that for the overall LF, being significant only at the $1.7\\sigma$ level, indicating that both luminosity-segregation and the morphology-density relation drive the observed trends in the overall LF. Luminosity-segregation is predicted by simulations for early-type galaxies in clusters, with bright early-type galaxies much more concentrated than their faint counterparts which follow the distribution of the dark matter mass profile (Springel \\etal \\cite{springel}).  In a study of 45 low-redshift \\mbox{$(0.04<z<0.18)$ }clusters, L\\'{o}pez-Cruz \\etal (\\cite{lopezcruz}) find that for the seven that are rich, dynamically-evolved clusters, characterised by the presence of a cD galaxy, their composite LFs are all well described by single Schechter functions with shallow faint-ends \\mbox{($\\alpha\\approx$-1.0)}. In contrast other clusters, often poorer, but in particular those not containing a cD galaxy, require two Schechter functions to fit their LFs, including a steep faint-end slope to model the dwarf population (see also Parolin, Molinari \\& Chincarini \\cite{parolin}). We thus indicate that the shallow faint-end slope observed in the high-density region of A\\,209 is related to the presence of the central dominant galaxy. cD galaxies are regarded as physically different to elliptical galaxies, with a different formation history, and are the product of dynamic processes which take place during the formation of their host cluster. They are thought to be built up through galactic cannibalism, or the accumulation of tidal debris. L\\'{o}pez-Cruz \\etal (\\cite{lopezcruz}) propose that the flatness of the faint-end slope in clusters containing cD galaxies results from the disruption of a large fraction of dwarf galaxies during the early stages of cluster evolution, the stars and gas of which are cannibalised by the cD galaxy, and the remainder redistributed into the intracluster medium.  \\subsection{Environmental Effects on the Red Sequence}  As well as having an affect on the TSLF of red sequence galaxies, the cluster environment could have an affect on the colour or slope of the red sequence itself, through the mean ages or metallicities of the galaxies. To examine this possibility the red sequence was fitted for the high- and intermediate-density regions independently. The red sequence was found to be \\mbox{$0.022\\pm0.014$\\,mag} redder in the high-density region than for the intermediate-density region by fixing the slope. In contrast no correlation between the slope of the red sequence and environment was observed. A similar effect is observed for bulge-dominated galaxies taken from the Sloan Digital Sky Survey (SDSS) (Hogg et al. 2003), in which the modal $^{0.1}[g-r]$ residual colour to the best-fitting C-M relation is \\mbox{0.01--0.02\\,mag} redder in the highest-density environments (corresponding to cluster cores) than in their low-density counterparts. In a study of 11 X-ray luminous clusters at $0.07<z<0.16$ Pimbblet \\etal (\\cite{pimbblet}) examine the red sequence as a function of cluster-centric radius and local galaxy density, and observe that the relation becomes progressively bluer as cluster-centric radius is increased out to 3\\,Mpc, and as the local surface density is decreased at a rate of $d(B-R)/d\\log_{10}(\\Sigma)=-0.08\\pm0.01$. In a study of the cluster A\\,2390 at $z=0.23$ Abraham \\etal (\\cite{abraham}) also observe the normalised colour of red sequence galaxies, $(g-r)_{r=19}$, to become increasingly blue with increasing cluster-centric radius, $r_{p}$, as \\mbox{$(g-r)_{r=19}=1.05-0.079\\log r_{p}$}.  These results indicate that the environment can affect the colour of a red sequence galaxy, so that red sequence galaxies are older and/or have higher metallicities in denser environments. To quantify this effect we consider a model red sequence galaxy as a 7.5\\,Gyr old  (at \\mbox{$z=0.209$}, corresponding to 10\\,Gyr at the present epoch), \\mbox{$\\tau=0.01$\\,Gyr}, solar-metallicity stellar population (Bruzual \\& Charlot \\cite{bruzual}). By varying its age and metallicity independently, we find that to reproduce the observed reddening in the C-M relation, galaxies in the high-density region must be on average 500\\,Myr older or 20\\% more metal-rich than their intermediate-density counterparts. By studying the spectra of 22\\,000 luminous, red, bulge-dominated galaxies from the SDSS, Eisenstein \\etal (\\cite{eisenstein}) indicate that red sequence galaxies in high-density regions are marginally older and more metal-rich than their counterparts in low-density environments, with both effects coming in at the same level.  This result is understandable in terms of cosmological models of structure formation, in which galaxies form earliest in the highest-density regions corresponding to the cores of rich clusters. Not only do the galaxies form earliest here, but the bulk of their star-formation is complete by $z\\sim1$, by which point the cluster core is filled by shock-heated virialised gas which does not easily cool or collapse, suppressing the further formation of stars and galaxies (Blanton \\etal \\cite{blanton99}; \\cite{blanton00}). Diaferio \\etal (\\cite{diaferio}) shows that as mixing of the galaxy population is incomplete during cluster assembly, the positions of galaxies within the cluster are correlated with the epoch at which they were accreted. Hence galaxies in the cluster periphery are accreted later, and so have their star-formation suppressed later, resulting in younger mean stellar populations. It should be considered however that this is a small-scale affect, and that this result in fact confirms that the red sequence galaxy population is remarkably homogeneous across all environments.  \\subsection{Galaxy Colours}  As well as considering the effect of the cluster environment on galaxy morphologies and LFs, its effect on star-formation has been examined by numerous authors (e.g. Balogh \\etal \\cite{balogh}; Ellingson \\etal \\cite{ellingson}; Lewis \\etal \\cite{lewis}). They find that star-formation is consistently suppressed relative to field levels for cluster galaxies as far as twice the virial radius from the cluster centre, and that for the majority of galaxies in the cluster core star-formation is virtually zero. For photometric studies it is possible to qualitatively measure the effect of the cluster environment on star-formation through measurement of the blue galaxy fraction --- the fraction of luminous \\mbox{($M_{R}^{*}<-20$)} galaxies whose colours indicate they are undergoing star-formation typical of late-type galaxies. We find that the blue galaxy fraction decreases monotonically with density, in agreement with other studies (e.g. Abraham \\etal \\cite{abraham}; Kodama \\& Bower \\cite{kodama}).  The observed trends of steepening of the faint-end slope, faintening of the characteristic luminosity, and increasing blue galaxy fraction, from high- to low-density environments, are all manifestations of the well known morphology-density relation (Dressler \\cite{dressler80}; Dressler \\etal \\cite{dressler97}), where the fraction of early-type galaxies decreases smoothly and monotonically from the cluster core to the periphery, while the fraction of late-type galaxies increases in the same manner. The observed trends in the composite LF simply reflect this morphology-density relation: the galaxy population in the cluster core is dominated by early-type galaxies and so the composite LF resembles that of this type of galaxy, with a shallow faint-end slope and a bright characteristic luminosity; whereas in lower density regions the fraction of late-type galaxies increases, and so the composite LF increasingly resembles that of the late-type TSLF, with a steep faint-end slope and a fainter characteristic magnitude (Binggeli \\etal \\cite{binggeli}).  Finally, we examined the effect of the cluster environment on galaxies through measuring the mean colour of luminous (\\mbox{$R<21$}) cluster galaxies as a function of their spatial position, as shown in Fig.~\\ref{colour}. This shows more clearly than any other result, the complex effects of the cluster environment and dynamics on their constituent galaxies. A\\,209 appears a dynamically young cluster, with a significant elongation in the SE-NW direction, the result of a recent merger with smaller clumps. To measure the effect of the cluster dynamics on the galaxy population, their properties should be measured against the parameter most easily related back to the cluster dynamics, their spatial position. As the cluster is significantly elongated, it should be possible to distinguish between whether the local density or cluster-centric distance is more important in defining the properties of galaxies, and indeed it is clear that the mean galaxy colour correlates most with the local density rather than the distance of the galaxy from the cluster core. It should be considered through that as the system appears elongated due to being the merger of two or more clumps, the notion of a cluster-centric radius loses much of its validity, as where does the centre of the system lie during a cluster merger? The location of the main cluster and the secondary merging clump appear confirmed by Fig.~\\ref{colour}, with the reddest galaxies concentrated around the cD galaxy (main cluster) and a more diffuse region 5\\,arcmin to the north coincident with the structure predicted from weak lensing analysis (Dahle \\etal \\cite{dahle}). The effect of the preferential SE-NW direction for A\\,209 is apparent in the presence of bright blue galaxies near the cD galaxy perpendicular to the axis and hence unaffected by the cluster merger, and an extension of red galaxies to the SE which may indicate the infall of galaxies into the cluster along a filament. This preferential SE-NW direction appears related to the large-scale structure in which A209 is embedded, with two rich (Abell class R=3) clusters A\\,222 at \\mbox{$z=0.211$} and A\\,223 at \\mbox{$z=0.2070$} are located \\mbox{$1.5^{\\circ}$} (15\\,Mpc) to the NW along this preferential axis.  Cluster dynamics and large-scale structure clearly have a strong influence on galaxy evolution, and it would be interesting to search for direct evidence of their effect on the star-formation histories of galaxies through spectroscopic observations of galaxies in the secondary clump, in the form of post-starburst signatures."
    },
    "0403/astro-ph0403626_arXiv.txt": {
        "abstract": "{We study the Large Magellanic Cloud Papillon Nebula (N\\,159-5), a conspicuous High Excitation Blob (HEB) lying in the star forming complex N\\,159. Using {\\it JHK} near-infrared photometry obtained at the ESO VLT with the ISAAC camera, we examine the stellar populations associated with the Papillon, tracing their history using stellar evolution models. Two populations are revealed: one composed of young, massive stars with an age \\ab\\, 3 Myr, and a second consisting of older lower mass stars of age spreading between 1 and 10 Gyr.  We analyze the properties of those populations and discuss their significance in the context of N\\,159.  We also estimate that if the star at the center of the Papillon is single its initial mass is \\ab\\,50\\,\\sm\\, and it is affected by an extinction $A_{V}$\\,\\ab\\,7 mag. \\\\ ",
        "introduction": "The compact \\h2 regions called High-Excitation Blobs (HEB) constitute a rare class of ionized nebulae in the Magellanic Clouds.  They are characterized by high excitation, small size, high density, and large extinction compared to typical Magellanic Cloud \\h2 regions. These objects are tightly linked to the early stages of massive star formation, when the stars begin to hatch from their parental molecular clouds. For this reason their study yields valuable information for a better understanding of the formation of massive stars.  \\\\  The \\object{N\\,159} complex \\citep{henize} lies south of the famous starburst site 30 Dor and has attracted special attention over the years. Its other designations are \\object{MC\\,77} \\citep{mcgee72}, \\object{LH\\,105} \\citep{lh}, and \\object{DEM\\,271} \\citep{dem}.  It shows several signs of ongoing star formation activity, such as infrared sources, cocoon stars, masers, and is also associated with the most important concentration of molecular gas in the LMC \\citep{jones86,brooks97, johansson98}. The molecular emission is in fact due to three distinct giant molecular clouds, known as N159-East, N159-West, and N159-South.  Molecular lines tracing high density regions are observed towards N\\,159-W and N\\,159-S in CS, CN, HCN, and HCO$^{+}$ \\citep{heikkila99}, while $^{13}$CO and upper-level $^{12}$CO transitions, and the [\\cii] emission line were mapped towards the three giant molecular clouds, including N\\,159-E \\citep{bolatto00}. The region we are interested in lies near N\\,159-E. \\\\  The Papillon Nebula, also called \\object{N\\,159-5}, to which the present study is devoted, is  the prototype of the HEB family \\citep{mhm82}, which now possesses several members, such as N\\,160A1, N\\,160A2, N\\,83B, N\\,11A in the LMC and N\\,88A and N\\,81 in the SMC. Recent {\\it HST} observations have resolved most of these objects and revealed turbulent media typical of newborn massive star regions marked by strong stellar winds interacting with the ambient ionized gas \\citep{mhm99a, mhm99b, mhm99c, mhm01a, mhm01b, mhm02a, mhm02b}. These observations also showed a large extinction due to local dust associated with ionized gas.  In a number of cases the exciting sources were uncovered as a small cluster of massive stars.\\\\  In particular, the WFPC2 {\\it HST} observations revealed that the featureless blob N\\,159-5 has in fact the morphology of a {\\it ``papillon''}, i.e. it is a butterfly-shaped ionized nebula with the ``wings'' separated by \\ab\\,2\\frac.3 (0.6 pc) \\citep[][ hereafter paper I]{mhm99c}.  Moreover, two subarcsecond features resembling a ``smoke ring'' and a ``globule'' were detected in the wings.  The images also showed a large number of subarcsecond filaments, arcs, ridges, and fronts carved in the ionized gas by the stellar winds from massive stars in the N\\,159 complex. However, no bright stars associated with the core of N\\,159-5 could be identified in the {\\it HST} images.  Could this be due to the extinction by dust large enough to hide the exciting stars? Since an $A_V \\geq 6$\\,mag was needed to bring an O8 star below our {\\it HST} sensitivity limit, we decided to perform high resolution near-IR imaging of the region with the ESO Very Large Telescope in order to address this issue and explore the properties of the stellar population of N\\,159.  \\\\ ",
        "conclusions": "The population of young massive stars, which was discussed earlier (Sect. 3.2), is spatially distributed over the whole field, while a sample of it,  represented by triangles in Figs.\\,\\ref{col-mag} and \\ref{col-col}, is  grouped in a cluster, marked as region B. This grouping is expected given the young age of these stars. How can though one explain the separation between this cluster and the other massive stars which even if they have similar ages are at a considerable distances from the cluster, for example \\ab\\,70\\frac\\, (\\ab\\,18\\,pc) for star \\#193, one of the most distant?  One explanation could be that massive star formation may have taken place simultaneously at different parts of the molecular cloud.  At these locations the molecular material has been fully dissociated and ionized, and we do not observe it now.  Alternatively, massive star formation   has  occurred in cluster B, and subsequently a number of the members have been ejected due to the dynamics of the cluster.\\\\  It has been shown that once an embedded cluster forms, three mass evacuation mechanisms work over different timescales to disrupt it \\citep{kroupa01}: $a)$ expulsion of embryonic gas (approximately during the first 0--5 Myr), $b)$ mass loss from evolving stars (significant after \\ab\\,3 Myr), and $c)$ stellar dynamical evaporation and ejection of stars (all times).  Binary-binary collisions are required to produce high velocity escapees to occur in low density clusters \\citep{leonard88,leonard90}, although simple calculations suggest that such interactions are rather unlikely.  Recently \\citet{vine03} have studied the dynamics of massive stars in young clusters containing gas and stars. They have shown that the location of massive stars outside the core of the cluster does not exclude their formation in the dense cluster core. The massive stars could have originated in the core, but escaped from that region during the gas expulsion phase.  Furthermore, the ejection of the OB stars must have happened during an earlier evolutionary stage when the cluster was most probably more compact than today \\citep{portegies99}. Assuming that star \\#193 has been kicked out of cluster B, an escape velocity of \\ab\\,5.5\\,km\\,s$^{-1}$ has been necessary for it to reach its observed position after a travel time of 3 Myr. This estimate is a lower limit due to projection on the sky of a three-dimensional configuration in space.  Higher velocities are quite possible since escapees can leave their birthplace with velocities up to 100 km\\,s$^{-1}$ or even larger \\citep{leonard90,kroupa95}.  \\\\  We note also that all the candidate massive stars are devoid of proper nebulosity, in contrast to the Papillon. This fact suggests that the Papillon is probably the youngest visible massive star formation event in the whole field. The strength of the molecular hydrogen emission detected towards the Papillon confirms its nature as a very young star formation region \\citep{israel88,kawara88}. In fact the observed luminosity of the H$_{2}$ line {\\it v = 1 -- 0 S(1)} towards the Papillon is two times larger than that observed at the Orion source \\citep{kawara88}. The massive star(s) powering the Papillon have not had enough time to disrupt the \\h2 region.  Moreover, the presence of nebulosity excludes the possibility for the Papillon of  ejection from cluster B. It is therefore conceivable that the Papillon lies somewhat above or below the mean plane of N\\,159. We believe that the Papillon is situated at the side nearer to us since it is visible in the optical. We estimate that star \\#371, which lies towards the center of the Papillon, has a mass of \\ab\\,50\\,M$_\\odot$, even though based on our current resolution we cannot exclude the possibility that the star is multiple.  Should other massive stars be embedded inside that nebula, much better spatial resolution and deeper exposures are required in order to uncover them. From our previous {\\it HST} observations we estimated an exciting star of type at least O8\\,V, \\ab\\,30 \\sm\\, \\citep{vacca96}, for the Papillon using the \\hb\\, flux measurement and assuming that the \\h2 region is ionization bounded (Paper I). The difference between the two mass estimates is due to the fact the \\h2 region is density-bounded at least towards us and that the flux correction for extinction is not straightforward. The latter point is probably the reason why the radio continuum observations, which are less affected by extinction, yield a higher Lyman continuum flux.  In an earlier work \\citep{mhm85}, we used the radio continuum observations at 843 MHz, obtained with a beam of 43\\frac\\,\\x\\,46\\frac, to derive a flux density of 55 mJy for N\\,159-5, after correcting for contamination by the surrounding field.  A resulting ultraviolet flux of \\ab\\,1.2\\,\\x\\,10$^{49}$ photons indicates an O7\\,V type star of \\ab\\,38\\,\\sm\\, \\citep{vacca96}.  Given the uncertainties involved, the stellar mass estimates based on the \\hi\\, emission from the nebula agree well with the \\ab\\,50 \\sm\\, derived from photometry using evolutionary models.  \\\\  An age of \\ab\\,3 Myr was derived for the massive star population using the evolutionary models and supplementary data on one of the members. We wish to note though that this age estimate may not be very accurate due to the degeneracy of the near IR colors of massive stars. In fact any isochrone between 1 and 10 Myrs would be consistent with our data. We favored the 3 Myr isochrone in order to meet the requirements of star Sk$-69^{\\circ}257$. \\\\   One can also estimate the number of stars which power the \\h2 region N\\,159 on the basis of radio continuum observations. \\citet{clarke76} measured a radio continuum flux density of 6.5 Jy at 408 MHz using the Molonglo telescope whose beam had a width of 2\\min.6\\,\\x\\,2\\min.9. The beamwidth is comparable with the size of our ISAAC field, and the target coordinates match  the position of the Papillon, while the reported pointing accuracy is 18\\frac\\,\\x\\,5\\frac. The derived Lyman continuum flux of 1.36\\,\\x\\,10$^{51}$ photons s$^{-1}$ corresponds to some 40 massive stars of type O5\\,V with an initial mass of \\ab\\,50\\,\\sm\\,\\citep{vacca96}. Taking a Salpeter-like initial mass function with slope $ x = - 1.5$, we can predict the presence of some  360 stars of mass about 10\\,\\sm\\, and  3240 stars of \\ab\\,2\\,\\sm. Where are these 40 O5\\,V stars? It is quite possible that they are among the stars we imaged but due to the degeneracies in the colors mentioned earlier they can only be clearly identified if spectroscopic observations were available. Moreover, some of them may be embedded in the molecular cloud and some situated outside our ISAAC field.\\\\     The color-mag diagram (Fig.\\,\\ref{col-mag}) also shows the presence of intermediate mass stars of \\ab\\,4--10\\,\\sm\\, on the main sequence formed together with high mass stars \\ab\\,3 Myr ago. This is in agreement with more detailed results on the Orion Nebular Cluster (ONC) based on a large body of data  (\\ab\\,3500 stars identified within \\ab\\,2.5 pc of the Trapezium, among which at least \\ab\\,1600 with photometric and spectroscopic data in the visible) \\citep{hillenbrand97}.  According to these studies, low-, intermediate-, and high-mass stars have formed together in the ONC a few Myr ago \\citep{palla99}. However, this may not be a universal trend since \\citet{herbst82}'s study of NGC\\,3293 led them to the conclusion that in a cluster low- and intermediate-mass stars form first, with the process continuing gradually until the high-mass stars appear. This result is in agreement with more recent findings on star formation in LMC clusters and associations. For instance, in the case of the R\\,136 cluster, situated in the LMC 30 Dor, \\citet{massey98} arrived to  the conclusion that intermediate-mass stars began forming some 6 Myr ago and continued up to the time when the high-mass stars formed, 1--2 Myr ago.  \\\\   An interesting question is  whether the young (\\ab\\,3 Myr) and old (\\ab\\,1--10 Gy) stellar populations have formed in the same region of space.  Although presently we do not have the necessary data to address this issue and cannot reach a firm conclusion, it is quite possible that both populations be spatially unrelated. The LMC is known to have a considerable depth, the old population can have formed in a different depth during much earlier star formation activities.  In order to get some insight about this question, we used the 2MASS data to probe a bare stellar field devoid of any particular nebular emission lying near the N\\,159 complex (radius 1\\min.22, centered on  $\\alpha$ = 05h\\,39m\\,00s, $\\delta = -69^{\\circ}$\\,47\\min\\,30\\frac\\,).  The corresponding HR diagram shows the absence of a young, unreddened population, but the presence of an old population resembling the one found towards N\\,159. Although this population is relatively smaller in number with respect to that of N\\,159, since 2MASS is not as deep as  our photometry, the presence of the old population is certain. The old population seems therefore to be a common background stellar component towards this part of the LMC.\\\\   The presence of low-mass pre-main sequence LMC stars in the above diagrams seems unlikely, even if those objects are characterized by large near IR colors, {\\it H -- K}\\,\\ab\\,1.5 mag \\citep{lada92,chabrier00}. A pre-main sequence star of \\ab\\,1 \\sm\\, has a luminosity of  log {\\it (L/$L_{\\odot})$}\\,\\ab\\,1 on its birthline, corresponding to an observed visual magnitude of \\ab\\,21, which is below our detection limit. An intermediate mass pre-main sequence star of 5\\,\\sm\\, has an effective temperature of \\ab\\,11,000 K and log {\\it (L/$L_{\\odot})$}\\,\\ab\\,3, corresponding to a reddened 16 magnitude star, occupying loci around {\\it J -- H}\\,\\ab\\,0.5 and {\\it H -- K}\\,=\\,0.5 mag \\citep{lada92}. There is  a few  number of sources with such colors in Fig.\\,\\ref{col-col}, given the color uncertainties at those magnitudes. Therefore, we cannot exclude the possibility that some of those points represent intermediate mass PMS stars. As for more massive objects, the concept of pre-main sequence is not applicable to stars above \\ab\\,6\\,\\sm\\, since the birthline and the ZAMS unify at those mass levels \\citep{palla93}. \\\\   Comparison between LMC N\\,159 and SMC N\\,81 points out dramatic differences between the environments of these two HEBs. The present work shows the Papillon as part of a rich complex containing a large molecular cloud and a cluster of young, massive stars, whereas our previous study of SMC N\\,81, based on ISAAC near IR observations \\citep{mhm03}, revealed a solitary star formation event.  Moreover, since the two compact \\h2 regions have several comparable characteristics,  if we assume  that they have gone through a similar formation process, then the HEB formation can take place in both very dense as well as  rather sparse environments. \\\\"
    },
    "0403/nucl-th0403062_arXiv.txt": {
        "abstract": "We perform molecular dynamics simulations of screening by bound target electrons in low energy nuclear reactions. Quantum effects corresponding to the Pauli and Heisenberg principle are enforced by constraints. We show that the enhancement of the average cross section and of its variance is due to the perturbations induced by the electrons. This gives a correlation between the maximum amplitudes of the inter-nuclear oscillational motion and the enhancement factor. It suggests that the chaotic behavior of the electronic motion affects the magnitude of the enhancement factor. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403632_arXiv.txt": {
        "abstract": "We present results from X-ray and radio observations of the recently discovered young Vela-like pulsar \\psr , which is coincident with the \\EG\\ \\gr\\ source GeV 2020+3658.  A 19.0-ks {\\it Chandra} ACIS-S observation has revealed a $\\sim$20$^{\\prime \\prime} \\times 10^{\\prime \\prime}$ pulsar wind nebula that is reminiscent of the equatorial tori seen around some young pulsars, along with thermal emission from an embedded point source ($kT_{\\infty} = 0.15 \\pm 0.02$\\,keV).  We name the nebula \\pwn .  Its spectrum is well fit by an absorbed power-law model with photon index $\\Gamma = \\errors{1.7}{0.3}{0.2}$, hydrogen column density n$_{\\rm H} = (\\errors{7.8}{1.7}{1.4}) \\times 10^{21} $\\,cm$^{-2}$, and an unabsorbed 0.3$-$10.0\\,keV flux of $(\\errors{1.9}{0.1}{0.3}) \\times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$.  We have spatially fit \\pwn\\ with a model that assumes a toroidal morphology, and from this we infer that the torus is highly inclined ($83^{\\circ} \\pm 1^{\\circ}$) to the line of sight.  A 20.8-ks {\\it Chandra} observation in continuous-clocking mode reveals a possible pulse detection, with a pulsed fraction of $\\sim$37\\% and an H-test probability of occuring by chance of $1.2 \\times 10^{-4}$.  Timing observations with the Arecibo radio telescope spanning two years show that \\psr\\ glitched sometime between MJDs~52616 and 52645 with parameters ${\\Delta \\nu}/{\\nu} = (2.587 \\pm 0.002) \\times 10^{-6}$ and ${\\Delta \\dot{\\nu}}/{\\dot{\\nu}} = (6.2 \\pm 0.3) \\times 10^{-3}$, similar to those of the largest glitches observed in the Vela pulsar.  \\psr\\ is heavily scattered (${\\tau}_{sc} = 17.7 \\pm 0.9$\\,ms at 1\\,GHz) and exhibits a significant amount of timing noise. ",
        "introduction": "Pulsar wind nebulae (PWNe) allow us to study the basic energetic processes of young, rotation-powered pulsars.  It is still unclear how the rotational kinetic energy of these objects is converted into a relativistic particle outflow.  The {\\it Chandra X-ray Observatory}'s unprecedented spatial resolution has revealed that the winds from young, energetic pulsars can lead to highly structured PWNe.  In some cases these PWNe take the form of an equatorial outflow with polar jets, where the best example of this is the Crab nebula \\citep{wht+00}.  These torus plus jet nebulae, of which there are roughly a half-dozen examples, are particularly interesting because a) they allow us to study the interaction of the PWN with its surrounding environment (the location of the bright torus is generally believed to mark the inner termination shock of the pulsar wind) and b) the well-defined morphology of the PWN can constrain the geometry of the pulsar.  This is crucial for testing models of high-energy pulsations, such as the outer-gap model \\citep[e.g.][]{rom96} and the polar-cap model \\citep[e.g.][]{dh96}.  \\psr\\ (spin period $P \\sim$103.7\\,ms, dispersion measure DM $\\sim$369\\,pc cm$^{-3}$, and flux density at 1400\\,MHz $S_{1400} \\sim$100\\,$\\mu$Jy) was discovered by \\citet{rhr+02} with the 305-m Arecibo radio telescope during a targeted search for radio pulsations from X-ray sources proposed as counterparts to unidentified \\EG\\ \\gr\\ sources \\citep*{rrk01}.  Timing observations of \\psr\\ revealed that although it is dim, it is young and energetic (characteristic age ${\\tau}_c \\equiv P/2\\dot{P} = 17$\\,kyr, $\\dot{E} \\equiv 4{\\pi}^2 I \\dot{P}/P^3 = 3.6 \\times 10^{36}$\\,ergs s$^{-1}$, assuming moment of inertia $I = 10^{45}$ g cm$^{-2}$), and a likely counterpart to the {\\it ASCA} X-ray source AX J2021.1+3651 and the \\EG\\ \\gr\\ source GeV J2020+3658.  This discovery bolsters the idea that many of the Galactic unidentified \\EG\\ sources \\citep[see][]{hbb+99} are young energetic pulsars that were not detected in large-scale radio pulsar surveys.  While the spin properties of \\psr\\ are very similar to those of the known \\gr\\ pulsars Vela (PSR~B0833$-$45, $P = 89.3$\\,ms, $\\dot{E} = 6.9 \\times 10^{36}$\\,ergs s$^{-1}$) and PSR~B1706$-$44 ($P = 102.5$\\,ms, $\\dot{E} = 3.4 \\times 10^{36}$\\,ergs s$^{-1}$), the 12\\,kpc distance inferred from its DM (which places it on the far edge of the outer spiral arm) using the NE2001 electron model of the Galaxy \\citep{cl02} would make it much more \\gr\\ efficient than either Vela or PSR~B1706$-$44.  In fact, it would have to be even more \\gr\\ efficient than the much-less energetic \\gr\\ pulsars Geminga and PSR~B1055$-$52 (assuming the NE2001 DM distance to PSR~B1055$-$52 of 0.72\\,kpc).  This is contrary to the observation that \\gr\\ efficiency ${\\eta}_{\\gamma} \\equiv L_{\\gamma}/\\dot{E}$ increases with age \\citep{thom99} and the theoretical expectation that it is inversely proportional to the square-root of the spin-down luminosity \\citep[${\\eta}_{\\gamma} \\propto {\\dot{E}}^{-1/2}$, e.g.][]{zh00a}.  However, since DM distances can be off by a factor of a few, \\psr\\ may be significantly closer (or farther) than its DM implies.  The similarity of \\psr\\ to the Vela pulsar, which is surrounded by a highly-structured pulsar wind nebula \\citep*[see][]{hgh01}, prompted us to make X-ray observations of the area surrounding \\psr.  Only $\\sim$30 of the $>$1300 known pulsars in the Galaxy are Vela-like \\citep[i.e. with $P \\sim 100$\\,ms, $\\dot{E} \\sim 10^{36}-10^{37}$ ergs s$^{-1}$, and characteristic ages $\\tau_c \\sim$10$-$20\\,kyr, see][]{kbm+03}.  Half of these are known to have associated PWN \\citep*{krh04}.  Here we present the results of X-ray observations made with the {\\it Chandra} ACIS-S detector.  These observations resolve the {\\it ASCA} source into a new PWN, which we name \\pwn , and a thermally emitting point source, likely the neutron-star surface.  {\\it Chandra} observations in continuous-clocking mode show a potential detection of X-ray pulsations from \\psr .  We also present radio timing observations made with the Arecibo telescope, which show that \\psr\\ has glitched and is highly scattered. ",
        "conclusions": "We have used the {\\it Chandra X-ray Observatory} to observe the young and energetic pulsar \\psr\\ and have detected a new PWN, \\pwn , as well as an embedded point source near the radio timing position of the pulsar.  The morphology of \\pwn\\ is perhaps best decribed as an equatorial torus, and we have derived geometrical parameters for the orientation of the system by spatially fitting the nebula.  Spectral fitting of the nebula shows that it is similar to the nebulae of other energetically-similar pulsars.  Spectral fitting of the point source shows that it is well fit by an absorbed black-body model or a black-body plus power-law model.  Radio timing observations of \\psr\\ have revealed a glitch similar to the largest glitches seen in the Vela pulsar. Analysis of countinuous-clocking data from {\\it Chandra} indicates that there may be weak X-ray pulsations from the source (pulsed fraction of $\\sim$37\\%) at the period predicted by the radio ephemeris.  \\psr\\ remains the most likely counterpart to the high-energy \\EG\\ \\gr\\ source GeV J2020+3651 (which is the 10th brightest \\gr\\ source above 1\\,GeV).  It is an excellent \\gr\\ pulsar candidate for the AGILE and GLAST missions, although a contemporaneous ephemeris will be essential, given the large amount of timing noise the pulsar exhibits.  Deeper observations with {\\it Chandra} and {\\it XMM-Newton} are needed to clarify the morphology of \\pwn\\ and to confirm the presence of pulsations.  Such observations are also able to confirm whether the faint outer jet hinted at here is real.  Currently, there only about a half dozen PWNe with torus plus jet morphologies, with the Crab and (likely) Vela nebulae as the best examples.  With a deeper observation, \\pwn\\ may prove to be the third best example.   Further work by our group will also include analysis of radio polarization data we have taken with Arecibo.  These data may help bolster the geometrical interpretation described here if the swing of the polarization angle across the pulse is measurable.  Proper-motion measurements could check if \\psr\\ is moving along the axis of the nebula, which is the case for both Vela and the Crab.  However, measuring the proper motion of \\psr\\ (which will likely be small since the distance is likely large) will be very difficult both through radio timing because of timing noise or through VLBI because the radio source is very faint.  We are also analyzing VLA observations (in A, C, and D arrays) of the source region to look for a radio PWN and/or SNR.  This is important for characterizing the surrounding medium, which has an important effect on the morphology of the PWN and its size."
    },
    "0403/astro-ph0403318_arXiv.txt": {
        "abstract": " ",
        "introduction": "The planetary nebula (PN) phase is one of the final stages of stellar evolution for 1-8$\\msun$ stars. A planetary nebula is created when an asymptotic giant branch (AGB) star sheds its hydrogen-rich envelope via stellar wind, exposing the underlying hot carbon-oxygen (CO) core. The surface of the remnant star has a surface temperature T$>$30000K and therefore it emits a significant number of photons with wavelengths, ($\\lambda$), less than 912 Angstroms. These photons with energy greater than 13.6eV are sufficiently energetic to ionize hydrogen in the ejecta. The free electrons collisionally excite the ions in the nebula. The excited atoms then radiatively de-excite creating emission lines. PNe appear green in a small telescope or binoculars because of the [OIII]$\\lambda$5007.  The energy input from the star is balanced by the energy emitted by lines from the nebulae. Photons from the star which photoionize the nebulae are the primary energy source. Emission lines carry this energy away. This energy balance determines the temperature of the nebulae and the relative numbers of each ionization stage of each atomic species.  PN emission lines can be used to determine to reasonable accuracy the nebular abundances of several atomic species. The density in most PNe is too low for ions to be collisionally de-excited. Therefore, these ions must emit one or more photons to get back to the ground state. By balancing the number of collisional excitations with the number of radiative deexcitations, the abundance of some ionic species can be determined, e.g., using the [OIII]$\\lambda$5007 line we can get the abundance of O$^{+2}$. To get the nebular abundance of an atom, an ionization correction factor is used to account for the unseen ionization stages. Therefore, the uncertainties in each element vary depending on how much of the element is in each stage. The procedure for determining elemental abundances is described in more detail in elementary textbooks on the subject (e.g. Osterbrock 1989).  Recently, detailed photoionization models have become widely available, allowing detailed calculations of the nebulae (i.e. CLOUDY). Such models improve our ability to determine nebular abundances.  The abundances of several important elements can be inferred for PNe. The list includes hydrogen, helium, carbon, nitrogen, oxygen, neon, sulfur, and argon. Some of these such as neon and argon can not be measured AGB phase which preceeds the PN phase.  The gas that makes up a PN has been processed via stellar evolution. This processing leads to the difference between the abundances in PNe and the differences in solar or HII region abundances. The abundances in the Sun and in HII regions are respectively believed to reflect the composition of the interstellar medium (ISM) at 5Gyrs and the current epoch, respectively. The material making up PNe has been processed by stellar evolution. This probably leads to enhancements in the abundances of He, C, and N.   The globular cluster M15 contains the well studied planetary nebula (PN) K648. This is one of the few galactic PNe with a reasonably well-determined distance. Therefore, fundamental properties such as the stellar luminosity can be determined with some confidence. Because of its globular cluster membership, many of the progenitor properties, such as the zero age main sequence (ZAMS) mass, can be inferred reliably.  Due to the importance of K648 as a halo PN, it has been the focus of several abundance studies, and all of these show the abundances of most metals to be depleted relative to the sun, consistent with a progenitor of low metallicity.  Carbon is an exception; studies which determine the ratio (by number) of C/O in K648 infer values that range from $4-11$ (Adams \\etal\\ 1984; Henry, Kwitter, and Howard 1996; Howard, Henry, and McCartney 1997), which is far above C/O in the Sun of 0.43 (\\nocite{ag89}Anders and Grevesse 1989, hereafter AG89). This is in fact much higher than the average C/O ratio of $\\approx 0.8$ for solar neighborhood PNe (\\nocite{rs94}Rola and Stasi{\\'n}ska 1994).  Low and intermediate mass stars that have left the main sequence, ascended the giant branch, and passed through the horizontal branch, then enter a thermally unstable phase where energy is generated by shell He and H-burning called the thermally pulsing asymtoptic giant branch (TP-AGB) stage, which is a very important yet not well understood phase [Detailed reviews of this stage can be found in \\nocite{i95} Iben (1995), \\nocite{l93} Lattanzio (1993), and \\nocite{ir83} Iben and Renzini (1983)]. During the TP-AGB stage the star alternates between a long stage where the luminosity is generated mostly by quiescent hydrogen shell burning, with a helium burning layer producing a minority of the energy, and a thermal runaway stage in the unstable helium burning layer (\\nocite{sh65} Schwarzschild and H\\\"{a}rm 1965, 1967; and \\nocite{w66} Weigert 1966).  The second stage results in expansion of the outer layers and an extinguishing of the H burning shell. This short stage, characterized by rapid changes, with helium burning dominating the energy generation, is known as a thermal pulse or a He shell flash.  TP-AGB stars exhibit large mass-loss rates ranging from 10$^{-7}-$10$^{-4}$ \\msolyr\\ . Indeed such high mass-loss rates are predicted to result in the ejection of the envelope, at which point the star leaves the AGB and becomes a planetary nebula central star (CSPN). The first models of CSPN tracks were made by \\nocite{p71} Paczynski (1971) who showed that the CSPNs evolve horizontally on the HR diagram when nuclear burning is still taking place and then as they cool the luminosity and temperature decrease. \\nocite{hs75} H\\\"arm and Schwarzschild (1975) showed that a CSPN could leave the AGB as either a helium burning or hydrogen burning star. The observational consequences of hydrogen and helium burning have been studied in the more refined models including mass loss showed that the subsequent evolution of the central star depends on whether or not the star leaves the AGB as a helium or hydrogen burner [Sch\\\"{o}nberner (1981, 1983) and \\nocite{i84} Iben (1984)].  Low mass stars (M$\\lesssim$ 3 M$_{\\odot}$) can experience two mixing episodes or ``dredge-ups''. During dredge-up, material that has been processed by nuclear burning is mixed into the surface layers. At the entrance to the giant branch, the convective region can extend into the core, leading to mixing of CNO products into the outer layers. Similarly as shown by \\nocite{i75} Iben (1975), after a thermal pulse on the AGB, the convective region can extend into the core, mixing He-burning products into the outer layers. These two mixing events are known as first and third dredge up, respectively (second dredge up will not concern us here). Therefore, a third dredge-up is a natural explanation of the high carbon abundance found in K648.  On the other hand, no carbon stars have been observed either in M15 or in any other globular cluster, although such stars should be the immediate progenitors of objects such as K648 if a third dredge-up occurs. Thus, the lack of carbon stars in M15 weakens the argument for a third dredge-up event.  One possible explanation for the absence of carbon stars is a delayed scenario in which the third dredge-up of carbon rich material changes the structure of the envelope during the following interpulse phase, ultimately increasing the mass-loss rate significantly and driving off the stellar envelope (Iben 1995). Thus envelope ejection is delayed until the interpulse phase following this dredge-up of carbon rich material.  Another explanation supposes that the envelope is removed during the quiescent He-burning stage that follows a thermal pulse (\\nocite{ren89}Renzini 1989 and \\nocite{rfp88}Renzini and Fuci-Pecci 1988).  The carbon then originates in a fast wind from the central star (CSPN).  In addition, the wind produces shock-heating in the nebula, which, if not properly accounted for during an abundance analysis, may lead to the inference of a spuriously high C/O ratio. In this case the envelope would be ejected immediately after a thermal pulse while helium shell burning still dominates the luminosity.  We refer to this mechanism as the prompt scenario.  In this paper we calculate detailed envelope models of thermally pulsing asymptotic giant branch star envelopes to test the predictions of the delayed mechanism, perform other calculations relevant to the prompt mechanism, and compare output of each with observations of K648. Section~2 describes the envelope code, section~3 presents the observational data and the results for the delayed and prompt models, and a brief discussion of our findings is given in section~4. ",
        "conclusions": "A multivariate data analysis of galactic planetary nebulae has been carried out with the goal of determining the effective dimensionality of PN parameter space as well as the importance of the numerous observational parameters in classifying PNe.  The parameters employed for this study are the abundance ratios He/H, N/H, O/H, Ne/H, N/O, Ne/O, and the spatial and kinematic parameters R, Z, and $V_{LSR}$. Our main results are:  \\begin{enumerate} \\item Planetary nebula abundance parameter space is two dimensional, i.e. the data can be described by two principal components.  The first and most important one is related to the products of stellar nucleosynthesis during the evolution of PN progenitors.  The second component is related to progenitor metallicity.  \\item The kinematic properties are not linearly correlated with any other parameters.  \\item The disk PNe separate into three clusters, each having a distinct location in the nucleosynthesis-metallicity plane.  A fourth cluster comprises many of the bulge PNe. The bulge PNe only separate from the disk PNe because of the inclusion of the parameters R and $V_{LSR}$.  \\item We have identified some unusual type I PNe with extremely low oxygen abundances. However, it is not clear if these are real.  \\item Bulge and disk PNe show the same distributions along the nucleosynthesis and metallicity components.  \\end{enumerate}  This analysis is a first attempt to characterize planetary nebula parameter space in an objective way. The future addition of parameters such as carbon abundance, nebular expansion velocity, morphology, binarity, and central star mass further improve our understanding of these objects.    \\newpage \\chapter{Thermally Pulsing AGB Models} \\label{modelchap}  Thermally pulsing asymptotic giant branch (TP-AGB) stars experience two distinct and repeating phases: the helium shell flash or thermal pulse (TP) and the time between pulses or the interpulse phase (IP). The duration of the thermal pulse ($\\tp$) is short compared to the duration of the interpulse phase ($\\tip$) with $\\tp /\\tip\\approx 0.01$. During the TP-AGB, these stars spend most of the time in the interpulse phase, burning both He and H burn quiescently in thin shells. During the interpulse phase, most of the luminosity is produced by the hydrogen burning shell. However, the helium shell in this configuration is unstable, and it will eventually result in a helium shell flash. Thermal pulses occur quasi-periodically, resulting in the luminosity of the helium burning shell increasing by several orders of magnitude. The increased luminosity causes the outer layers of the star to expand, including the hydrogen burning shell. The luminosity generated by hydrogen burning during the shell flash decreases to essentially zero due to the expansion. After the thermal pulse the star resumes its interpulse configuration. This cycle repeats until the envelope is lost via mass-loss.  The two distinctive phases of the TP-AGB have important consequences for the surface abundances of the star. At the end of a thermal pulse it is possible for the convective envelope to penetrate into regions where He has been partially burned into $\\ctw$ and minor amounts of $\\oxy$, mixing these products to the stellar surface. This process is called the third dredge-up and is believed to be the process by which carbon stars are produced. In the more massive TP-AGB stars it is possible for the base of the convective envelope to get hot enough to burn $\\ctw$ to $\\cth$ and $\\nit$ and $\\oxy$ to $\\nit$ in a process known as hot-bottom burning. This process is believed to produce the nitrogen rich PNe. Hot bottom burning can take the carbon dredged-up into the envelope and process it into nitrogen. This is one of the possible solutions to the so-called carbon star mystery.  The third dredge-up occurs at the end of a thermal pulse and mixes helium and carbon rich material into the stars outer layers. During the thermal pulse a convective shell appears in the helium burning zone. This convection zone spans the mass range from the base of the interpulse He burning shell to just below the position of the convective envelope during the preceding interpulse phase. In this zone these important helium burning reactions take place: \\begin{itemize} \\item{$\\he +\\he\\longrightarrow ^8{\\rm Be}+\\gamma$} \\item{$^8{\\rm Be}+\\he\\longrightarrow\\ctw +\\gamma$} \\item{$\\ctw +\\he\\longrightarrow\\oxy +\\gamma$} \\end{itemize} with the first two being the most important. The third reaction does not occur very often. Therefore, the convective shell at the end of the pulse is composed primarily of helium and carbon. During the transition period between the thermal pulse and the resumption of the interpulse phase the convective shell disappears and the convective envelope can penetrate into this region. The mixing of carbon and helium rich material to the surface layers of the star at the end of a thermal pulse is known as the third dredge-up to distinguish it from the first two dredge-ups which we will describe later."
    },
    "0403/astro-ph0403542_arXiv.txt": {
        "abstract": "Observations with the SZ effect constitute a powerful new tool for investigating clusters and constraining cosmological parameters. Of particular interest is to investigate how the SZ signal correlates with other cluster properties, such as the mass, temperature and X-ray luminosities. In this presentation we quantify these relations for clusters found in hydrodynamical simulations of large scale structure and investigate their dependence on the effects of radiative cooling and pre-heating. ",
        "introduction": "}  Galaxy clusters are the largest gravitationally bound objects in the Universe. They have typical masses of $10^{14}-10^{15}h^{-1}$~M$_\\odot$ and contain hundreds of galaxies within radius of a few Mpc. The intra-cluster medium (ICM) is filled with hot ionized gas, typically at temperatures 1--15 keV, which produces strong X-ray emission and causes spectral distortions in the CMB spectrum via the Sunyaev--Zel'dovich (SZ) effect (see Ref.~\\cite{sunyaev:1972}).  Numerical simulations indicate that the non-baryonic dark matter component in clusters, which is the dominant fraction of their mass, is remarkably self-similar for systems in  approximate state of equilibrium, see e.g. Ref.~\\cite{navarro:1997b}. However, the baryonic gas component does not share the same degree of self-similarity. This is more evident from observations of the $L_X$--$T$ relation in clusters, which is much steeper than is predicted by simple self-similar scaling laws, specially for low-mass systems (see e.g. Refs.~\\cite{edge:1991,xue:2000}). This deviation from self-similarity has been interpreted as due to non-gravitational processes, such as radiative cooling and heating, that raise the entropy of the gas, see e.g. Ref.~\\cite{ponman:1999}.  The purpose of this study is to investigate the correlation between intrinsic properties of clusters found in simulations, and to compare them with analytical scaling laws. We use high-resolution hydrodynamical simulated clusters to assess the impact of radiative cooling and pre-heating on these scaling relations. Here we focus only on scalings involving the SZ integrated signal, $Y$, at redshift zero. The effects on the mass--temperature and X-ray scaling relations have already been analyzed in detail in Refs.~\\cite{muanwong:2001,muanwong:2002} for this same set of simulations. A more detailed analysis of the SZ scaling relations can be found in Ref.~\\cite{dasilva:2003}, whose results supersede those presented at this conference. ",
        "conclusions": "We have studied the effect of our models of cooling and pre-heating on the SZ cluster population by identifying clusters in the simulation boxes and computing their characteristic properties. We correlated the SZ luminosity with other cluster properties and derived scaling relations at redshift zero. The non-radiative simulation reproduces well the scalings predicted by the self-similar model, whereas the inclusion of cooling and pre-heating generally changes the slope and normalization of the scaling relations for those runs. The integrated $Y$-signal is found to be tightly correlated with the mass-weighted temperature and total mass of the cluster.   The author wishes to thank A. Liddle, P. Thomas, S. Kay and O. Muanwong for many fruitful discussions and acknowledges the use of the computer facilities of the Astronomy Unit at Sussex, UK, and CALMIP at Toulouse, France. The simulations used were carried out on a Cray-T3E at EPCC as part of the VIRGO Consortium collaboration."
    },
    "0403/astro-ph0403068_arXiv.txt": {
        "abstract": "We present a spectroscopic survey of the giant stellar stream found in the halo of  the Andromeda galaxy. Taken with  the DEIMOS multi-object spectrograph  on the  Keck2  telescope, these  data  display a  narrow velocity  dispersion of  $11\\pm3\\kms$, with  a steady  radial velocity gradient of $245 \\kms$ over the $125 \\kpc$ radial extent of the stream studied so  far.  This implies  that the Andromeda galaxy  possesses a substantial  dark matter  halo.  We  fit the  orbit of  the  stream in different galaxy potential models. In  a simple model with a composite bulge, disk and  halo, where the halo follows  a ``universal'' profile that is  compressed by  the formation of  the baryonic  components, we find that  the kinematics  of the stream  require a total  mass inside $125\\kpc$ of  $M_{125} =  7.5^{+2.5}_{-1.3} \\times 10^{11}  \\msun$, or $M_{125} >  5.4 \\times  10^{11} \\msun$ at  the 99\\%  confidence level. This is the first galaxy in  which it has been possible to measure the halo  mass distribution  by such  direct dynamical  means over  such a large distance range. The resulting orbit shows that if M32 or NGC~205 are  connected with  the stream,  they must  either trail  or  lag the densest  region of  the stream  by more  than  $100\\kpc$. Furthermore, according to the best-fit orbit,  the stream passes very close to M31, causing  its demise as  a coherent  structure and  producing a  fan of stars that will  pollute the inner halo, thereby  confusing efforts to measure  the properties of  genuine halo  populations.  Our  data show that several  recently identified  planetary nebulae, which  have been proposed as evidence for the existence  of a new companion of M31, are likely members of the Andromeda Stream. ",
        "introduction": "Stellar streams represent the  visible remnants of the merging process by which the halos of galaxies are built up. By studying these streams we  can attempt  to unravel  the formation  of galactic  halos, seeing when, how  and how many  small galaxies arrived and  were incorporated into large galaxies \\citep{helmi99,  johnston01}.  Streams are also of great interest as  probes of the large-scale mass  distribution of the dark halos they reside  in \\citep{johnston99, ibata01a, zhao99}.  This utility  stems from  the fact  that streams  from  low-mass disrupting stellar systems trace the orbit of their progenitor, giving a means to constrain the tangential motion of the stars in the stream.  The stars must  move along  the  stream,  and the  magnitude  of the  tangential velocity  must be such  that, when  the star  is integrated  along the orbit,  it ends  up  with the  same  velocity as  the  stars that  are currently downstream on the same orbit.  Giant stellar  streams may well  be common structures  around galaxies \\citep{pohlen, malin}.   Indeed, the  most conspicuous feature  in the halo of the Milky Way is the giant rosette stream originating from the Sagittarius  dwarf galaxy,  which contains  approximately half  of the high  latitude ($|b|>30^\\circ$)  intermediate age  stars  at distances greater than  $15\\kpc$ \\citep{ibata01a, ibata02,  majewski}.  It would appear that  the Milky Way has  not incorporated into the  Halo a more massive galaxy than the Sagittarius dwarf over the last $\\sim 7\\Gyr$.  This naturally  leads to  the question of  whether the MW  has unusual feeding  habits.    To  answer  this,  we  have   undertaken  a  large photometric study of M31, using the  wide field cameras at the INT and CFHT telescopes to  resolve stars over the entire  disk and inner halo of  that  galaxy  \\citep{ibata01b,ferguson02,mcconnachie}.   This  has given us an unprecedented panoramic  view of the large scale and small scale structure of a disk galaxy. The analysis of this huge dataset is still in progress, but it  has already yielded some surprising results regarding the incidence of substructure  in the halo of M31.  The most prominent  of these  substructures  is a  stream-like over-density  of stars near  the minor axis of  M31 \\citep{ibata01b}, at  first sight a facsimile of  the Sagittarius  Stream around the  Milky Way.   The red giant branch (RGB) stellar density in the halo increases on average by a  factor  of  two  in  the on-stream  regions  and  is  statistically significant at the 50-sigma  level.  Interestingly, this stream points toward the  Andromeda satellites M32  and NGC205, and is  aligned with the outer  isophotes of NGC205, suggesting a  relationship between the Andromeda Stream and these two dwarf galaxies.  If this interpretation is correct, the  stream has to be the  result of previous interactions with M31, as there is  otherwise not enough time to spatially separate it from either of the two dwarf galaxies.  The proximity  of M31 provides us  with an opportunity  to undertake a spectroscopic  survey of  individual stars  within the  stream; indeed Andromeda offers the  only extragalactic giant galaxy in  which such a study  can be  undertaken  with current  instrumentation,  as in  more distant systems such halo substructure  would be smeared into very low surface brightness features.  \\begin{figure*} \\ifthenelse{\\UseFigs=1}{ \\includegraphics[angle=270,width=\\hsize]{M31_StreamOrb.fig02.ps}}{} \\caption{For  convenience, the  left-hand panel  reproduces  the chart shown previously  in Figure~1. The  right-hand panel shows  a sideways view of  the stream, drawn to  the same scale as  the left-hand panel, which displays  the line of sight  depth of the  fields, together with the positions of M32 and NGC~205.   The disk of M31 is highly inclined to  our line  of sight  ($12\\degg5$); the  thick line  is  a schematic representation   of   a   disk   of  radius   $27\\kpc$   inclined   at $12\\degg5$.  Evidently the  stream orbits  close to  the plane  of the Andromeda galaxy.} \\end{figure*}  One useful  property of the  stream is that  it is on a  highly radial orbit. It passes  very close to the centre of  M31, where a comparison of the RGB tip of the stream with that of M31 itself shows them to lie at  the same  distance  \\citep{mcconnachie}; whereas  in the  furthest field that it has been detected  to date, the peak of its giant branch is 0.27~mag  fainter, indicating that it is  $106\\pm20\\kpc$ behind M31 \\citep{mcconnachie}.  This  fortuitous alignment close to  the line of sight, allows  us to measure  directly the potential gradient  over $> 100\\kpc$, and hence measure the halo mass.  The  layout  of this  paper  is  as  follows: section~2  presents  the spectroscopic survey of the Andromeda  Stream, the results of which are used in  section~3 to constrain  the mass of  the dark matter  halo of M31,  and finally  in section~4  we  discuss the  significance of  the results  and present the  conclusions of  this study.  Throughout this work, we  assume a distance of  $780\\kpc$ to M31  \\citep{stanek} and a systemic radial velocity of $-300\\kms$ \\citep{devaucouleurs}. ",
        "conclusions": "In  this first  kinematic study  of the  giant stellar  stream  in the Andromeda galaxy,  we have  been able to  measure the  radial velocity gradient along  the stream from the outermost  field currently probed, $125\\kpc$ distant from the centre of M31, down to an inner field $\\sim 20\\kpc$ from the centre of  that galaxy.  Over this huge distance, the (projected)   radial  velocity  changes   by  $245\\kms$,   implying  a de-projected  velocity  difference of  $\\sim  280\\kms$. This  velocity gradient is  used to  obtain a zeroth  order analytic estimate  of the mass of the  halo, which for simple halo-only galaxy  models such as a logarithmic halo  or an NFW halo,  implies a mass  inside $125\\kpc$ of $M_{125} =  7.6 \\pm 1.2 \\times  10^{11} \\msun$ and $M_{125}  = 6.4 \\pm 1.3  \\times   10^{11}  \\msun$,   respectively.  In  both   cases,  the uncertainty associated  with a possible tangential velocity  of M31 of $300\\kms$, is $\\sim 5$\\%.  We also  investigate more realistic solutions, allowing  the stream to have a non-radial orbit, and taking  a galaxy model that is the sum of a disk,  bulge and dark halo  \\citep{klypin02}. The dark  halo of this model is a  perturbation on a spherical NFW model,  to account for the adiabatic contraction  of the dark  matter as the  baryonic components form.  If we disregard the distance data in Fields 12 and 13, since we cannot be  certain that  the stream is  present in those  regions, the most likely mass of this  galaxy model is $M_{125} = 7.5^{+2.5}_{-1.3} \\times  10^{11}\\msun$, with  a lower  limit of  $M_{125} =  5.4 \\times 10^{11} \\msun$  (at 99\\% confidence). This result  is fully consistent with the zeroth order analytic estimates discussed above, and suggests that the derived mass is not very sensitive to the adopted mass model. Furthermore, there is a reasonable agreement of the resulting rotation curve  with the  kinematics of  previously-observed disk  tracers (see Figure~6),  despite  the fact  that  the model  was  only  fit to  the kinematics of the  stream.  This agreement is due in  part to the fact that we  have taken previously-fitted  models for the disk  and bulge, but the halo contribution to the total rotation curve dominates beyond $\\sim 20\\kpc$, and it is at  these large distances that we have fitted the model to the stream  data.  This confers further confidence on the derived mass  model. The uncertainty on  the mass estimate  due to the possible  tangential  velocity  of  M31  is  likely  not  very  large, approximately 7\\% for a tangential  velocity of $300\\kms$.  One of the main uncertainties  in the present  analysis is the flattening  of the halo, which  we have not  explored, as the  current data set  does not provide  sufficient  constraints.    Future  studies,  fitting  N-body simulations  to a  larger kinematic  sample  of stream  stars, can  be expected  to improve  the mass  estimate and  also constrain  the halo shape.  \\begin{figure} \\ifthenelse{\\UseFigs=1}{ \\includegraphics[angle=0,width=\\hsize]{M31_StreamOrb.fig07.ps}}{} \\caption{The spreading of orbits after a close encounter with M31. The thick-line orbit reproduces the best-fit orbit in the realistic galaxy potential previously shown in  Figure~4, integrated from the centre of Field~6.  The  thin-lines show similar orbits, starting  from the same spatial  position, but with  velocities perturbed  by a  random offset drawn  from  a Gaussian  distribution  of  dispersion $11\\kms$.   This fanning-out of stars  on nearby orbits will lead  to the disappearance of  the  Andromeda  Stream,  and   implies  that  it  is  a  transient phenomenon. The small  square on the minor axis  shows the position of the \\citet{brown} ACS field (not  drawn to scale!). All other markings are as in Figure~1.} \\end{figure}  It is  only recently that measurements  of the mass of  M31 beyond the edge of its gaseous disk have been possible. \\citet{courteau} analysed the  velocities of  7 Andromeda  satellites, finding  a total  mass of $13.3\\pm 1.8 \\times 10^{11} \\msun$. In contrast, using a larger sample of 10 satellite galaxies, 17 globular clusters and 9 planetary nebulae as test  particles, \\citet{evans00a} found that the  most likely total mass  of M31 is  $12.3^{+18}_{-6} \\times  10^{11}\\msun$, approximately half   of  their   Milky  Way   estimate  of   $19^{+36}_{-17}  \\times 10^{11}\\msun$.  With improved radial  velocities of the M31 satellites they were later able to reduce their M31 mass uncertainties, finding a value   of  the   total   mass  of   $\\sim  7.0^{10.5}_{-3.5}   \\times 10^{11}\\msun$  \\citep{evans00b}, which  is fully  consistent  with our result of  $M_{125} = 7.5^{+2.5}_{-1.3}  \\times 10^{11}\\msun$ (defined within a radius of $125\\kpc$).  However, a definitive statement of the relative  masses  of  M31  and   the  Milky  Way  awaits  an  improved measurement for the Milky Way.  The best  fit orbit in the best  fit potential is prograde  and in the region where it  is currently observed, it lies close  to the plane of M31.   Thus it  appears  that the  orbit  of the  Andromeda stream  is peculiar  in  being  extremely  radial,  passing  very  close  (within $2\\kpc$) of  the centre  of M31.  This  requires very  special initial conditions.  The  stream stars, which are  spatially narrowly confined in Fields~1 to 8, will  diverge dramatically upon passing close to M31 to form  a low-density fan-like  structure, since orbits  that deviate only slightly  from the  orbit displayed in  Figure~4 on  the plunging part  of their  course  will  take very  different  paths after  being accelerated around the centre of M31, as demonstrated in Figure~7. The fanning-out of the stream is  likely to confuse efforts to measure the metallicity  and  age  of  the  M31 halo;  for  instance,  the  recent discovery  by \\citet{brown}  of a  young  halo component  in M31  from main-sequence fitting  of an  extremely deep ACS  field may be  due to stream contamination (see Figure~7).  Thus  the stream  may be  in the  process of  vanishing as  a coherent structure, providing a supply of metal-rich stars into the halo.  This also suggests  that the  stream was removed  from its  progenitor less than an orbital period ago (the pericentre to pericentre period of the continuous-line orbit in Figure~4 is $1.8\\Gyr$), as we would otherwise not observe  the structure as a  stream.  The ephemeral  nature of the stream  implies that  the progenitor  must have  survived  until $\\sim 1.8\\Gyr$ ago. As we discuss further below, the progenitor was probably of  low mass, implying  that the  rate of  decay of  its orbit  due to dynamical friction was  slow, so it followed (or  continues to follow) an orbit close to the current orbit of the stream.  However, any dwarf galaxy on the derived orbit  must have experienced extreme tides as it repeatedly passed close to the centre  of M31.  One option is that the progenitor was a very dense  dwarf galaxy that was sufficiently robust to survive  the huge tides.  This  brings to mind M32  as a candidate, though  detailed  numerical  modeling  is  required  to  examine  this possibility.   The alternative option  is that  the progenitor  of the stream deflected off another halo object, sending it plunging into the current orbit, analogous to  the suggestion by \\citet{zhao98} to explain the longevity of the Galactic satellite Sagittarius.  However, the connection with M32 presents some difficulties.  Although M32  appears  to  reside  in  the stream,  its  velocity  is  markedly different. For M32  to be associated with the  stream would require it to be at a different phase  in the orbit (either lagging or trailing). Furthermore, the low velocity dispersion of the stream would appear to preclude  M32 as  its  progenitor (which  has  $\\sigma_v \\sim  50\\kms$ outside of  the nucleus, \\citealt{vandermarel94}),  though this cannot be  confirmed  without  a  detailed  dynamical study.   The  case  for association of NGC~205 with the stream also appears weaker given these kinematic measurements, since the best fit orbit does not overlap with it in phase-space.  Future studies  may allow us to examine this issue in more  detail by  following the stream  beyond the  region currently probed with kinematics.  The  velocities of the stream stars presented here also  shed light on the  recent identification of  a possible new companion to M31, And VIII \\citep{morrison}.  The radial velocities of those  planetary  nebulae,  which   are  located  close  to  M32  (see Figure~1),  appear  to  have  radial  velocities  consistent  with  an extrapolation of the Stream, as they lie close to the straight line in the right-hand panel of Figure~1.  This would suggest that And VIII is most likely  part of the Andromeda  Stream, situated in  the region of highest  over-density reported  by \\citet{ibata01b}.   However,  it is interesting to note that the radial velocities of the \\citet{morrison} planetary nebula sample tend to increase towards the North (showing an apparent  positive  gradient in  the  right-hand  panel of  Figure~1), whereas  the Stream  stars  have a  negative  velocity gradient.   The connection between the two  structures therefore merits to be examined more   carefully.   During   the  refereeing   process,  a   study  by \\citet{merrett}  was presented which  also investigates  the planetary nebulae around  M31.  However, though  their survey has  also detected PNe in  a region at the base  of the stream (near  our Field~8), their interpretation  is  inconsistent  with  the kinematics  of  the  stars reported here.   The direction of motion  of the orbit  they derive is opposite to ours,  and the path of their  orbit, which intercepts many M31-disk  PNe is  substantially different  to the  orbit that  we have fitted.  It  is possible  that their finding  reveals the  presence of another kinematic structure in Andromeda.  The  majority of the  stream stars  that were  surveyed have  a narrow velocity  dispersion  of  $11\\pm  3\\kms$, though  slightly  skewed  to positive velocities.  The  fact that our line of  sight looks down the stream (see Figure~2),  so that we probably see stars  over a range of distance along the line of sight, and hence at different phases in the orbit,  will tend to  render the  observed velocity  dispersion higher than  the  intrinsic velocity  dispersion.   This  indicates that  the progenitor was  most likely  a low mass  dwarf galaxy.  The  Milky Way satellite  Sagittarius, which  has a  velocity dispersion  of $11\\kms$ \\citep{ibata97}, also has a gigantic stellar stream, but with a larger velocity dispersion of $20\\kms$  \\citep{yanny}. However, it is unclear at  present whether  the  lower velocity  dispersion  measured in  the Andromeda stream  compared to the Sagittarius stream  implies that its progenitor was of lower mass than Sagittarius, or not."
    },
    "0403/astro-ph0403297_arXiv.txt": {
        "abstract": "{We present integrated colours, integrated spectral energy distributions, and absorption-line indices, for instantaneous burst solar-metallicity binary stellar populations with ages in the range $1 - 15\\,$Gyr. By comparing the results for populations with and without binary interactions we show that the inclusion of binary interactions makes the appearance of the population substantially bluer -- this is the case for each of the quantities we have considered. This effect raises the derived age and metallicity of the population. Therefore it is necessary to consider binary interactions in order to draw accurate conclusions from evolutionary population synthesis work. ",
        "introduction": "The majority of current studies of the integrated light of stellar populations using the evolutionary population synthesis (EPS) method -- the most direct approach for modelling galaxies -- have tended to focus solely on the evolution of single stars. However, observations tell us that upwards of 50\\% of the stars populating galaxies are expected to be in binary or higher-order multiple systems (Duquennoy \\& Mayor 1991; Richichi et al. 1994, for example). Binary evolution, if the component stars are close enough to exchange mass, can drastically alter the evolution path of a star as expected from single star evolution. Binary interactions can also create some important classes of objects, such as blue stragglers (BSs: Pols \\& Marinus 1994), and subdwarf B stars (sdBs, also referred as extreme horizontal branch [EHB] stars: Han et al. 2002; 2003). Therefore binary stars have the potential to play an important role in determining the overall appearance of any realistic stellar population and their evolution should be accounted for in population synthesis models. A few EPS studies have previously made advances in this direction (e.g. Pols \\& Marinus 1994; Cervino et al. 1997) but only for specialised cases and not to the degree that we present here.  In this letter we assume that all stars are born in binaries and born at the same time, i.e. an instantaneous binary stellar population (BSP). We then model any interactions within these binaries in our EPS models to investigate the effects on the integrated colours, integrated spectral energy distributions (ISEDs), and absorption feature indices. The outline of the paper is as follows: we describe our EPS models and algorithm in Sect. 2, present our results and some discussion in Sect. 3, and then finally in Sect. 4 we give our conclusions. ",
        "conclusions": "We have simulated realistic stellar populations composed of 100\\% binaries by producing $2 \\times 10^5$ binary systems using a Monte Carlo technique. We computed the integrated colours, ISEDs and Lick/IDS absorption feature indices for these instantaneous burst BSPs with and without binary interactions. In comparison we find that modelling binary evolution, and the additional classes of stars that this produces, leads to bluer integrated colours, ISEDs and Lick/IDS absorption feature indices, and therefore makes the derived age and metallicity raise. In this letter we have only considered the effects of binaries for solar metallicity BSPs -- more detailed studies will be given later."
    },
    "0403/astro-ph0403012_arXiv.txt": {
        "abstract": "New precision in measuring extragalactic distances using supernovae has confirmed with high probability an accelerating increase in redshift with distance.  This has been interpreted as implying the existence of dark energy in an expanding and accelerating, flat universe.  A more logical explanation of these observations follows directly from an observation made by Erwin Schr\\\"odinger in 1939 that in a closed Friedmann universe $every$ quantum wave function changes with spacetime geometry.  Double the size of the universe and both the wavelengths of photons and the sizes of atoms double.  When the evolution of atoms and photons are combined, the meaning of Hubble redshift is reversed.  Redshift is characteristic of contracting universes.  The magnitude-redshift curve for a contracting universe has exactly the accelerating form recently observed and is in excellent quantitative agreement with the data of Riess et al. 1998, Knop et al. 2003, and others.  An observed maximum redshift of 1.3  gives a minimum age estimate for the universe of $114$ billion years.  The time until collapse is estimated to be 15 billion years or less. ",
        "introduction": "Supernovae provide an excellent opportunity for measuring extragalactic distances. The observations of these exploding stars using the best earth and space-based telescopes, a physical understanding of supernovae evolution, and sophisticated analyses of their spectra have led to the surprising conclusion that the expansion of the universe seems to be accelerating (Perlmutter \\& Schmidt 2003).  The upward trend of the magnitude versus redshift curve at high redshift is clear in Figure 1 (Riess et al. 1998).  This departure from a linear or decelerating trend has been explained as eternal expansion with a positive cosmological constant and a current acceleration of expansion.  This conclusion is based on a conventional interpretation of redshift as the result of photon wavelength increasing in an expanding Friedmann universe.  This was one idea originally suggested by Hubble \\& Tolman (1935), but it is interesting to note that they concluded their paper with  \\begin{quote} It also seemed desirable to express an open-minded position as to the true cause of the nebular red-shift, and to point out the indications that spatial curvature may have to play a part in the explanation of existing nebular data. \\end{quote}  The conclusion drawn here is that spatial curvature indeed plays a critical part in explaining cosmic redshift.  It is the combination of changes in photons during their transit time and changes in atoms between the time of photon emission and later when they provide the wavelength standard, that defines redshift. ",
        "conclusions": "Recent supernovae data may be interpreted as providing experimental confirmation of the theoretical predictions of the coevolution of spacetime geometry and quantum wave functions made by Schr\\\"odinger (1939), Sumner (1994), and Sumner \\& Sumner (2000).  This explanation of the magnitude-redshift relationship comes directly and unambiguously from the roots of the best-tested physical theories of the 1930's.  \\bigskip \\medskip  The author thanks Dawn Sumner for many insights and helpful suggestions.  \\newpage"
    },
    "0403/astro-ph0403224_arXiv.txt": {
        "abstract": "We construct models of core-collapse supernovae in one spatial dimension, including rotation, angular momentum transport, and viscous dissipation employing an $\\alpha$-prescription.  We compare the evolution of a fiducial 11\\,M$_\\odot$ non-rotating progenitor with its evolution including a wide range of imposed initial rotation profiles ($1.25<P_0<8$ s, where $P_0$ is the initial, approximately solid-body, rotation period of the iron core). This range of $P_0$ covers the region of parameter space from where rotation begins to modify the dynamics ($P_0\\sim8$ s) to where angular velocities at collapse approach Keplerian ($P_0\\sim1$ s). Assuming strict angular momentum conservation, all models in this range leave behind neutron stars with spin periods $\\lesssim10$ ms, shorter than those of most radio pulsars, but similar to those expected theoretically for magnetars at birth.  A fraction of the gravitational binding energy of collapse is stored in the free energy of differential rotation.  This energy source may be tapped by viscous processes, providing a mechanism for energy deposition that is not strongly coupled to the mass accretion rate through the stalled supernova shock. This effect yields qualitatively new dynamics in models of supernovae. We explore several potential mechanisms for viscosity in the core-collapse environment: neutrino viscosity, turbulent viscosity caused by the magnetorotational instability (MRI), and turbulent viscosity by entropy- and composition-gradient-driven convection.   We argue that the MRI is the most effective. We find that for rotation periods in the range $P_0\\lesssim5$ s, and a range of viscous stresses, that the post-bounce dynamics is significantly affected by the inclusion of this extra energy deposition mechanism; in several cases we obtain strong supernova explosions. ",
        "introduction": "\\label{section:introduction}  Employing the now-standard suite of microphysics and a full solution to the Boltzmann equation for all neutrino species, the neutrino mechanism of core-collapse supernovae fails  in one spatial dimension (Rampp \\& Janka 2000, 2002; Liebend\\\"{o}rfer et al. 2001a,b; Mezzacappa et al.~2001; Thompson et al.~2003; Liebend\\\"{o}rfer et al.~2004). The initial \"bounce shock,\" formed as the core density reaches $\\sim$$2\\times10^{14}$\\,g\\,cm$^{-3}$ and nuclei undergo a phase transition to free nucleons, stalls almost immediately due to a combination of neutrino losses, the ram pressure of the infalling mantle, and the dissociation of nuclei into neutrons and protons behind the shock. The revival of the shock to an asymptotic energy of order $10^{51}$\\,erg is the focus of modern supernova theory.  Any mechanism for this revival must fundamentally rely on the transfer of gravitational binding energy to the post-shock mantle.   As its name makes clear, the \"neutrino mechanism\" (Bethe \\& Wilson 1985) relies on neutrino interactions (primarily $\\nu_e n\\leftrightarrow pe^-$ and $\\bar{\\nu}_e p\\leftrightarrow n e^+$) to transfer binding energy to the shocked matter in the post-bounce epoch. In spherical models, the stalled bounce shock remains trapped forever and continues to accrete the overlying  progenitor. Although some models show that multi-dimensional effects might be necessary for success of the neutrino mechanism (Herant et al.~1994; Burrows et al.~1995; Janka \\& M\\\"{u}ller 1995, 1996; Fryer et al.~1999; Fryer \\& Warren 2002), multi-dimensional models employing more sophisticated neutrino transport fail to explode (Janka et al.~2002; Buras et al.~2003) -- albeit marginally.  In constructing one-dimensional models, the focus of the supernova community has been largely on refining details of the neutrino mechanism.  These details include a fuller description of the nuclear physics and the neutrino interactions and, over the last ten years, a precise handling of neutrino transport  (Mezzacappa \\& Bruenn 1993a,b,c; Burrows et al.~2000; Rampp \\& Janka 2000, 2002; Liebend\\\"{o}rfer et al. 2001a,b; Mezzacappa et al.~2001; Thompson et al.~2003). Precision neutrino transport is important because the degree of neutrino heating behind the shock controls the subsequent dynamics and because the optical depth of this region is $\\lesssim0.1$. For this reason, the most recent one-dimensional models are fully spectral and multi-angle (Rampp \\& Janka 2000, 2002; Liebend\\\"{o}rfer et al. 2001a,b; Mezzacappa et al.~2001; Thompson et al.~2003); neutrino transport also presents the most significant technical hurdle in calculating realistic models of core collapse in two or three spatial dimensions (Livne et al.~2003).  It is possible that yet more refined treatments of the radiation hydrodynamics or subtle changes in the microphysics of neutrino interactions may yield explosions in one-dimensional models via the neutrino mechanism. However, it is also possible that the solution to the supernova problem does not lie in such technical improvements or in an accumulation of five and ten percent effects, but instead in new physics and order unity phenomena not previously identified or fleshed out.  Such  improvements might include new massive pre-collapse progenitors with very different core entropy and lepton fraction, new finite temperature high-density nuclear equations of state (including exotic meson species, quark matter, phase transitions, etc. and the associated neutrino microphysics), or the interplay between rotation and magnetic fields. In this paper, we take a step towards addressing the last of these issues in detail.  When the core of a rotating massive star collapses, some of the gravitational binding energy of collapse is stored in a strong shear profile.  This free energy reservoir may be tapped by viscous processes. In this way, the binding energy of collapse, temporarily trapped in differential rotation, may be transferred to the matter as thermal energy. This is qualitatively different from the standard neutrino mechanism, which relies solely on neutrino interactions to communicate the binding energy of the collapsed iron core to the rest of the star.  We find that dissipation of shear energy can significantly affect the dynamics of supernovae; models that would otherwise fail succeed as a result of viscous dissipation. The most promising source of viscosity is magnetic stresses generated by the magnetorotational instability (MRI; Akiyama et al.~2003) and perhaps magneto-convection (Duncan \\& Thompson 1992; Thompson \\& Duncan 1993). Although magnetic fields in excess of $10^{15}$ G are generated, no large iron core seed field is required for the success of this mechanism, since any seed field grows exponentially on a rotation period.  In \\S\\ref{section:motivate}, we provide some context and motivation for the present study and we review some previous research on core collapse with rotation. In \\S\\ref{section:model}, we describe the technical specifications of our one-dimensional models including rotation, viscous dissipation of shear energy, and the corresponding transport of angular momentum. Section \\ref{section:viscous} describes mechanisms for viscosity, including neutrino interactions, the MRI, and hydrodynamic convection. In \\S\\ref{section:rotation}, we present results from rotating models of core collapse including the centrifugal force, but without viscous dissipation.  We also provide a general discussion of the conditions for explosion and an explanation for why these models and non-rotating models fail. We compare the magnitude of the viscous processes discussed in \\S\\ref{section:viscous}  in \\S\\ref{section:compare}. In \\S\\ref{section:dissipation}, we present results from calculations of core collapse including viscous dissipation.  We  show that explosions are obtained, we explain why these models succeed, and we delineate the successful parameter space. In \\S\\ref{section:summary}, we summarize our results and discuss the implications of our findings. ",
        "conclusions": "\\label{section:summary}  In this paper we forward the idea that some of the gravitational binding energy of core collapse may be trapped in differential rotation and that this free energy can be tapped on a viscous timescale ($\\sim1$ s in the gain region) and deposited as thermal energy in the gas. We identify magnetic stresses generated by the magnetorotational instability as the most likely origin of this viscous energy dissipation. We have constructed models of core collapse, adding rotation and viscosity to our existing algorithm for spectral, multi-angle radiation hydrodynamics. We find that viscous dissipation can enhance the heating in the gain region sufficiently to yield explosions in models that would otherwise fail. Initial spin periods of the roughly solid-body iron core should be $P_0\\lesssim5$ seconds for this mechanism to be robust.  We emphasize several important points about this result. First, our simulations are one-dimensional and as such do not include the convective motions that have been shown to at least aid, if not enable, explosions (Herant et al.~1994; Burrows, Hayes, \\& Fryxell 1995; Janka \\& M\\\"{u}ller 1996; Fryer et al.~1999; Janka et al.~2002; Buras et al.~2003).   We expect that multi-dimensional models with {\\it slower} initial progenitor rotation periods, including the effects of viscous dissipation consistently, will also yield explosions. Second, the effects of dissipation depend on where, when, and how quickly dissipation acts. Improved understanding of these issues will require MHD simulations and further progenitor modeling. Third,  the values of  $P_0$ we require to significantly effect the dynamics of core collapse are fully consistent with the rotating progenitor models of Heger et al.~(2000).  However, there is considerable uncertainly in these progenitor rotation rates.  In particular, the models of Heger et al.~(2003) and Heger et al.~(2004), which attempt to include angular momentum transport by magnetic torques consistently in one dimension, find much slower rotation rates for the iron core just prior to collapse.   For low mass progenitors in the range $12-15$ M$_\\odot$, the rotation periods obtained by  Heger et al.~(2004) are larger than our $P_0\\lesssim5$ s fiducial bound. Heger et al.~(2004) do, however, find that the rotation period decreases with increasing progenitor mass.  Thus, without a detailed study beyond the scope of this paper, we cannot make a definite statement as to whether rotational effects are important in their core-collapse progenitor models.  Finally, our limit on $P_0$ is, at face value, inconsistent with the observed distribution of spin periods of young pulsars, which are in the range of many tens to many hundreds of milliseconds (Kaspi \\& Helfand 2002). There is, of course, a danger in such comparisons since it is very difficult to establish the true birth periods of neutron stars observationally.  It is also unclear if the spindown of very young neutron stars is well-understood theoretically as there are a number of ways to remove angular momentum from the neutron star during the explosion phase (as in the models presented here with viscous dissipation) or during the Kelvin-Helmholtz cooling epoch (\\S\\ref{section:motivate}).  Given these uncertainties, we reiterate that the range of rotation rates considered here are expected theoretically in models of magnetar formation (Duncan \\& Thompson 1992; Thompson \\& Duncan 1993; Kouveliotou et al.~1999), a possibly large subset of all supernovae.  Our work has highlighted two primary impacts of rotation on the mechanism of core-collapse supernovae: 1) Most importantly, rotational kinetic energy tapped by viscous processes can help drive explosions (\\S\\ref{section:dissipation}), and 2) Rotation lowers the effective gravity in the core, increasing the radius of the stalled shock and the size of the gain region (\\S\\ref{section:rotation}, Fig.~\\ref{plot:mdshbp}).  Since ejection is inhibited by the deep potential well, this helps facilitate explosion. In addition to these, there are several multi-dimensional effects of rotation on core-collapse that remain to be fully explored (Burrows, Ott, \\& Meakin 2003):  1) Rotation might generate vortices that can dredge up heat from below the neutrinospheres and thereby enhance the driving neutrino luminosities. 2) Rotation can result in a large pole-to-equator anisotropy in the mass accretion rate after bounce due to the centrifugal barrier along the poles. Finally, 3) rotation results in pole-to-equator anisotropies in the driving neutrino flux that have never before been accurately calculated. Since the polar neutrino flux should be enhanced and the polar mass flux should be lowered (point 2), explosion along the poles with bipolar morphology seems that much more likely (Kotake, Yamada, \\& Sato 2003; Burrows \\& Goshy 1993; Burrows, Ott, \\& Meakin 2003).  Thus, such bipolarity, and the consequent optical polarization of the debris, are not exclusive signatures of MHD-driven explosions (Akiyama et al. 2003; Symbalisty 1984), but may be a natural consequence of any rapidly rotating models.  In addition to the effects of rotation on the supernova mechanism, we have also shown how the observational signature of core collapse might be affected by rotation by identifying the systematics in the neutrino luminosity and average energy as a function of time, initial rotation rate, and progenitor (\\S\\ref{section:rotation}, Fig.~\\ref{plot:ltet}). If the neutrino radiation from the next galactic supernova is observed in modern neutrino detectors like SuperKamiokande and SNO, we may hope to distinguish rapidly rotating core collapse via its neutrino signature.  In addition to the difference highlighted in Fig.~\\ref{plot:ltet} we note that rapidly rotating models may leave an appreciable amount of shear energy ($\\sim10^{52}$ erg) trapped in the inner core, which is stable to the MRI and convection.  The timescale to dissipate this energy is uncertain, but could be $\\sim 1-10$ seconds or longer (eq.~\\ref{nuspruit}; Spruit 2002) as the PNS cools.  The timescale for dissipation of this energy will likely differ from the Kelvin-Helmholtz timescale (Burrows \\& Lattimer 1986; Pons et al.~1999). Thus, rapidly rotating PNSs may have a unique long-timescale neutrino signature.  Finally, we have assessed the relative role of neutrino viscosity, the turbulent viscosity caused by the magnetorotational instability, and the viscosity of purely hydrodynamic convection. We find that, for the rotation rates considered here, the MRI and convection may be comparable in importance, but that both dominate microscopic neutrino viscosity everywhere and at all times during the collapse and accretion epochs.  Turbulent magneto-convection should be generic for all modestly rotating progenitors and of critical importance to the dynamics of models with fairly rapid rotation, $P_0\\lesssim5$ s (Duncan \\& Thompson 1992; Thompson \\& Duncan 1993). Thus, in models with these rotation rates, magnetohydrodynamic effects should not be ignored. The simplicity of the algorithm employed here, which assumes spherical symmetry in a system that so clearly violates it and an $\\alpha$ viscosity, makes it clear that more work must be done to ascertain the full effects of MHD in rotating core collapse."
    },
    "0403/astro-ph0403154_arXiv.txt": {
        "abstract": "We present the {H{\\sc i}} data for 5 spiral galaxies that, along with their H$\\alpha$ rotation curves, are used to derive the distribution of dark matter within these objects. A new method for extracting rotation curves from {H{\\sc i}} data cubes is presented; this takes into account the existence of a warp and minimises projection effects. The rotation curves obtained are tested by taking them as input to construct model data cubes that are compared to the observed ones: the agreement is excellent. On the contrary, the model data cubes built using rotation curves obtained with standard methods, such as the first-moment analysis, fail the test. The {H{\\sc i}} rotation curves agree well with the H$\\alpha$ data, where they coexist. Moreover, the combined H$\\alpha$ + {H{\\sc i}} rotation curves are smooth, symmetric and extended to large radii.  The rotation curves are decomposed into stellar, gaseous and dark matter contributions and the inferred density distribution is compared to various mass distributions: dark haloes with a central density core, $\\Lambda$ Cold Dark Matter ($\\Lambda$CDM) haloes (NFW, Moore profiles), {H{\\sc i}} scaling and MOND. The observations point to haloes with constant density cores of size $r_{core} \\sim r_{opt}$ and central densities scaling approximately as $\\rho_0 \\propto r_{core}^{-2/3}$. $\\Lambda$CDM models (which predict a central cusp in the density profile) are in clear conflict with the data. {H{\\sc i}} scaling and MOND cannot account for the observed kinematics: we find some counter-examples. ",
        "introduction": "One of the tantalising and yet unanswered questions of contemporary astrophysics is that concerning the nature of dark matter (DM). While its existence has been inferred for several decades, only in the recent past its amount has been carefully assessed on various astrophysical scales. On galaxy scales, the classical and still most powerful tool to unveil DM is that of rotation curves. Compared to the pioneering work of \\citet*{Ru:80} and \\citet{Bo:81}, modern studies take advantage from greatly improved observational techniques, facilitating much more significant validations of cosmological theories, such as Cold Dark Matter (CDM). Although the prime astrophysical importance of the juxtaposition of observations and models was assessing the amount of DM in galaxies \\citep{PS:88}, the timed question has become its {\\it distribution}. This issue directly pertains to the nature of this component, since a precise knowledge of the density profile $\\rho(R)$ is likely to be decisive: the DM nature is likely to determine the halo properties.  Yet, a wealth of observations have not been able to settle but the most fundamental issues, such as the validity of the CDM scenario on galactic scales (e.g. \\citealt{NFW:96}, \\citealt{Mo:98}). Some of the shortcomings on the observational side were related to the limitations of optical rotation curves (internal rms, limited radial extent, dust contamination), while radio studies were affected by the well-known limited spatial resolution. Nevertheless, progress has been made, due to the increasing number of investigations on dwarf spirals and low-surface brightness (LSB) galaxies. These studies cast serious doubts on one of the fundamental properties of CDM haloes, namely their 'cuspiness' in the galaxy centres, and infer the presence of large core radii in the DM density profiles (see e.g. \\citealt{Fl:94}, \\citealt{Mo:94}, \\citealt{Bur:97}, \\citealt{Kr:98}, \\citealt{Mc:98}, \\citealt{SB:00}, \\citealt{dB:01}, \\citealt{Sa:01}, \\citealt{dBB:02}, \\citealt*{Sa:03}, \\citealt*{W:03}, \\citealt{Si:03}). However, due to the many steps in the data analysis, there can be subtle systematic errors that could distort the results or in any case render the results very poorly constrained. This has triggered off the recent debate concerning the reliability of the data and how well the mass models are really constrained. There are claims that the observations could actually be consistent with the dark matter density profiles predicted by the CDM simulations, not only by considering the {H{\\sc i}} data alone (\\citealt{vdB:00} and \\citealt{vdBS:01}), but also by combining H$\\alpha$ and {H{\\sc i}} data (\\citealt{P:02} and \\citealt{S:03}). This is the reason why particular care should be taken in choosing a suited sample and in performing the data analysis. Note that recent simulations (e.g. \\citealt{N:03}) do not converge to a well-defined value of the inner slope down to the resolution limit (about 1 kpc), even though the slope of the dark matter density profile (defined as $ -d{\\rm ln}\\rho/d{\\rm ln}r$) at 1\\% of the virial radius is still about 1.2 for a typical galaxy. Notice also that the observational results on spiral galaxies show a discrepancy with the standard $\\Lambda$CDM predictions well beyond 1 kpc, i.e. well beyond the resolution limit of the simulations.  In this paper we study a {\\it sample} of galaxy rotation curves ideal for deriving the properties of the dark matter haloes around galaxies; it consists of five late-type bulge-less normal (high surface brightness) spiral galaxies; compared to the sample of \\citet{BS:01} we have the same resolution in the central regions, but a larger spatial extension allowing us to determine the size of the core radii. One of the main concerns is that the rotation curve has both a high spatial resolution and a large extension, i.e. beyond the optical radius; this is typically achieved by combining optical (H$\\alpha$) and radio ({H{\\sc i}}) data (see e.g. \\citealt{dBB:02} and \\citealt{S:03}). The former provide the necessary high resolution (1$\\arcsec \\dots 2 \\arcsec$) while the latter allow us to trace the potential out to large radii, typically 2--3 times the optical radius. It has been known for a long time that {H{\\sc i}} data usually suffer from the lack of resolution, which can affect the results, especially in the innermost parts; this problem is nicely reviewed by \\citet{vdB:00}.  The difficulty in deriving reliable rotation curves from {H{\\sc i}} data cubes resides in the fact that a data cube is 4-dimensional (RA, dec, flux density and radial velocity), while the rotation curve is 2-dimensional (rotation velocity vs. radius). This means that starting from the data cube we have to take into account two types of considerations: one concerning the {\\it positions} in which the rotation curve should be traced, i.e. how to reduce the spatial dimensions from (RA, Dec) to galactocentric radius. This is usually achieved by taking the position-velocity diagram along the major axis or by fitting concentric ellipses around the centre (tilted-ring modelling of the velocity field). The other consideration concerns the {\\it rotation velocities} that should be attributed to these galactocentric radii: at the positions defined above we want to associate only one of the possible velocities that could in principle be derived from the spectra at single points (flux density vs. radial velocity), i.e. the radial component of the rotation velocity at that radius. This is usually done via different methods, most of them assuming symmetry of the profiles, like the first-moment analysis (the intensity weighted mean), or the single Gaussian fitting; therefore, when the profiles are not symmetric, i.e. in galaxies with a high inclination and/or a poor to intermediate resolution, these methods cannot be applied.  \\begin{table*} \\caption{Physical parameters of the galaxies. M$_I$ is the I-band absolute magnitude and M$_{\\rm HI}$  is the total {H{\\sc i}} mass, $r_{\\rm HI}$ is the {H{\\sc i}} radius (defined as the radius at which the surface density drops below 1 M$_{\\odot}$ pc$^{-2}$) and $r_{d}$ is the exponential scale length. $r_f$ is the farthermost radius with data and $r_{opt}$ is the optical radius (defined as 3.2$r_d$) and RR us the radio resolution scale. The dynamical mass M$_{dyn}$ is determined at $r_{f}$. \\vspace{0.3cm}} \\begin{tabular}{l l l l l l} \\hline \\hline Galaxy name                      & ESO 116-G12        & ESO 287-G13       & ESO 79-G14 & NGC 1090 & NGC 7339    \\\\ \\hline M$_I$ (mag)                      & -20.0              & -21.7             & -21.4      & -21.8    & -20.6       \\\\ M$_{\\rm HI}$  (M$_{\\odot}$)      &1.5$\\times 10^{9}$  & 1.1$\\times 10^{10}$&3.5 $\\times 10^{9}$ &8.5$\\times 10^{9}$ &5.5$\\times 10^{8}$    \\\\ Inclination ($^{\\circ}$)         & 74                 &78                 &84          &64        &79           \\\\ $r_{\\rm HI}/r_{d}$          & 6.7                &7.4                & 4.8        &8.8       &4.1          \\\\ $r_f$ (kpc)                 & 11.4               & 24.9             & 17.7       &     29.7   & 5.4        \\\\ $r_{opt}$ (kpc)             & 5.4               & 10.5            & 12.4        &     10.9   & 4.9         \\\\ $r_f$/RR                    & 5.4               & 5.8            & 5.0        & 10.6      &  4.6          \\\\ M$_{dyn}$ (M$_{\\odot}$)     & 3.3$\\times 10^{10}$ & 1.9 $\\times 10^{11}$ & 1.3$\\times 10^{11}$ & 1.8$\\times 10^{11}$ & 3.1$\\times 10^{10}$ \\\\ \\hline \\label{physical} \\end{tabular} \\end{table*}  We have developed a new method for both of these considerations: a) concerning the positions, we traced the rotation curve along the ridge of the warp (this step is discussed in \\citealt{V:04}), when the velocity field was not sufficiently sampled (i.e., for the galaxies ESO 116-G12, ESO 287-G13, ESO 79-G14 and NGC 7339) and the tilted-ring modelling of the velocity field was therefore not possible; b) for the velocities we only considered the velocity side of the profiles opposite to the systemic velocity, and then corrected for the effects that artificially broaden the profile, i.e.: the turbulence of the ISM, the instrumental velocity resolution and the beam-broadening. The present paper concentrates on step b).  When both steps a) and b) were applied we called our method WAMET (WArped Modified Envelope Tracing method); when only step b) was applied and a velocity field was constructed (i.e., when the velocity field was sufficiently sampled, the case of NGC 1090) we called our method MET (Modified Envelope Tracing method). These approaches were tested by constructing artificial data cubes based on geometrical models of the {H{\\sc i}} disc that were iteratively compared to the observed data cubes. The MET/WAMET method proves to provide a better initial estimate of the rotation curve then more traditional methods, such as the first-moment analysis and the single-Gaussian fitting. Modelling the data cubes results in a powerful way to test and improve the derived rotation curves, as shown e.g. by \\cite{Gea:03}.  This work is structured as follows: in Section 2 we present the sample and optical data that were used in this study, in Section 3 we describe the {H{\\sc i}} observations and the data reduction and Section 4 shows the data analysis; in Section 5 we present the results of data analysis, the new method for extracting the rotation curves is shown in Section 6 and the comparison between the observed and the modelled data cubes is presented in Section 7. In Section 8 we introduce the different mass models that were considered; the results of the mass decompositions are shown in Section 9 and in Section 10 we draw the conclusions concerning the derivation of the rotation curves and the dark matter haloes around the galaxies of our sample.  Throughout this paper we adopted a value for the Hubble parameter $H_0$=75 km s$^{-1}$ Mpc$^{-1}$.   \\begin{table*} \\centering \\begin{footnotesize} \\centering \\caption[]{ Optical and radio resolution scales (indicated respectively as OR and RR), and $r_{f}$, in units of the scale radii of the Burkert and NFW density profiles. \\vspace{0.3cm} \\label{rs}} \\begin{tabular}{l c c c c c c} \\hline \\hline Galaxy       &OR/$r_{core}$     & RR/$r_{core}$ &OR/$r_{s}$  &RR/$r_{s}$   &   $r_{f}$/$r_{core}$  & $r_{f}$/$r_s$     \\\\ \\hline ESO\\,116-G12 & 0.03              & 0.49  & 0.02   & 0.29         & 2.7        & 1.6                    \\\\ ESO\\,287-G13 & 0.01              & 0.18  & 0.02   & 0.29         & 1.0        & 1.7                    \\\\ ESO\\,79-G14  & 0.04              & 0.43  & 0.02   & 0.24         & 2.1        & 1.2                    \\\\ NGC\\,1090    & 0.04              & 0.32  & 0.03   & 0.23         & 3.4        & 2.5                    \\\\ NGC\\,7339    & 0.02              & 0.15  & 0.01   & 0.08         & 0.6        & 0.3                    \\\\ \\hline \\end{tabular} \\end{footnotesize} \\end{table*} ",
        "conclusions": "  \\begin{itemize} \\item[$\\bullet$]{Cored haloes uniquely fit the rotation curves, with core radii of order $r_{opt}$ and central densities of order $10^{-24}$ g cm$^{-3}$.}  \\item[$\\bullet$]{NFW haloes cannot represent the dark matter haloes around the galaxies of our sample: they fit badly and a significant number of data points is inconsistent with the fits.}  \\item[$\\bullet$]{The situation is even worse for the Moore model, as implied by its steeper inner profile.}  \\item[$\\bullet$] { There are cases in which for radii in the range $(1-4)r_d$ the Burkert and NFW halo rotation curves are similar; this may explain why in the literature cases can be found (e.g. \\citealt{J:03}) in which both models fit equally well a rotation curve of small extension and limited resolution.}  \\end{itemize}  We thus reach the conclusion that the galaxies of the present sample are uniquely successfully fitted by cored haloes, with a core size comparable to the optical radius. This suggests the existence of a well-defined scale length in dark matter haloes, linked to the luminous matter, which is totally unexpected in the framework of CDM theory (see \\citealt{SB:00})."
    },
    "0403/astro-ph0403648_arXiv.txt": {
        "abstract": "Thanks  to its extraordinary brightness, the X-ray afterglow of GRB030329 could be studied by {\\it XMM-Newton} up to two months after the prompt Gamma-ray emission.  We present the results of two {\\it XMM-Newton} observations performed on May 5 and 29, as well as an analysis of the {\\it Rossi-XTE} data of the early part of the afterglow, discussing in particular the stability of the X--ray spectrum and presenting upper limits on the presence of X--ray emission lines. ",
        "introduction": "GRB030329 is an exceptional Gamma--ray burst for various reasons: it  had a very large fluence of $\\sim$10$^{-4}$ erg cm$^{-2}$ (30--400 keV, \\cite{Ricker:2003}), in the top 1\\% of all observed GRBs; its redshift is  z=0.1685 (\\cite{Greiner:2003},\\cite{Caldwell:2003}), which makes it the second nearest GRB; its optical transient was observed at magnitude 13 one hour after the explosion (\\cite{Price:2003},\\cite{Torii:2003}); it is the first GRB unambiguously associated with a supernova (\\cite{Stanek:2003}, \\cite{Hjorth:2003}). The detailed studies in all wavebands made possible by the brightness of this event are yielding an unprecedented understanding on the jet structure, GRB energetics and circumburst environment. In particular, the X-ray afterglow could be studied at late times, with a sensitivity which was not achieved for previous bursts.  The spectral shape and the time evolution of the X--ray afterglow has been already reported and discussed in comparison with preliminary measurements of the optical afterglow by \\cite{Tiengo:2003}. ",
        "conclusions": ""
    },
    "0403/astro-ph0403362_arXiv.txt": {
        "abstract": "We reexamine arguments advanced by Hayashi \\& Matsuda (2001), who claim that several simple, physically motivated derivations based on mean free path theory for calculating the viscous torque in a quasi-Keplerian accretion disk yield results that are inconsistent with the generally accepted model. If correct, the ideas proposed by Hayashi \\& Matsuda would radically alter our understanding of the nature of the angular momentum transport in the disk, which is a central feature of accretion disk theory. However, in this paper we point out several fallacies in their arguments and show that there indeed exists a simple derivation based on mean free path theory that yields an expression for the viscous torque that is proportional to the radial derivative of the angular velocity in the accretion disk, as expected. The derivation is based on the analysis of the epicyclic motion of gas parcels in adjacent eddies in the disk. ",
        "introduction": "Accretion disks around a central compact object are ubiquitous in astrophysics. In such systems, the plasma accreting onto the central object typically takes the form of a swirling disk. The disk-like nature of accretion is due to the angular momentum originally carried by the accreting material, and also to the spin of the central object, which provides a preferred axis of rotation. At a given radius in a quasi-Keplerian accretion disk, the average motion of the gas parcels is an approximately circular Keplerian orbits, combined with a radial sinking motion as the parcel gradually moves deeper into the potential well of the central mass. The radial velocity that is far smaller than the Keplerian (azimuthal) velocity. A detailed description of quasi-Keplerian accretion disks can be found in Pringle (1981), among several other references.  The specific angular momentum (angular momentum per unit mass) carried by a parcel of gas in a quasi-Keplerian accretion disk is very close to the Keplerian value $\\sqrt{G M R}$, where $R$ is the radius measured from the central object with mass $M$. It follows that a parcel of plasma has to lose angular momentum in order to sink from a larger orbit into a smaller one and eventually cross the event horizon. In the absence of strong winds or jets, the angular momentum must therefore flow {\\it outwards} in such an accretion disk in order to enable accretion to proceed. In disks of turbulent fluid, the flow of angular momentum is due to the exchange of plasma parcels in neighboring annuli with different angular velocities. This fluid picture may describe certain astrophysical cases such as the rings of Saturn, but in many other cases of astrophysical interest, the coupling between adjacent annuli is provided not by ``parcel interchange,'' but rather by the magnetic field. In either the fluid or magnetohydrodynamical (MHD) scenarios, the angular momentum flow is expressed as a torque between neighboring annuli. Our focus here is on the fluid picture, and on the validity of the various attempts to analyze quantitatively the angular momentum transport associated with the interchange of parcels.  There are several previous papers in which derivations of the viscous torque in fluid dynamical situations are developed based on simple, physically motivated arguments (e.g., Frank, King, \\& Raine 1985, 1992, 2002; Hartmann 1998). However, some of these do not yield an expression for the viscous torque that is proportional to the radial derivative of the angular velocity. This is a problem because the viscosity is fundamentally due to the ``rubbing'' of matter in adjacent radial annuli in the disk, and consequently the viscosity should {\\it vanish} in the case of solid body rotation with $\\Omega(R)=\\rm constant$, where $\\Omega(R)$ is the angular velocity in the disk at radius $R$. Hayashi \\& Matsuda (2001; hereafter HM) recognized this point, and attempted to clear up some of the confusion by carefully examining the previously published derivations of the viscosity that were based on the mean free path approach. They concluded that the mean free path approach inevitably leads to an inward rather than outward flow of angular momentum in the disk, which is unphysical. However, we argue that the reasoning of HM was flawed because they did not consider the epicyclic nature of the parcel trajectories. We present a simple derivation based on mean free path theory, combined with the actual epicyclic motion of the gas parcels, that in fact yields a physically reasonable expression for the viscous torque between neighboring annuli in a quasi-Keplerian accretion disk.  We discuss the standard fluid dynamical formulation of viscous torques in an accretion disk in \\S~2. We then discuss the problems with the various previous derivations that attempted to utilize a mean free path approach to compute the viscous torque in \\S~3. In \\S~4 we present a simple physical derivation based on analysis of the epicyclic (ballistic) motion of gas parcels in adjacent eddies, and we demonstrate that this approach yields the expected form for the torque in terms of the gradient of the angular velocity. We present our final conclusions in \\S~5. ",
        "conclusions": "The derivation presented in \\S~4 clearly employs mean free path theory, since we assumed that the fluid parcels travel an average radial distance $\\lambda$ before being damped by the surrounding gas. We have thus shown that there does exist a simple derivation, based on mean free path theory, that yields an expression for the viscous torque ${\\cal G}$ (eq.~[\\ref{eq26}]) that is directly proportional to the radial derivative of the angular velocity, $d\\Omega/dR$, in agreement with our physical expectation. This expression for ${\\cal G}$ can be used in equation~(\\ref{eq4}) to proceed further in deriving the structure of the accretion disk, as in Pringle (1981). Our results provide a simple but important unification of the ``parcel interchange'' viscosity model with the MHD viscosity model, in which the coupling that transports the angular momentum is provided by the magnetic field rather than by fluid turbulence (Frank, King, \\& Raine 2002). In the MHD model, the torque is found to be proportional to the gradient of the angular velocity, $d\\Omega/dR$, in agreement with the results we have obtained here by applying the mean free path approach to the case of fluid turbulence."
    },
    "0403/astro-ph0403681_arXiv.txt": {
        "abstract": "{Multi-frequency radio polarimetric observations of the diffuse Galactic synchrotron background enable us to study the structure of the diffuse ionized gas via rotation measure maps. However, depolarization will introduce artifacts in the resulting rotation measure, most notably in the form of narrow, elongated ``depolarization canals''. We use numerical models of a non-emitting Faraday rotating medium to study the RM distribution needed to create depolarization canals by depolarization due to a finite beam width, and to estimate the influence of this depolarization mechanism on the determination of RM. We argue that the depolarization canals indeed can be caused by beam depolarization, which in turn is a natural consequence when observing a turbulent medium with limited resolution. Furthermore, we estimate that beam depolarization can induce an additional error of about 20\\% in RM determinations, and considerably less in regions that are not affected by depolarization canals. ",
        "introduction": "The Galactic magnetic field and turbulence play a major r\\^{o}le in structuring the Galactic disk. The recognition of the dynamical importance of turbulence in the interstellar medium (ISM) (e.g. Armstrong et al. \\cite{ARS1995}) let our picture of the ISM evolve from a simple three-phase gas to a highly complex medium. Turbulence is believed to play a crucial r\\^{o}le for molecular cloud and star formation (see e.g. Larson \\cite{LAR1981}; Mac Low \\& Klessen \\cite{MAK2003}), however, its origin is widely debated. On a larger scale turbulence leads  to chemical mixing (de Avillez \\& Mac Low \\cite{DAM2002}) and contributes to the vertical pressure balance in the Galactic disk and to the heating of the ISM (Minter \\& Balser \\cite{MAB1998}).  Magnetic fields are tightly linked to the turbulent nature of the ISM. Their large-scale components help confining cosmic rays to the Galaxy, and thus affect chemical processes in molecular clouds. They probably contribute to generating turbulence in the disk (Sellwood \\& Balbus \\cite{SEB1999}) and definitely affect the vertical disk structure. The ratio between mass and magnetic flux in the ISM is an essential piece of information necessary to understand the r\\^{o}le of magnetic fields for molecular cloud and eventually star formation (see e.g. Crutcher \\cite{CRU1999}).  Density and magnetic field structure in the ISM can be probed using multi-frequency radio polarimetry to study Faraday rotation\\footnote{Faraday rotation is the rotation of the plane of linear polarization of radiation when it propagates through a magneto-ionic medium, due to the birefringence of the medium to circularly polarized radiation. The rotation of the polarization plane through the medium is $\\Delta\\phi = RM \\lambda^2$, with rotation measure $RM = 0.81\\int B_{\\parallel}n_eds$, where $B_{\\parallel}$ is the component of the magnetic field parallel to the line of sight in $\\mu$G, $n_e$ is the thermal electron density in cm$^{-3}$ and $ds$ is the path length in pc. } in the magneto-ionic medium. Polarized extragalactic  point sources and pulsars have been used to determine rotation measures (RM), but these give only measurements at certain lines of sight. Observing the diffuse Galactic synchrotron background instead yields structure in polarized emission on many scales, along many contiguous lines of sight. However, depolarization complicates the interpretation of the resulting RM maps. Depolarization occurs if the telescope beam is larger than the scale of structure in polarization (beam depolarization), along the line of sight if the radiation is emitted and Faraday rotated largely in the same medium (depth depolarization or internal Faraday dispersion), or if the frequency bandwidth is so large that the polarization angle changes significantly within the band (bandwidth depolarization), see e.g. Burn (\\cite{BUR1966}) or Sokoloff et al.\\ (\\cite{SBS1998}). In polarization observations of the Galactic synchrotron background, often it is difficult if not impossible to estimate the separate contributions of beam and depth depolarization.  Depolarized regions often exhibit canal-like features (e.g Duncan et al.\\ \\cite{DHJ1997}, \\cite{DRR1999}, Gray et al.\\ \\cite{GLD1999}, Uyan\\i ker et al.\\ \\cite{UFR1999}, Haverkorn et al.\\ \\cite{HKB2000}, Gaensler et al.\\ \\cite{GDM2001}). Both beam and depth depolarization can cause the canals, but require a different underlying RM distribution. The goal of this paper is to determine if the observed depolarization canals can be due to beam depolarization, and how beam depolarization affects the computed RM. To this effect, we use numerical models of the warm magneto-ionic ISM to compare to observations of the diffuse Galactic synchrotron background. The models are irradiated with uniform synchrotron background emission, which is Faraday rotated in the modeled medium. Then, the polarized radiation is smoothed to simulate a finite telescope beam, where after RM's are computed from the smoothed emission maps. The structure in polarized intensity and RM emerging from the smoothed modeled radiation field is then compared to the observations. This allows us to simulate purely the effect that beam depolarization has on polarimetric observations of the ionized ISM.  This work was motivated by the observations briefly presented in Sect.~\\ref{s:obs}. In Sect.~\\ref{s:nummod} we describe the numerical models, the addition of the synchrotron background and the simulation of beam depolarization. Section~\\ref{s:beamdepol} discusses the signature of beam depolarization in the maps of polarized intensity, while in  Sect.~\\ref{s:error} the effect of beam depolarization on RM is studied.  \\begin{figure} \\centering \\includegraphics[width=0.5\\textwidth]{figure1.ps} \\caption{Observed polarized intensity at 349~MHz in a field centered at $(l,b) = (161\\degr, 16\\degr)$. High intensity is shown white, up to a maximum of $\\sim$~13~K.} \\label{f:obs_p} \\end{figure} \\begin{figure} \\centering \\includegraphics[width=0.5\\textwidth]{figure2.ps} \\caption{Map of RM in the observed field. The grey scale runs from $-15$ to 15~rad~m$^{-2}$. The map shows an RM value at every position, regardless of the quality of the linear fit of $\\phi$ against $\\lambda^2$.} \\label{f:obs_rm} \\end{figure} ",
        "conclusions": "\\label{s:conc}  Multi-frequency radio observations of the Galactic synchrotron background allow us to probe the structure of the magneto-ionized component of the ISM from RM maps. However, the quality of the RM determinations depends strongly on the degree of local depolarization due to intrinsic structure in the ISM, and thus in the RM. We investigated the effect of beam depolarization on the quality of the fitted RMs with the help of numerical models. We find: \\begin{enumerate} \\item The models show depolarization canals due to beam depolarization across sharp gradients in RM. This suggests that narrow gradients of a few rad~m$^{-2}$ can be a common feature in a magneto-ionized medium. Combined with other observational evidence, we therefore conclude that the depolarization canals seen in our 350~MHz observations are caused by beam depolarization due to RM gradients. While the width of the canals is an observational artifact, their presence and length mirror underlying structure in the Faraday-rotating medium. \\item RM determinations close to canals can be incorrect by several rad~m$^{-2}$, depending on the absolute RM difference between both edges of the canal. This is a consequence of a non-linearity in $\\phi(\\lambda^2)$ introduced by finite beam width. \\item RM determinations in smooth regions or in regions separated from canals by at least a beam are accurate to $0.05\\Delta\\RM$, where $\\Delta\\RM$ is the width of the RM distribution, or $\\delta\\mbox{RM} = 0.15$ rad~m$^{-2}$ for the observations mentioned here. \\end{enumerate}  In future research, improved models will enable us to investigate what structures in the ISM actually cause the elongated gradients of RM, and will serve as a calibration what structure information we can extract from observational maps."
    },
    "0403/astro-ph0403044_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:intro}  In astrophysics in general, and cosmology in particular, any given object of interest can have many important length and time scales.  An excellent example of this is the process of galaxy formation.  When studying the assembly of galaxies in a cosmological context, one wants to resolve a large enough volume of the universe to capture enough large-scale structure (a box with length on the order of several megaparsecs\\footnote{1 parsec = 3.26 light years = $3.0857 \\times 10^{18} cm$}). However, in order to adequately resolve structure in an individual galaxy one wants to have resolution two orders of magnitude smaller than the ultimate size of the objects of interest (a dwarf galaxy is on the order of $\\sim$ 1 kiloparsec).  This typical cosmological problem requires roughly five orders of magnitude of dynamical range, which is prohibitively expensive when done using a single grid.  Many other astrophysical phenomena, such as the study of molecular clouds and the formation of stars and galaxy clusters, require similarly large dynamical range.  Many scientists have adopted Lagrangean techniques such as smoothed particle hydrodynamics (SPH)\\cite{Monaghan92} to address these issues.  However, this type of method suffers from several drawbacks, including poor shock resolution and fixed mass resolution in regions of interest.  The use of grid-based techniques with structured adaptive mesh refinement avoids many of these problems, and additionally allows the use of higher-order hydrodynamics schemes.  In this paper we present \\enzo, an MPI-parallel 3D Eulerian adaptive mesh refinement code. Though it was originally designed to study cosmological structure formation, the code is extensible and can be used for a wide range of astrophysical phenomena. For more information on the performance and optimization of the \\enzo\\ code, see the contribution by James Bordner in this volume.  The \\enzo\\ web page, which contains documentation and the source code, can be found at \\url{http://cosmos.ucsd.edu/enzo/}. ",
        "conclusions": "\\label{sec:summary}  In this paper we have presented \\enzo\\, a cosmology code which combines collisionless N-body particle dynamics with a hydrodynamics package based on the piecewise parabolic method, all within a block-based adaptive mesh refinement algorithm.  Several other physics packages are implemented, including multiple models for gas cooling and ionization, a uniform ultraviolet background model for gas heating, and a prescription for star formation and feedback.  \\enzo\\ is being released to the public as a community astrophysical simulation code.  This code is being modified and documented to be as widely useful as possible, and can be found at \\url{http://cosmos.ucsd.edu/enzo/}.  Active development is taking place, centering on the addition of magnetohydrodynamics and a diffusive radiative transfer algorithm.  Further information concerning the performance of the \\enzo\\ code (including a package of performance monitoring and visualization tools) is described in a contribution by James Bordner in this volume, and on the \\enzo\\ website.  \\newpage"
    },
    "0403/astro-ph0403272_arXiv.txt": {
        "abstract": "We present evidence for the presence of a transition in the accretion properties of radio-loud sources. For a sample of radio galaxies and radio--loud quasars, selected based on their extended radio properties, the accretion rate is estimated from the black hole mass and nuclear luminosity. The inferred distribution is bimodal, with paucity of sources at accretion rates, in Eddington units, of order $\\sim 10^{-2}$ - assuming a radiative efficiency of 10 per cent - and possibly spanning one-two orders of magnitude.  Selection biases are unlikely to be responsible for such behavior; we discuss possible physical explanations, including a fast transition to low accretion rates, a change in the accretion mode/actual accretion rate/radiative efficiency, the lack of stable disc solutions at intermediate accretion rates or the inefficiency of the jet formation processes in geometrically thin flows. This transition might be analogous to spectral states (and jet) transitions in black hole binary systems. ",
        "introduction": "Radio-loud active galactic nuclei (AGNs) are characterized by the presence of relativistic jets, whose energetics at least in some sources is comparable to that produced by the nuclear engine.  Why and how some systems are able to form powerful jets remains one of the unsolved issues in AGN research, and currently limits our understanding of the relation between the active nucleus, its host galaxy and the larger scale environment. It is possible that all AGNs pass through a jet phase (evolutionary interpretation) or that only certain objects develop powerful jets. For further progress to be made, it is crucial to determine if and how the nuclear activity, in terms of accretion properties onto the central black hole (BH) and/or the black hole spin, is related to the presence and formation of jets. MHD processes appear to be most likely involved in the extraction (and collimation) of energy to form jets from the accretion disc \\citep{blandpay82} and/or black hole spin \\citep{blandzna77}. A significant amount of theoretical work has been carried out recently in order to understand the effectiveness of these processes -- which is largely related to the strength and topology of the magnetic field -- in relation to the properties/regime of accretion (e.g. \\citealt{ghosh97}; \\citealt{livio99}; \\citealt{balbushawley02}; \\citealt{meier02}; \\citealt{merloni02}; \\citealt{livio03}).  A further relevant element tightly related to this issue is the evidence that jets appear in two rather neatly distinct flavors, the more powerful ones exhibiting FR~II radio morphology, the weaker ones FR~I morphology \\citep{fanarof74}. At the zeroth order, the unified model for radio--loud AGNs identifies these two typologies of radio galaxies (RGs) with the parent populations (i.e. the misaligned counterparts with respect to the jet direction) of radio-loud quasars (RLQs) and BL~Lac objects respectively, whose observed properties are dominated by the highly anisotropic emission from relativistic emitting jets \\citep{blandrees78}. This association is supported by the similar isotropic properties as well as the relative number density of such (parent and beamed) populations (\\citealt{urry95} for a review).  For the high power systems the lack of broad emission lines in FR~II RGs has been interpreted as due to the presence of material obscuring the line of sight at large angles with respect to the jet axis (torus).  The existence of an obscuring torus has found direct support from the detection of scattered broad emission lines and from infrared studies which have revealed quasar nuclei hidden by rest frame visual extinction in several FR~IIs (e.g.  \\citealt{antonucci84}; \\citealt{hill96}).   Although this scenario is successful at the zeroth order, issues are still open.  On a rather robust observational ground, it has emerged that the fraction of broad lined objects increases with the source luminosity, at least at low luminosities, and thus - in flux limited samples - with redshift (however this appears not to be the case when weak narrow line objects are excluded, \\citealt{willott01}). Furthermore the quasar narrow line ([OIII]) luminosity and submillimeter emission are larger than those of radio galaxies (e.g. \\citealt{willott02a} for a recent overview).  The two main interpretations of such findings invoke either a dependence of the torus opening angle on source luminosity (`receding torus model', \\citealt{lawrence91}; \\citealt{hill96}) or the presence of a population of - low luminosity - FR~IIs with low levels of nuclear accretion emission (\\citealt{hine79}; \\citealt{laing94}; \\citealt{baum95}; \\citealt{jackson99}; \\citealt{willott00}; \\citealt{chiab}; \\citealt{grimes03}). Another puzzling aspect is the recent suggestion that there might exist a so far undetected population of QSOs with FR~I radio morphology \\citep{blundrawl01}.  Furthermore, despite of the clue offered by the findings by \\citet{ledlow96} that the radio power at the transition between FR~Is and FR~IIs is correlated to the host galaxy magnitude, it is still unclear what causes the FR~I-FR~II dichotomy. Possibilities include the different ways in which a jet interacts with the ambient medium (\\citealt{bick95}; \\citealt{gkwi01}) and/or the different nuclear intrinsic properties of the accretion and jet formation and the jet content (\\citealt{baum95}; \\citealt{meier99}; \\citealt{rey96}).  While little observational evidence has been gathered in support of any of these interpretations, the use of correlations between the global large scale and nuclear properties for the sample analyzed by \\citet{ledlow96} intriguingly suggests that a fundamental parameter determining the FR~I-FR~II dichotomy could be the mass accretion rate in Eddington units, which possibly regulates the accretion mode and also the characteristics of nuclear outflows (\\citealt{ghisellinicelotti2001}).  Therefore, while the environment could play a role in determining the radio morphology and the jet appearance, it is of fundamental importance to determine any connection between the accretion properties and formation of jets.  In this context the aim of our work is to shed light on the process of formation and evolution of jets with respect to the nuclear AGN properties. In particular we intend to assess the role of the central black hole mass ($M_{\\rm BH}$) and the mass accretion rate (in Eddington units, $\\dot{m}\\equiv \\dot M / \\dot M_{\\rm Edd} $) in determining the nuclear and extended properties for a sample of radio-loud objects (RGs and RLQs) as complete as possible.  Clearly, $\\dot{m}$ relies on an estimate of $\\dot{M}$, which will be (necessarily) inferred from the bolometric accretion luminosity of the AGN, $L_{\\rm bol}$. The unknown radiative efficiency of the accreting flow will constitute a crucial parameter in interpreting our findings.  The outline of the paper is the following.  In Section~2 we introduce the considered samples of sources; in Section~3 we discuss the analysis and assumptions adopted to obtain the central black hole mass, the bolometric luminosity and the mass accretion rate in Eddington units.  The results and their discussion are presented in Section~4 and 5, respectively, while in Section~6 we summarize our conclusions. For comparison with previous work in this field a zero cosmological constant universe with $H_{0}$=75 km s$^{-1}$ Mpc$^{-1}$ and $q_{0}$=0.5 is adopted [Assuming a $\\Lambda$CDM cosmology ($\\Omega_{\\Lambda} =0.7, \\Omega_{\\rm M} = 0.3$) does not alter our results, and quantities can be easily re-estimated from Table~1].  \\section[]{Sample selection}  Our goal is to investigate the nuclear properties of a well defined sample of radio--loud objects, whose selection is unbiased with respect to the nuclear properties themselves. Our starting point are the 298 RGs and the 53 RLQs from the 3CR catalogue \\citep{spinrad85}.  For each source, we must be able to estimate $M_{\\rm BH}$ (from the stellar luminosity of the host galaxy, $M_{\\rm bulge}$, see \\S 3.1), as well as the bolometric luminosity of the active nucleus.  While $M_{\\rm bulge}$ is easily measured for RGs, in RLQs the stellar light is often overwhelmed by the non-thermal contribution of the AGN. From the complete sample of 53 3CR RLQs we have therefore selected the sub-sample for which the host galaxy is clearly resolved \\citep{lehnert99}.  This includes 31 objects with 0.3$\\la z \\la$2.0 and total radio luminosities at 178~MHz in the range 10$^{34}$ $< L_{178}<$ 10$^{36}$ erg~s$^{-1}$~Hz$^{-1}$.\\footnotemark \\footnotetext{Twelve of the 3CR RLQs were not observed by Lehnert et al. because of scheduling problems.}  On the other hand, while $L_{\\rm bol}$ is relatively easy to estimate (modulo an appropriate bolometric correction) for RLQs, as the emission from the central AGN is bright and easily measurable, the situation is more delicate for RGs. To obtain the best possible constraints on $L_{\\rm bol}$, we select the 79 3CR radio galaxies for which high resolution Hubble Space Telescope (HST) images led to either a measurement or an upper limit of the optical unresolved nuclear luminosity\\footnote{3CR~386 was included in the Chiaberge et al.'s study. However, ground based observations subsequently confirmed the putative nucleus to be a foreground star (Marchesini, Ferrarese \\& Celotti, in prep). We will exclude the galaxy from further consideration.}  (Chiaberge, Capetti \\& Celotti 1999, 2000, 2002). This sample includes 80\\% of the 3CR RGs with $z \\leq 0.3$. Of these, 26 are FR~Is (10$^{31}$ $< L_{178}<$ 10$^{35}$ erg~s$^{-1}$~Hz$^{-1}$) and 53 FR~IIs (10$^{32}$ $< L_{178}<$ 10$^{35}$ erg~s$^{-1}$~Hz$^{-1}$).  While not complete, this sample is homogeneous (all galaxies belong to the radio-flux limited 3CR sample) and unbiased with respect to nuclear properties (see discussion in \\S5).  We will consider this our main sample. We also note the significant overlapping of the sources both in $L_{178}$ and $z$ (see Fig.~\\ref{fig4}).  In order to assess the incidence of possible selection biases, enlarge the statistics, and include RLQs with redshifts overlapping with the redshift range of the RGs (thus reducing the redshift-luminosity degeneracy and any evolutionary effect, see Fig.~\\ref{fig4}), we also consider three additional samples. The first comprises the 28 low luminosity RGs (mostly FR~I) from the B2 sample (B2RG) with published high resolution HST images \\citep{B2_02}. These galaxies populate the low redshift range $z < 0.15$ and span 10$^{29}$ $< L_{178}<$ 10$^{33}$ erg~s$^{-1}$~Hz$^{-1}$. The second sample includes the 12 RLQs from \\citet{McLureDunlop2001} (MDRLQ). This is a rather heterogeneous sample with redshift 0.15$\\la z \\la$0.3; the common thread being the availability of HST images for all objects. The third sample includes three high redshift (1.9$\\la z \\la$2.4) RLQs imaged with HST by \\citet{hutch2002}.  The total radio luminosity-redshift plane for all of the sources considered in this work is shown in Fig.~\\ref{fig4}.  \\begin{figure} \\center \\includegraphics[width=8cm]{figure4.ps} \\caption{Total radio luminosity $L_{178}$ at 178~MHz versus redshift. Circles are FR~Is (empty ones are B2RGs while filled ones are 3CR~RGs), squares are FR~IIs and triangles are RLQs (empty ones are 3CR~RLQs; light filled ones are MDRLQs and the dark filled ones are the three RLQs from \\citealt{hutch2002}).} \\label{fig4} \\end{figure} ",
        "conclusions": "We considered the nuclear properties of radio-loud objects, comprising a complete sample of radio galaxies (both FR~I and FR~II) and radio-loud quasars. We estimated the mass of the central supermassive black hole, the nuclear accretion luminosities, and the corresponding mass accretion rates.  We found that while BH masses span a relatively large range, they are clustered around $10^8-10^9$ M$_{\\odot}$ and there is no systematic difference in their values between FR~I, FR~II radio galaxies and radio-loud quasars, in agreement with previous findings (e.g. \\citealt{odowd02}). Biases introduced by the (radio) criteria adopted in defining the samples do not allow us to address whether a lower limit on the BH mass or an upper envelope in radio luminosity for a given mass exist in radio-loud sources (\\citealt{McLureDunlop2002}; \\citealt{laor00}; \\citealt{dunl03}; but see \\citealt{ho2002}; \\citealt{oshlack2002}; Woo \\& Urry 2002a, 2002b).  However, the distributions in bolometric luminosity, and (even more so) in the corresponding mass accretion rate appear to be bimodal. FR~I, low excitation FR~II and some high excitation FR~II galaxies occupy the low end of the $L_{\\rm bol}$ and $\\dot{m}$ distributions: $10^{41}\\la L_{bol} \\la 10^{43.5}$ erg s$^{-1}$ or $10^{-6}\\la \\dot{m} \\la$ few $10^{-4}$. On the other hand, radio-loud quasars, broad line radio galaxies and some high excitation FR~II galaxies occupy the high end of the distributions: $10^{45}\\la L_{\\rm bol} \\la 10^{47}$ erg s$^{-1}$ or $10^{-2}\\la \\dot{m} \\la$ 1. The region in between is characterized by a marked deficiency of sources, especially if high excitation galaxies are indeed intrinsically similar to broad line ones.  The relative completeness and homogeneity of our sample, which includes low and high power sources selected on the basis of the extended radio properties, allowed to reveal such bimodal behavior, which had not been noticed in previous studies (\\citealt{odowd02}; \\citealt{woo02a}). This is so pronounced to appear despite of the uncertainties in the inferred quantities. No clear selection bias which could induce such result has been found. In a future work, we will test this bimodal behavior through more accurate determination of the BH masses (via the $M_{\\rm BH}-\\sigma$ relation) and luminosities for a smaller subsample of sources including a representative sample of FR~II sources (Marchesini, Ferrarese \\& Celotti, in preparation). The range of $\\dot m$ corresponding to the transition/gap between FR~Is and FR~IIs ($\\dot{m}\\sim 10^{-2}-10^{-1}$ for $\\epsilon =0.1$) is similar to the value of $\\dot m \\sim$0.06 corresponding to the dividing line between FR~I and FR~II radio galaxies discussed by \\citet{ghisellinicelotti2001}, which was inferred via an independent procedure based on the properties of the \\citet{ledlow96} RG sample via global correlations between the extended radio emission and the accretion one. In particular the magnitudes of the host galaxies reported by \\citet{ledlow96} were converted into BH masses (again via the $M_{\\rm bulge} - M_{\\rm BH}$ relation), while the radio luminosities were associated to nuclear accretion rates, by adopting the correlation between radio power and narrow line luminosity \\citep{willott99}, and the latter converted into ionization luminosity. The absence of a gap in the findings by \\citet{ghisellinicelotti2001} might however arise as the relation by \\citet{willott99} was (necessarily) adopted by these authors also for FR~Is, which tends in fact to overestimate the FR~I line luminosity (as also shown by \\citealt{wills04}).  A dichotomy is not seen if the radio luminosity is substituted to $L_{\\rm bol}$ (in this case the distribution appears continuous with significant overlap between the different classes), making our findings even more intriguing.  A discontinuity in the $\\dot{m}$ distribution implies a transition/gap in the accretion which is not reflected in the extended radio emission and possibly, if a similar jet radiative efficiency would characterize FR~Is and FR~IIs, in the jet power. This would also imply a significantly different ratio between the accretion and jet powers in low and high luminosity systems, questioning whether the $\\dot m$ taken at face value in low power sources is sufficient to give raise to the associated jets. We also notice that the radio luminosity corresponding to the transition in $L_{\\rm opt}$ between LEGs and BLRGs -- according to the established radio-line correlations (e.g. \\citealt{willott99}) -- reflect the break in the radio sources luminosity function (e.g. \\citealt{dunlpeak}).  Our findings indicate that the whole radio-loud population can be characterized by thermal vs non-thermal dominated (optical) emission (as suggested by Chiaberge et al. 2002).  The population of FR~II radio galaxies is in this respect complex and - as expected - the source behavior is related to the nuclear line properties.  In other words, the nuclear structure and the central activity define the parent radio galaxies of RLQs and BL~Lacs and such activity is not uniquely connected to the extended radio power and morphology. This might suggest different time-scales for the evolution of the radio and optical activity (see \\citealt{ciras03}) and/or a temporal evolution from the higher to the lower power sources.  Some possibilities accounting for the found bimodal behavior have been discussed.  These involve a relatively rapid transition in the accretion rate; a transition in the accretion mode, implying a change in radiative efficiency and/or the onset of significant mass loss, with transition values of $\\dot{m}$ in agreement with the rates at which radiative inefficient flows are expected to be possible; the presence of an unstable (e.g.  radiation pressure dominated) accretion regime or an inefficient jet formation process associated to geometrically thin flows at intermediate $\\dot m$. While no definite conclusion can be drawn on the origin of the found transition/gap, some of these hypothesis might be (dis)-favored by examining with more accurate measures the effective absence of sources over a significant range at intermediate $\\dot m$, i.e. the presence of a significant gap in the observed properties.  A further aspect which we are exploring is whether the dichotomic behavior holds also for radio-quiet sources, which would imply that the last of the above possibilities cannot be the (only) explanation.  Several pieces of evidence already indicated that in low power radio galaxies and BL~Lac objects the luminosity resulting from dissipation in the accretion flow is extremely low, compared to that typically inferred from RLQ activity (as e.g. from the low line luminosity, limits on the nuclear emission, e.g. \\citealt{urry95}; \\citealt{celotti98}; \\citealt{odowd02}; \\citealt{wang02}). Furthermore, stringent limits came with the finding that the radiated output is orders of magnitude lower than that expected from the (Bondi) accretion rate estimated from the properties of hot gas at the accretion radius for a 10 per cent efficiency (\\citealt{fabian95}; e.g. M87, \\citealt{dimatteo03}; IC4296, \\citealt{pellegrini03}).  Our findings, consistent with the above results, appear thus to indicate the presence of a relatively rapid transition and a gap between the high radiative efficient accretion regime and the accretion properties of low power sources.  Clearly, these results are also consistent with the possibility, suggested by \\citet{dimatteo03}, that the low disc accretion rate might have a time dependent (on-off) character (even between low activity and almost `inactive' states).  Note that if indeed a large fraction of the radio, X--ray and optical emission in some low power objects is dominated by non-thermal radiation from the base of the jet (as proposed), this might imply $\\dot{m}$ lower than $\\sim 10^{-5}-10^{-6}$. The findings from the line emission however indicate that at least in a relevant number of sources the non-thermal jet and 'disc' emission should be of similar order.  Finally, it is worth mentioning that the found behavior could indeed be analogous to the transitions between spectral states present in black hole X-ray binary systems and microquasars, possibly in terms of the value of the transitional $\\dot m$, the nuclear spectral properties and the onset and nature of the jet production (e.g. \\citealt{fender}).  As already amply stressed in the literature in relation to the understanding of accretion at low rates, there is great need to test for the presence of reflection features in the X-ray spectra of such systems, as well as variability and polarization measures, as it has already been possible for the Galactic Center (\\citealt{baganoff01}; \\citealt{aitken00}). Theoretical and numerical work might however lead to a deeper understanding first."
    },
    "0403/astro-ph0403428_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "None of the models was able to exactly reproduce the optical polarization data on the Crab pulsar. The general features, however, (ie.~fast changes of position angle and minima in polarization degree associated with two peaks) find qualitative explanation within the two-pole caustic model."
    },
    "0403/astro-ph0403237.txt": {
        "abstract": "{ % We present a study of the early build-up of dust in high redshift galaxies. The study is based on the analysis of  38  Damped  Ly$\\alpha$ systems (DLAs) for which we derive the fraction of iron atoms in dust form, $f_\\mathrm{Fe}$. The sample is representative of metal-poor galaxies in the redshift range $ 0.6 \\leq z \\leq 3.4$ selected on the basis of their absorption \\ion{H}{i} column density ($N(\\ion{H}{i})  \\geq  2 \\times 10^{20}$ atoms cm$^{-2}$). We find that the dust fraction increases with metallicity,   from  $f_\\mathrm{Fe} \\simeq 0$ at [Fe/H]$\\sim -2$ dex, up to $f_\\mathrm{Fe} \\simeq 0.9$ at solar metallicity; the increase is fast below [Fe/H] $\\simeq -1$ dex and mild at higher metallicities. We also find some evidence for an increase of $f_\\mathrm{Fe}$ with cosmic time; a large fraction of the systems  younger than $\\approx 3$ Gyr has $f_\\mathrm{Fe} \\la 0.5$. % These results indicate the dust-to-metal ratio increases in the course of chemical evolution, at variance with the hypothesis of an approximately constant dust-to-metal ratio, commonly adopted in models of galactic evolution. This hypothesis is consistent with local and high-redshift data only  when the   metallicity is relatively high  ([Fe/H]$\\ga -1$ dex). The  results of this work suggest that the main mechanisms of dust formation  may be rather sensitive to the level of metallicity attained by a galaxy in the course of its chemical evolution. A metallicity-dependent dust production by SNe II seems to be the most promising mechanism for explaining the rise of $f_\\mathrm{Fe}$ at [Fe/H] $\\la -1$ dex. ",
        "introduction": "Interstellar dust plays a key role in a variety of astrophysical processes relevant to galactic evolution and affects, with its presence, many measurements of basic physical properties of stars and galaxies. % % For these reasons, understanding the process of dust formation and evolution on a galactic scale is of great interest in several areas of astrophysical research. % In the course of galactic evolution the abundance of dust is expected to track that of the metals synthesized and ejected  from  the stars. Therefore, the dust-to-gas ratio, $\\cal D$,  is expected to correlate   with the metallicity, $\\cal Z$ or, in other words, the fraction of heavy elements in dust form, $f_\\mathrm{d} \\approx {\\cal D/Z}$, is expected to stay approximately constant  (e.g. Franco \\& Cox 1986). % In the local universe the existence of a   $\\cal{Z - D }$ correlation is supported by empirical evidence. % For instance, Issa et al. (1990) found a  good %$\\cal{Z - D }$ correlation %${\\cal Z \\propto D}^{0.85\\pm0.1}$ for  6 spirals of the Local Group and the Magellanic Clouds, %from the analysis of %metallicities estimated in \\ion{H}{ii} regions and estimating $\\cal D$ from the visual extinction per hydrogen atom. %, i.e. %${\\cal D} \\simeq A_V/N(\\ion{H}{i})$. Schmidt \\& Boller (1993) found %${\\cal Z \\propto D}^{0.63\\pm0.25}$, with a similar correlation for 23 local dwarf galaxies estimating $\\cal D$ %metallicities estimated from oxygen abundances and from  IRAS dust mass and \\ion{H}{i} integrated-emission mass. From a similar analysis Lisenfeld \\& Ferrara (1998)   found a good %$\\cal{Z - D }$ correlation for 28 dwarf irregular galaxies, but not for 16 blue compact dwarfs; they attribute these differences % and attributed this behaviour to variations of the star-formation   and mass-outflow rates between these different types of galaxies.  The empirical $\\cal{Z - D }$ correlation found in the local universe is, to date, the main observational input to models of galactic evolution which incorporate the dust component. %has been {\\em assumed} in several models of galactic evolution %(e.g. Franco \\& Cox 1986 [check], Edmunds \\& Eales 1998 [check]). %Silva et al. 1998 [check]). A self-consistent treatment of the metal and dust content in the models %of galactic  chemical evolution requires taking into account the processes of grain formation, accretion and destruction (e.g. Dwek \\& Scalo 1980), increasing in such a way the number of the %model parameters with few, if any, observational constraints. The detailed treatment of all these processes can be performed using numerical models. Following this approach, and requiring that accretion and destruction time-scales evolve in a similar fashion, Dwek (1998)  found that the interstellar ${\\cal D/Z}$ mass ratio is  approximately constant in the Milky Way. %independent of Galactocentric radius. % More elementary, analytical models have been recently proposed  to reproduce the dust cycle in galaxies (e.g. Hirashita 1999, Edmunds 2001, Inoue 2003). An approximately constant ${\\cal D/Z}$ mass ratio is generally predicted by these  studies, with the exception of Inoue (2003), who proposes models with increasing  ${\\cal D/Z}$    in the course of evolution. The observed linear  $\\cal{Z - D }$ relation in the local universe is interpreted by Inoue as a sequence of constant galactic age. % Measuring the dust-to-metal ratios in galaxies caught in their early stages of evolution is fundamental to probe the validity of these models. This can be done by studying the dust content of high-redshift galaxies, a field of research that is receiving a special attention given its potential impact on our understanding of the early universe (e.g. Silva et al. 1998, Hirashita \\& Ferrara 2002, Morgan \\& Edmunds 2003).  Quasar absorption-line systems %with associated metal lines can be used to investigate the relation between dust   and metallicity at high redshift on observational grounds.  This is true for the Damped Ly $\\alpha$ (DLA) systems, the class of absorbers with the highest \\ion{H}{I} column density ($N(\\ion{H}{I}) > 10^{20.3}$ atoms cm$^{-2}$), which are believed to originate in galaxies  located along the quasar line of sight (Wolfe et al. 1986). % Measuring the metallicity and the dust abundance of DLA systems makes it possible  to probe the  ${\\cal D - Z}$ relation in primeval galaxies. The metallicity  is determined from  column-density measurements, which can be rather accurate  (e.g. Lu et al. 1996, Prochaska \\& Wolfe 1999, Molaro et al. 2000, Ledoux et al. 2003). Element abundances do not suffer from severe ionization corrections in DLA systems (e.g. Vladilo et al. 2001), and the metallicity level can be estimated pretty well, at least for elements not affected by dust depletion (Pettini et al. 1999). Measuring the  dust abundance, however, is less straightforward.  In principle, the amount of dust   in DLA systems  can be estimated from the reddening  of the  background quasar. This technique, however, cannot be applied to individual systems given the uncertain, variable continuum of the quasar. As a consequence, quasar extinction has only been estimated  statistically, by  comparing samples with and without foreground DLAs (Pei et al. 1991). The dust-to-gas ratio in DLA systems estimated in this way is between 5\\% and 20\\% of the Galactic value, in broad agreement with the low level of metallicity of DLA systems. However, this type of analysis does not give information on the evolution of the dust abundance with the metallicity. % Also the dust emission properties cannot be employed to derive the dust abundance in DLA systems. In fact, attempts to identify the emission of intervening galaxies in the quasar field have been successful only for a dozen of absorbers at $z \\la 1$ (Le Brun et al. 1997, Turnshek et al. 2001; see also Kanekar \\& Chengalur 2003), but not in the redshift interval where DLAs are most commonly identified and investigated ($z \\ga 2$).  So far, the only method to estimate the dust abundance of individual DLAs relies on the comparison of the elemental abundances of refractory and non-refractory elements. Since the pioneer investigation on Cr and Zn by Pettini et al. (1994), the evidence for differential depletion of refractory elements relative to the volatile element Zn has accumulated over the years and is now well established  (Hou et al. 2001, Prochaska \\& Wolfe 2002). An additional argument supporting the connection between depletion and dust is the correlation between depletion and fraction of molecular hydrogen (Levshakov et al. 2000, Ledoux et al. 2003).  The existence of a correlation between differential depletion and metallicity, reported by Ledoux et al. (2003), suggests that the dust content of DLA systems   evolves in the course of chemical evolution. However, in order to study the evolutionary properties of the dust we need to convert the differential depletions into dust-to-gas or dust-to-metal ratios. % %Studies of the differential depletions taken at face value %commonly neglect the effects of nucleosynthesis and %of  Zn depletion. In a previous work we presented a method for estimating these ratios in DLA systems, %from the differential depletion, taking also into account the depletion of Zn self-consistently (Vladilo 1998; Paper I). In that work we derived a $\\cal{Z - D }$ correlation  between the metallicity corrected for dust effects and the dust-to-gas ratio. These two quantities, however, are not independent when the observed metallicity is very low. % (see Eq. 20 in paper I). A much safer approach consists in studying the evolution of the {\\em dust-to-metal ratio}, rather than that of the {\\em dust-to-gas ratio}. The former  is derived from the column densities of two metals and is   less affected by error propagation than the latter, which requires, in addition, the use of the \\ion{H}{i} column density.  A systematic study of the evolution of the dust-to-metal ratio in DLA systems is still missing. This type of study requires converting the differential depletions measured at face value into the {\\em fraction of metals in dust form}, $f_d$, taking into account (i) variations of the dust composition that may occur in different types of galactic environments (e.g. galaxies with different metallicities) and (ii) nucleosynthetic effects that may affect the volatile/refratory abundance ratio (e.g. evolution of the Zn/Fe ratio). In a recent work we presented a methodology aimed at tackling these problems (Vladilo 2002a; Paper II). This methodology was successfully applied to recover (i) the (stellar) abundance pattern of the Small Magellanic Cloud starting from interstellar measurements affected by dust depletion; (ii) the $\\alpha$/Fe ratio in DLA systems corrected for dust effects (Vladilo 2002b; Paper III).  In the present paper we   apply this  new method to derive the fraction of iron in dust form, $f_\\mathrm{Fe}$, for   the   sample of DLA systems with measurements of the Zn/Fe ratio. This sample now includes 38 absorbers, more than twice the sample of Paper I. We use the results of the most recent studies of Zn abundance in metal-poor stars % (Mishenina et al. 2002, Gratton et al. 2003, Nissen et al. 2003) to derive an educated guess for the intrinsic Zn/Fe ratio in DLA systems which, at variance with Paper I, is now a free parameter of the method. % Finally, for the first time in this type of investigation,  the robustness of the results is tested against possible bias   introduced by the observations or by the method itself. % The paper is organized as follows. In the next Section we present the method of analysis and the sample. In Section 3 we present the results and test their robustness. In Section 4 we discuss the implications of the results in terms of formation and evolution of dust in primeval galaxies. The conclusions are summarized in Section 5. ",
        "conclusions": "We have investigated % the relation between dust-to-metal ratio % in DLA systems % as a function of metallicity and cosmic time with the aim of % casting light on the evolutionary properties of the dust-to-metal ratio in DLA systems with the aim of casting light on the early build-up of dust in galaxies. The results that we obtain are representative of %indicative of the behaviour of dust in high-redshift, metal-poor galaxies selected on the basis of their \\ion{H}{i} absorption column density, i.e. $N(\\ion{H}{i}) \\geq 2 \\times 10^{20}$ atoms cm$^{-2}$.  As an estimator of the dust-to-metal ratio in DLA systems we have used the fraction of iron atoms in dust form, $f_{\\mathrm Fe}$. One advantage of this approach is that iron is a refractory element which traces the presence of dust even in harsh interstellar environments. The dust fraction  $f_{\\mathrm Fe}$ has been derived by  comparing   the iron abundance  with the abundance of the volatile element Zn. As in Paper I,   we take into account the fact that also Zn can be incorporated in dust form. However, the   methodology adopted here features important improvements over the one followed in Paper I: (1) the intrinsic Zn/Fe  ratio in DLA systems is now a free input parameter; (2) the dust chemical composition is allowed to vary in different galactic environments (see Paper II); (3) we prefer  the dust-to-metal ratio, rather than the less reliable dust-to-gas ratio, to study the evolution of the dust content. % % Also the possible dependence of dust depletions on physical parameters has % been investigated. % The sample analysed includes 39 DLA systems with median % redshift $z \\simeq 2.???$ and The main results  of the present   study can be summarized as follows.  The frequency distribution of iron depletions, $\\delta({\\mathrm Fe}) = \\log (1-f_{\\mathrm Fe})$, peaks at $\\delta({\\mathrm Fe}) \\sim 0$, and decreases regularly with increasing $|\\delta({\\mathrm Fe})|$. {\\em A significant fraction of DLA systems has depletions lower than the lowest values measured in the Galactic ISM}, i.e. those of the Galactic halo gas. The frequently quoted claim that depletions of DLA systems \"are typical of Galactic warm gas\" is correct only for the remaining systems. High values of depletion, typical of Galactic cold disk gas, are not found.  {\\em The dust fraction $f_\\mathrm{Fe}$  increases with metallicity}, albeit with a significant scatter and a % This is in line with the trend between differential depletions % (not corrected for nucleosynthetic effects) and metallicity reported by % Ledoux et al. (2003). However,   we also find  evidence for % a possible {\\em flattening of the trend when [Fe/H] $\\ga -1$ dex}. % and % an increase of the scatter  with % increasing metallicity. An empirical relation between dust fraction and metallicity  of the type $f_\\mathrm{Fe} \\approx f_{\\mathrm{Fe},\\circ} -10^{\\alpha(\\mathrm{[Fe/H]}-\\mathrm{[Fe/H]}_{\\circ})}$, with $f_{\\mathrm{Fe},\\circ} \\simeq 0.9$, $\\alpha \\simeq -1.2$, and $\\mathrm{[Fe/H]}_{\\circ} \\simeq -1.7$, provides a good fit to the data.  % The data  follow  an empirical law of the type %\\begin{equation} %f_\\mathrm{Fe} \\approx %f_{\\mathrm{Fe},\\circ} -10^{\\alpha(\\mathrm{[Fe/H]}-\\mathrm{[Fe/H]}_{\\circ})} %\\label{f_law} %\\end{equation} %with %$f_{\\mathrm{Fe},\\circ} \\simeq 0.9$, %$\\alpha \\simeq -1.2$, and %$\\mathrm{[Fe/H]}_{\\circ} \\simeq -1.7$ %  The comparison between the dust fraction  $f_\\mathrm{Fe}$ and the look-back time, $t_l$ (derived from the absorption redshift), reveals a weak evidence for a rise of $f_\\mathrm{Fe}$ with cosmic time. In spite of a large scatter, the data are consistent with a fast rise of $f_\\mathrm{Fe}$ starting at about 2 Gyr after the Big Bang and lasting a few billion years. Indeed, the comparison of the frequency distributions of depletions for the two sub-samples with  $t_l \\geq 10.5$ Gyr and $t_l < 10.5$ Gyr suggests that the dust fraction does increase with cosmic time.  The above conclusions are rather stable for different choices of the input parameters required for the determination of $f_\\mathrm{Fe}$. In particular, they are not affected by changes of the adopted guess of the intrinsic Zn/Fe ratio in DLA systems,  [Zn/Fe]$_\\mathrm{g}$, in line with the measurements of the same ratio in metal-poor stars. The large amount of measurements recently reported by three independent groups consistently indicate that the   Zn/Fe ratio in metal-poor stars is very close to solar and probably constant in the metallicity interval typical of DLA systems. In any case,  {\\em assuming that the intrinsic Zn/Fe ratio decreases with chemical evolution}, as reported by some authors, {\\em the existence of the trends with metallicity and cosmic time would even be reinforced}.   We have tested the robustness of the results in light of the selection bias which are known to affect the measurements of the \\ion{H}{i}, \\ion{Fe}{ii}, and \\ion{Zn}{ii} column densities. From the analysis of a mock sample of DLA systems with assigned values of dust fraction and metallicity we conclude that the bias effects are small and therefore that {\\em the trend between dust fraction and metallicity represents a genuine property of DLA systems. This is also true for the distribution of depletions, which is peaked at low values and shows a deficiency of systems with depletions typical of Galactic cold disk gas. } The dust obscuration bias may in part be responsible for such deficiency, but probably only for the systems of high metallicity.  The presence of different interstellar phases with a range of physical parameters in DLA systems may contribute to the scatter of the dust fractions measured at each metallicity. %As a first step to interpret the present results we considered a dependence %of the dust abundance on the physical state of the gas. However, the depletions  do not show correlations with physical parameters measured in DLA systems, such as the spin temperature % (Kanekar \\& Chengalur 2003) or the  \\ion{C}{ii}/\\ion{C}{i} ionization ratio. %(Liszt 2002). % or the molecular fraction %(Ledoux et al. 2003). We propose that {\\em the evolutionary trends of the dust fraction in DLA systems  are  mostly governed by the chemical evolution} of the galaxies where the absorptions originate, {\\em rather than by the evolution of their  physical properties}.  The extremely low dust fraction of iron  measured at [Fe/H]$ \\lsim -2$ dex suggests that dust formation is very inefficient during the earliest stages of galactic chemical evolution. We considered the possibility that the main mechanisms of dust formation  in galaxies may have a metallicity-dependent efficiency. Indeed, previous   work already indicates that the dust formation efficiency in the ejecta of SNe II and in the winds of late-type giants may increase with metallicity. We argue that, in addition, the injection of iron-rich dust by type SNe Ia may yield a rise of $f_\\mathrm{Fe}$ in the course of evolution. The metallicity-dependent dust formation in SN II seems to be the most promising candidate for explaining the rise of $f_\\mathrm{Fe}$ at [Fe/H] $\\la -1$ dex.  In order to reproduce the observed evolution of the dust content of DLA systems a simultaneous treatment of the main processes of dust formation, accretion and destruction is required. The observed increase of $f_\\mathrm{Fe}$  with metallicity is contrary to the expectations of current chemical evolution models, which postulate or infer an approximately constant dust-to-metal ratio, with the exception of the simple, analytical  models by Inoue (2003), which predict an increase of the dust-to-metal ratio $\\cal{D/Z}$, starting from very low values. % The present results may represent a unique guideline for incorporating    the dust component in well-tested models, % of galactic chemical evolution, particularly for the early stages of evolution. The successful implementation of the observed trends between $f_\\mathrm{Fe}$, metallicity and cosmic time in realistic models will provide fresh clues  on the formation redshift and nature of DLA galaxies."
    },
    "0403/astro-ph0403102_arXiv.txt": {
        "abstract": "Much recent discussion about dark matter has been centered on two seemingly independent problems: the abundance of substructure in dark matter halos, and the cuspiness of the halos' inner density profile. We explore possible connections between the two problems by studying the gravitational scattering effects due to subhalos on the phase-space distribution of dark matter particles in the main halos. Our series of controlled numerical experiments indicates that the number and mass density of subhalos can be high enough to cause the collisionless dark matter particles in the inner part of a main halo to diffuse, flattening the main halo's inner cusp within a few dynamical times.  Depending on the masses and concentration of the subhalos, the inner density profile of the whole system (main plus sub halos) can either steepen or flatten.  Subhalo accretion can therefore introduce significant scatter in the inner density profiles of dark matter halos, offering a possible explanation for the range of profiles seen in both observations and cosmological simulations. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/hep-ph0403092_arXiv.txt": {
        "abstract": "\\hspace*{\\parindent} We describe string production after angled brane inflation. First, we point out that there was a discrepancy in previous discussions. The expected tension of the cosmic string calculated from the four-dimensional effective Lagrangian did not match the one obtained in the brane analysis. In the previous analysis, the cosmic string is assumed to correspond to the lower-dimensional daughter brane, which wraps the same compactified space as the original mother brane. In this case, however, the tension of the daughter brane cannot depend on the angle ($\\theta$). On the other hand, from the analysis of the effective Lagrangian for tachyon condensation, it is easy to see that the tension of the cosmic string must be proportional to $\\theta$, when $\\theta \\ll 1$. This is an obvious discrepancy that must be explained by consideration of the explicit brane dynamics. In this paper, we will solve this problem by introducing a simple idea. We calculate the tension of the string in the two cases, which matches precisely. The cosmological constraint for angled inflation is relaxed, because the expected tension of the cosmic string becomes smaller than the one obtained in previous arguments, by a factor of $\\theta$. ",
        "introduction": "String theory is perhaps the most promising scenario where quantum gravity is included by the requirement of additional dimensions and supersymmetry. Models with more than four dimensions are interesting, since the idea of the large extra dimension\\cite{Extra_1} may solve or weaken the hierarchy problem.\\footnote{Denoting the volume of the $n$-dimensional compact space by $V_n$, the observed Planck mass is obtained by the relation $M_p^2=M^{n+2}_{*}V_n$, where $M_{*}$ denotes the fundamental scale of gravity.} In scenarios of large extra dimensions, the standard model fields are expected to be localized on a wall-like structure, while the graviton propagates in the bulk. In the context of string theory, the most natural embedding of this picture would be realized by a brane construction. These models are of course very interesting from the cosmological viewpoint. Inflation with such a situation is still an interesting problem\\cite{low_inflation, matsuda_nontach}. A model of defect inflation with large extra dimension is discussed in \\cite{matsuda_defectinfla}. Other cosmological issues,  in which defects play important roles, such as baryogenesis with a low fundamental scale are discussed in \\cite{low_baryo, Defect-baryo-largeextra, Defect-baryo-4D}. We think cosmological considerations with nonstatic brane configurations, such as brane defects or brane Q balls\\cite{BraneQball, MQCD_defect_matsuda} are very important, since we expect that future cosmological observations will reveal the cosmological evolution of the brane Universe. If one wants to know what kinds of brane defects are allowed to affect our present Universe, one should know how they were formed. In the original brane inflation scenario\\cite{brane-inflation0}, the inflationary expansion is driven by the potential between D-branes and anti-D-branes evolving in the bulk space of the compactified dimensions. Then the brane inflationary scenario for branes at a fixed angle is studied in ref.\\cite{angled-inflation}, where the slow-roll condition is improved by a small angle. The end of brane inflation is induced by the brane collision where brane annihilation (or recombination) proceeds through tachyon condensation\\cite{tachyon0}. During brane inflation, the tachyon is trapped in the false vacuum, which may result in the formation of lower-dimensional branes after brane inflation. The production of cosmological defects during this process is discussed in ref.\\cite{angled-inflation, angled-defect}, where it is concluded that cosmic strings are copiously produced in this scenario, but the formation of domain walls is seriously suppressed. Ref.\\cite{Majumdar_Davis}, however, indicates that all kinds of defects would be produced and the conventional problems of cosmic domain walls and monopoles should arise. Later, in refs.\\cite{D-brane-strings, Halyo, BDKP-FI}, brane production was reexamined and the conclusions were different from \\cite{angled-defect} and \\cite{Majumdar_Davis}. \\footnote{The actual phenomenological possibilities of cosmological brane defects in the brane world models are not confined within the above arguments. For example, in ref.\\cite{Defect-baryo-largeextra}, the enhancements of baryon number violation and baryogenesis that are mediated by cosmological defects are discussed for models with localized matter fields. Moreover, in ref.\\cite{Alice-string} and later in refs.\\cite{MQCD_defect_matsuda, incidental}, it is argued that the fields that parameterize the position of the branes can fluctuate in four-dimensional space-time, and the fluctuation forms a cosmological defect. In ref.\\cite{Alice-string}, the position of a brane in the fifth dimension is used to parameterize the Alice string in the effective four-dimensional space-time, and it was shown that the brane must smear in the core so that it resolves the anticipated singularity. Then in refs.\\cite{MQCD_defect_matsuda} and \\cite{incidental}, it is shown that the relative position of the branes in higher-dimensional bulk space could fluctuate in the four-dimensional space-time and could parameterize cosmological defects such as strings, domain walls, and Q balls.}  In this paper, we will examine the tension of the cosmic strings  produced after angled brane inflation. Here we summarize the naive expectations and known facts. \\begin{itemize} \\item The model of angled brane inflation is expected to be a model of D-term inflation in the four-dimensional effective Lagrangian. \\item A cosmic string that is produced after D-term inflation is examined in ref.\\cite{D-term-string}, and is shown to satisfy the BPS conditions in a generic situation. The exact form of the tension of the D-term string is also calculated. \\end{itemize} Thus, one would expect that the cosmic string that is produced after angled brane inflation is the D-term string that satisfies the BPS condition. If the above expectations are correct, the tension of the cosmic string should depend on the angle. On the other hand, however, the previous arguments on string formation after angled brane inflation did not support this expectation. For example, in refs.\\cite{angled-defect, D-brane-strings}, it is discussed that; \\begin{itemize} \\item The produced daughter brane must wrap around the same compactified space as the original mother brane. Thus, the calculated tension of the daughter brane does not depend on the angle. \\end{itemize} This result does not support the above expectation. Moreover, one can calculate the tension of the cosmic string from the effective Lagrangian for tachyon condensation. The tension obtained from the tachyon effective Lagrangian does depend on the angle, which again does not support the argument in refs.\\cite{angled-defect, D-brane-strings}.  In this paper, we will show that one can easily understand why there was such a discrepancy between the two calculations. Our result is that the tension of the produced cosmic string does depend on the angle, which is different from previous discussions. The final state of the produced daughter brane does not wrap the same compactified space as the original mother brane. However, here we should note that our result does not contradict to the previous arguments. We will consider the details of the recombination process of the original branes, and find that the daughter brane that wraps the same compactified space as the original mother brane would appear as the precedent state of the final brane configuration. The precedent brane configuration is then deformed into the final state, where the daughter brane is extended between recombined branes. The tension of the daughter brane now depends on the angle, which is consistent with the expectation from the effective Lagrangian.  In Sec. 2, we begin with a short review of an angled brane inflation model and the cosmic string production after inflation\\cite{angled-defect}. From the viewpoint of the effective Lagrangian, it is easy to see that the tension of the strings produced  must depend on the angle $\\theta$. In the previous arguments\\cite{angled-defect}, however, it was assumed that the daughter brane wraps the same compactified space as the mother brane. On the basis of this assumption, the tension of the cosmic string was calculated and did not depend on $\\theta$. However, if one pays further attention to the process of recombination of the mother branes, one can easily understand why the final state of the brane produced does not wrap the same compactified space as the mother brane. The tension of the cosmic string does depend on the angle $\\theta$, which is of course consistent with the analysis in the effective Lagrangian. The cosmological constraints are modified and will be relaxed by a factor of $\\theta$. ",
        "conclusions": "We have examined string production after angled brane inflation. First, we pointed out that there has been a discrepancy between the expected tension of the cosmic string calculated from the viewpoint of the four-dimensional effective Lagrangian and the one obtained from the brane tension of the daughter brane. If the cosmic string is a lower-dimensional daughter brane that wraps the same compactified space as the original mother brane, the tension of the string cannot depend on the angle. On the other hand, from the analysis of the effective Lagrangian, it is easy to see that the tension of the cosmic string must be proportional to $\\theta$ for $\\theta \\ll 1$. This is the problem we have solved in this paper. Our idea is shown in fig.\\ref{fig:split}. We have calculated the explicit form of the tension of the cosmic string, and compared the two results.  The cosmological constraint for angled inflation will be relaxed because the expected tension of the cosmic string becomes smaller than the one obtained in the previous arguments by a factor of $\\theta$."
    },
    "0403/astro-ph0403452_arXiv.txt": {
        "abstract": "We present new optical imaging and spectroscopy of the peculiar, edge-on Sbc galaxy UGC~10043. Using the WIYN telescope, we have obtained $B$, $R$, and \\HA+[\\NII] images, together with DensePak integral field spectroscopic measurements of the stellar \\CaII\\ infrared triplet and the H$\\alpha$ and [\\NII] lines from the ionized gas. The imaging observations show that the inner bulge of UGC~10043 ($a\\le$\\as{7}{5}) is elongated perpendicular to the galaxy major axis. At larger $r$, the bulge isophotes twist to become oblate and nearly circular, suggesting the bulge is triaxial. The bulge shows no clear evidence for rotation about either its major or minor axis. The inner, southwestern quadrant of the bulge is girdled by a narrow dust lane parallel to the minor axis; unsharp masking reveals that this minor axis dust lane may be part of an inner polar ring, although we find no unambiguous kinematic evidence of orthogonally rotating material. The stellar disk of UGC~10043 has a rather low optical surface brightness [$\\mu(0)_{R,i}\\sim$23.2 mag arcsec$^{-2}$], a small scale height ($h_{z}=395$~pc for $D=33.4$~Mpc), and a mild integral sign warp. A dusty, inner disk component  that appears tilted relative to the outlying disk is also seen. The \\HA\\ and [\\NII] emission lines in UGC~10043 resolve into multiple velocity components, indicating the presence of a large-scale galactic wind with an outflow velocity of $V_{\\rm out}\\gsim$104~\\kms. \\HA+[\\NII] imaging reaffirms this picture by revealing ionized gas extended to $|z|\\sim$3.5 kpc in the form of a  roughly biconical structure. The [\\NII]/\\HA\\ line intensity ratio increases with increasing distance from the plane, reaching values as high as 1.7. Unlike most galaxies with large-scale winds, UGC~10043 has only a modest global star formation rate ($\\lsim1M_{\\odot}$ yr$^{-1}$), implying the wind is powered by a rather feeble central starburst. We discuss evolutionary scenarios that could account for both the structural complexities of UGC~10043 and its large-scale wind. The most plausible scenarios require a major accretion or merger event at least a few Gyr ago. ",
        "introduction": "} UGC~10043 is a nearby galaxy ($D\\approx$33.4~Mpc; see Table~1) that has been classified as an Sbc spiral seen almost directly edge-on (e.g., de Grijs \\& van der Kruit 1996). However, recent high-resolution {\\it Hubble Space Telescope} ($HST$) images obtained through $F300W$ and $F814W$ filters (mid-ultraviolet and $I$-equivalent bands) by  Windhorst et al. (2002) highlighted some morphological peculiarities of this object, including an extremely thin disk that appears very dim in the mid-ultraviolet, a bulge with prolate isophotes, and a  narrow dust lane running parallel to the minor axis.  Vertical elongation is not a common feature of spiral bulges (e.g., Bertola, Vietri, \\& Zeilinger 1991). Indeed, most bulges are significantly flattened by rotation, and many appear ``disky'', with scale heights only slightly larger than their host disks (e.g., Kormendy 1993). A notable exception is the case of certain early-type spirals where ``orthogonally decoupled bulges'' have been recently discovered (Bertola et al. 1999; Sarzi et al. 2000; Corsini et al. 2003; see also Arnaboldi et al. 1993; Reshetnikov, Fa\\'undez-Abans, \\& de Oliveira-Abans 2002). Another class of galaxies displaying vertically elongated inner components are the polar ring galaxies (e.g., Whitmore et al. 1990), although in this latter case, the ``prolate bulge'' is actually an S0 or similar galaxy, whose rotation is orthogonal to an orbiting ring or disk of stars and gas (see the review by Sparke 2002).  A common trait of both the orthogonally decoupled bulge systems and the polar ring galaxies is that both classes of objects contain a structural component whose angular momentum vector is roughly parallel to the major axis of the host galaxy. This in turn suggests these systems could not have formed from the collapse of a single protogalactic cloud, but instead have undergone a ``second event'' after a precursor of the system was significantly evolved (e.g., Schweizer, Whitmore, \\& Rubin 1983; Bertola et al. 1999; Bournaud \\& Combes 2003). Another group of objects belonging to the class of ``multi-spin galaxies'' are the minor-axis dust lane ellipticals. These are observed to have modest rings of gas and dust orbiting about their short axis (e.g., Hawarden et al. 1981; Bertola 1987; Oosterloo et al. 2002). The peculiar minor axis dust lane of UGC~10043 causes its bulge to bear an interesting resemblance to these latter systems, while its elongated bulge and massive stellar disk suggest a more extreme version of these systems, and a possible link to the polar ring or orthogonally-decoupled bulge galaxies.  Polar rings, orthogonally decoupled bulges, and minor axis dust lanes are all relatively rare, but these systems are nonetheless of considerable interest from a galaxy evolution perspective. For example, the ability to measure the kinematics of a galaxy along two orthogonal axes permits unique constraints on the shape of the potential and of the dark matter halo (e.g., Tohline, Simonson, \\& Caldwell 1982; Sackett \\& Sparke 1990; Sparke 1996; Reshetnikov \\& Sotnikova 1997). Multi-spin galaxies serve as important evidence that there is no single recipe for the formation of a disk+spheroid system, and demonstrate that such galaxies are likely to be built  and modified by a variety of processes---all of which need to be accounted for in a complete picture of galaxy formation. In addition, we now have significant observational evidence that ``second events'' have played a key role in the evolution of nearly all massive galaxies (e.g., Freeman \\& Bland-Hawthorn 2002 and references therein). However, because the conditions required to form a stable, multi-spin galaxy from an accretion or merger event are rather restricted, these systems provide particularly important laboratories for exploring how initial conditions and the properties of the participating galaxies influence the outcome of these episodes. Studies of multi-spin galaxies may therefore provide important clues on the roles accretion and minor mergers played in shaping a much wider array of observed galaxy types (e.g., Bettoni \\& Galletta 1991).  The morphological similarities between UGC~10043 and certain multi-spin galaxies strongly suggest a relationship to these objects, and imply that UGC~10043 has had a rather different formation history from ``prototypical'' Sbc galaxies such as the Milky Way or NGC~891. To investigate this possibility further, we have obtained new optical imaging and spectroscopic observations of UGC~10043. This paper describes the analysis of these data, and uses them as a first step toward building a picture of the evolutionary history of this galaxy. ",
        "conclusions": "} Overall, the structural and kinematic complexities of UGC~10043, including its peculiar, vertically-elongated bulge and the presence of a large-scale wind,  suggest a complex evolutionary history, most likely requiring a major ``second event'' (accretion, infall, and/or merging). While these types of processes  are now believed to play a general  role in galaxy formation, UGC~10043 appears to be part of a special group of disk+bulge systems (that includes the polar ring/disk galaxies and orthogonally-decoupled bulge galaxies) whose formation requires rather specific combinations of conditions (e.g., environment, merger geometry, impact velocity, gas cross-section, dark halo concentration) to account for both the production and maintenance of their current structures. As such, these galaxies uniquely preserve and showcase important clues to their formation histories. And since the parameter space for producing such systems is far more limited than in the more general case, studies of galaxies like these can help to lift some of the degeneracies in our understanding of how a variety of internal and external factors  influence the present-day structures and compositions of  galaxies. In this context,  the unique properties of UGC~10043 make it a particularly important laboratory  for testing various classes of evolutionary models. Drawing on relevant theoretical modelling and discussions from the literature, we now sketch some possible formation scenarios for UGC~10043. We outline the strengths and weaknesses of each scenario in hopes of motivating additional observations and numerical modelling  to further test and refine these ideas.  \\subsubsection{Accretion of a Disk around a Naked Spheroid} There is no obvious mechanism by which a disk galaxy with a vertically-elongated bulge could form either by collapse of a single protogalactic cloud, or by secular evolution of an existing disk (e.g., through development of a bar instability). For these reasons, Bertola \\& Corsini (2002) have suggested that such galaxies are likely to represent cases where disks formed around pre-existing, naked spheroids (see also Binney \\& May 1996). In this picture, the decoupled, inner bulge cores seen in several such systems would be by-products of the disk-building event (see also Arnaboldi et al. 1993).  Although we do not yet know if it has a decoupled inner core, UGC~10043 offers an important constraint on this type of model, since its stellar disk is thinner, more regular, and more extended than the disks of any of the other spirals in which vertically elongated bulges have been identified (compare Figure~\\ref{fig:wiynimages} with the images in e.g., Sarzi et al. 2000; Reshetnikov et al. 2002). The small scale height of the UGC~10043 stellar disk and the relatively long characteristic dynamical time scales of its outer regions imply it is a dynamically old entity that has not undergone any dramatic heating or perturbations within at least the last few Gyrs (e.g., Reshetnikov \\& Combes 1997). This is also suggested by the observed vertical and radial color gradients, which are likely to arise from dynamical processes acting over a significant fraction of a Hubble time.  The idea of ``dressing a naked spheroid'' would seem to account nicely for certain properties of UGC~10043, including the apparent disjointedness of bulge and disk (Section~\\ref{optmorph}), and the similarities of its bulge to an elliptical galaxy (Sections~\\ref{bulgedecom}, \\ref{faberjackson}). In most galaxies, disk-building is also expected to augment an existing bulge (e.g., Kormendy 1993) and may account for the observed bulge color gradient (Section~\\ref{cmap}), or perhaps even its triaxial nature. It can be considered a more extreme version of the mechanism believed responsible for the formation of dust lane ellipticals (e.g., Bertola 1987; Sage \\& Galletta 1993; Oosterloo et al. 2002). Indeed, this type of model has been invoked to explain the origin of the Centaurus~A system (e.g., Malin, Quinn, \\& Graham 1983). However, in stark contrast with UGC~10043, the dusty, minor-axis disk of Centaurus~A is dramatically warped and twisted, presumably the  result of violent shearing and subsequent precession of the accreted disk (Sparke 1996).  A modified version of the above scenario would be the idea of disk growth around a spheroid having a smaller, pre-existing, gas-poor disk. Variations of this latter picture have been invoked as a means to account for the origin of  NGC~5084 (Zeilinger et al. 1992; Carignan et al. 1997), NGC~4550 (Rix et al. 1992), and other examples of S0-like galaxies possessing gas and/or extended disks (Bertola, Buson, \\& Zeilinger 1992; Reshetnikov \\& Sotnikova 1997; Kannappan, Jansen, \\& Barton 2004).  Unfortunately, for the case of UGC~10043,  this picture seems to be faced with some challenges. First, the material accreted onto either a naked spheroid or an S0-like galaxy would have to be predominantly gaseous, since any stellar constituent would be dissipationless and unable to cool into a thin layer. In addition, the gas would need to be located predominantly outside the stellar disk of the donor galaxy in order for it to be readily tidally stripped. Numerical simulations show that gas fractions of up to $\\sim$40\\% may be readily accreted from a donor provided that the gas is sufficiently extended (Bournaud \\& Combes 2003).  However, the neutral gas content of UGC~10043 (\\mhi$=5.9\\times10^{9} M_{\\odot}$) is typical of the  gas contents of giant spiral galaxies; this implies that to produce a galaxy as gas-rich as UGC~10043 would require a rather gas-rich predecessor (\\mhi$\\gsim 5\\times10^{10}M_{\\odot}$ to provide for both the current ISM and stellar content), with a weak enough potential such that its gas could be readily stripped by a galaxy of UGC~10043's mass ($M_{\\rm tot}\\sim10^{11}M_{\\odot}$). The giant low surface brightness spirals might be a suitable class of donor, but these are relatively rare, at least in the local universe (e.g., Matthews et al. 2001 and references therein).  Another challenge for the accretion scenario is that numerical simulations show that only rings or annuli, not complete disks, are directly formed via accretion (e.g., Rix \\& Katz 1991; Reshetnikov \\& Sotnikova 1997; Bournaud \\& Combes 2003). Therefore subsequent perturbations and viscous evolution would be required to produce radial spreading into a bona fide disk. This process is expected to take a minimum of several orbital periods (i.e., at least several Gyr; Rix \\& Katz 1991)---although this is not necessarily problematic, given other dynamical arguments for an old age for the UGC~10043 disk. In addition, Rix \\& Katz (1991) have shown that as an accreted gas disk grows radially, it will also decrease in scale height as it transitions from a pressure-supported to a rotationally-supported regime. This may offer a natural explanation for the observed radial color gradients in UGC~10043. However, in order to produce the thin {\\it stellar} disk observed in UGC~10043 would seem to demand that significant cooling of the gas disk took place before significant star formation ensued. In addition, in order to avoid strong twisting and differential precession over time,  the accreted disk would likely need to be self-gravitating (Sparke 1986). Numerical checks on whether the expected timescales for these various processes could reproduce the observed UGC~10043 disk would be valuable.  An alternate version of the accretion picture is that the UGC~10043 disk was built from multiple dwarf galaxies or subgalactic clumps, much as in the standard hierarchical build-up paradigm for normal disk galaxies. This may offer a natural explanation for the large observed radial color gradient in UGC~10043 (see Section~\\ref{cmap}) as well as its tilted inner disk and outer warp (see Sections~\\ref{optmorph}, \\ref{cmap}). Nonetheless, the idea of hierarchical build-up has at least one serious drawback for the case of UGC~10043---namely, how to  build such a thin disk without excessively heating it (e.g., T\\'oth \\& Ostriker 1992).  Some hope is offered by the recent numerical simulations of Huang \\& Carlberg (1997) and Val\\'azquez \\& White (1999), which have shown that while a merging galaxy on a prograde orbit can cause substantial heating and damage to a disk, an interloper on a retrograde orbit will cause far less damage (owing to the weaker coupling between the infalling material and motion in the existing disk), and instead induce a coherent tilt. Heating effects are further lessened if the accreted satellites have small relative masses, and accretion primarily occurs onto the dark halo. However, the accretion of a number of dwarfs of subgalactic clumps would be needed to build a disk as massive as that of UGC~10043, and it is unlikely that every infall event would have occurred on a retrograde orbit or onto the halo. In addition, if the potentials of ``naked spheroids'' contain two preferred planes (one along the short and one along the long axis), and if this build-up mechanism were effective, we might expect vertically elongated bulges to be observed in a much more significant fraction of present-day Sb-Sc spirals.  \\subsubsection{Capture} A second possibility is effectively the inverse of the above accretion scenario---namely that the bulge of UGC~10043 may be a secondary structure, acquired well after the disk had formed and settled. The culprit might be a small, gas-poor elliptical galaxy plunging in at a very low impact velocity (less than a few tens of \\kms) along its pole. This is analogous to the mechanism proposed for the formation of collisional ring galaxies, including the bulge-dominated collisional ring galaxies known as Hoag-type objects (Appleton \\& Struck-Marcell 1996 and references therein).  Interestingly, numerous pure disk (Sd) galaxies with properties nearly identical to the disk of UGC~10043 are known, and many examples can be found in the Flat Galaxy Catalogue of Karachentsev, Karachentseva, \\& Parnovsky (1993). One example is UGC~825 (Matthews \\& van Driel 2000; see their Figure~7). UGC~825 is strikingly similar to the UGC~10043 disk not only morphologically (having a ``needle-like'' disk with a dusty center), but also in terms of  surface brightness, luminosity, linear size, rotational velocity, and \\HI\\ content. Small, compact ellipticals with luminosities similar to the UGC~10043 bulge are also known, although these seem to be relatively rare in the local universe (e.g., Wirth \\& Gallagher 1984; Nieto \\& Prugniel 1987).  That a pure disk galaxy might be able to ``capture'' a bulge has the qualitative appeal of accounting for the appearance of UGC~10043 as a disk galaxy whose bulge has been ``inserted'' post factum (Figure~\\ref{fig:unsharp}); it  could also simultaneously account for several of the unusual properties of the bulge (large vertical extent relative to the disk, triaxial shape, lack of evidence for significant rotation) that imply it could not have formed secularly. Nonetheless, it remains highly questionable whether this process could work at all in practice, as it would seem to require an extraordinary set of circumstances, including a gas-rich disk and a small, compact elliptical present in the same region of space, both on suitable orbits, and possessing an extremely low relative velocity, rarely observed in real galaxy groups. Initially, the victimized disk would be expected to develop a bowl-like shape and ring structures in the stars and gas---neither of which is evident in UGC~10043---although for a collision of roughly equal-mass galaxies, these signatures are expected to diminish significantly after $\\sim10^{8}$~years. (Gerber, Lamb, \\& Balsara 1996).  Since the bulge and disk of  UGC~10043 share the same radial systemic velocity, one further requirement for a ``capture'' scenario is that the bulge cannot simply be caught in the act of ``passing through''. However, Appleton \\& Struck-Marcell (1996) propose that an intruder might be strongly decelerated via dynamical friction from a massive halo in the target galaxy.  Very thin disks such as UGC~825 are expected to possess massive dark matter halos (e.g., Zasov, Makarov, \\& Mikhailova 1991; Gerritsen \\& de Blok 1999), although numerical simulations are needed to test whether they could supply sufficient dynamical friction to effectively halt an intruder.  \\subsubsection{Origin via a Major Merger} Finally, let us consider the possibility that both the bulge and disk of UGC~10043 formed through a single event involving the merger of two moderate mass disk galaxies. This model has been explored by Bekki (1998) as a means of accounting for the origin of polar ring galaxies whose rings are too massive, too gas-rich, and too radially extensive to be readily explained via mass transfer or capture of a dwarf satellite (see also Iodice et al. 2002; Bournaud \\& Combes 2003). In this picture, a gas-rich ``intruder'' plummets through its ``victim'', leading to the formation of a central S0-like component, a stellar halo, and an orthogonally rotating ring or annulus of gas and stars via a combination of violent relaxation and dissipative processes. To produce a large ($r>$10 kpc) disk also requires the victim to be a gas-rich system (Bournaud \\& Combes 2003).  At least qualitatively, this scenario appears able to account for certain characteristics  of UGC~10043. First, this type of event is expected to drive gas to small radii, where it could trigger a nuclear starburst. In addition, it offers a way of accounting for some of the unusual properties of the bulge (vertically elongated isophotes; lack of evidence for major axis rotation). Through dissipative cooling, it also may be able to produce a gaseous structure with very small scale height. Unfortunately, like the scenarios described above, this one too seems to meet with some shortfalls and the uncomfortable requirement of rather specialized initial conditions. As discussed by Bournaud \\& Combes (2003), to produce a very massive polar disk requires that the victim itself be a very massive, gas-rich disk (with \\mhi\\ roughly twice that of the resultant system) and quite stellar-poor, so as to not produce an observable stellar halo. Thus the victim would need to be something resembling a giant low surface brightness galaxy. An additional constraint is similar to that cited for the ``capture'' scenario---namely that to form a disk rather than a ring, the impact speed of the intruder would need to be less than 30-40~\\kms---a rare condition in typical galaxy groups. Neither of these conditions rules out the merger picture, but it does imply that galaxies like UGC~10043 should be quite rare. Another concern is whether such a violent event could reproduce the complex, triaxial bulge structure of UGC~10043 and the significant color differences between bulge and disk. Once again, numerical work is needed to further explore this possibility.  \\subsection{Origin of the Enhanced Nuclear Star Formation and Galactic Wind\\protect\\label{Wakamatsu}} Independent of the details of how UGC~10043 was assembled, another question surrounding this galaxy concerns how it has evolved since formation, and what has led to the production of its large-scale wind and centrally-concentrated star formation. Although most of the types of formation mechanisms described above are predicted to drive gas to the inner regions of the galaxy (e.g., Reshetnikov \\& Sotnikova 1997; Bournaud \\& Combes 2003), the gas consumption rate required to feed the wind of UGC~10043 ($>1M_{\\odot}$ yr$^{-1}$ within the inner few kpc of the galaxy) suggests that the wind is likely to be a transient and recently-triggered phenomenon. Since the dynamically settled appearance of the disk of UGC~10043 suggests it is a relatively old structure, it is therefore not clear that the event that formed the UGC~10043 system is the same one that instigated the current  nuclear star formation and wind.  One possible solution to  this puzzle may be the mechanism proposed by Wakamatsu (1993). Wakamatsu suggested that if the central component of a polar ring galaxy is sufficiently dense, its oval potential will severely disturb the orbits of the gas at small and intermediate galactocentric radii in the disk, leading to the production of shock waves. He argues that repeated shock waves occurring as the disk rotates would cause loss of orbital energy and angular momentum, and in turn drive gas to the nuclear regions of the galaxy, much in the same manner as an inner bar. While we have shown that UGC~10043 is not a true polar ring (or polar disk) galaxy, its vertically-elongated bulge may nonetheless present a potential similar to the one envisioned by Wakamatsu (1993). Moreover, to match observations of  polar ring galaxies such as NGC~4650A, Wakamatsu assumes that the polar ``ring'' in his model galaxies is actually a disk-like entity, making the overall situation seemingly rather analogous to the UGC~10043 system. Qualitatively his scenario seems to predict certain key traits of UGC~10043: its wind, its dusty inner disk, and possibly the formation of a distinct inner bulge core, although quantitatively the agreement is less good. Typical boundary radii for the shock regions estimated by Wakamatsu are $\\sim4h_{b}$, where $h_{b}$ is the exponential scale length of the S0 disk (or bulge, in our case). For UGC~10043, exponential fits to the outer bulge isophotes yield $h_{b}\\approx$1~kpc, suggesting the shock boundary will extend to roughly 4~kpc. The radial extent of the nuclear star-forming region in UGC~10043 was inferred to be $\\lsim$2~kpc (Section~\\ref{radio}), corresponding rather closely with the observed edge of the bulge, but roughly a factor of two smaller than predicted by Wakamatsu's model. Nonetheless, further investigation of this type of model via numerical modelling would be of interest for exploring the inner gas dynamics of UGC~10043 and other related galaxies.  Another galaxy that may provide a test of link between Wakamatsu's scenario and the formation of galactic winds is the polar ring galaxy NGC~660 (Alton et al. 2000). NGC~660 is classified as a ``superwind'' galaxy, and it is also a relatively rare example of a polar ring-like structure surrounding a gas-rich disk galaxy, rather than an early type host, making it a likely environment for the production of shocks. In addition, the morphology of the superwind galaxy NGC~1482 (Veilleux \\& Rupke 2002) suggests it too may be a related object. To our knowledge, no kinematic or structural decoupling has yet been identified in NGC~1482, but inspection of DSS images reveals hints of a ring or extended disk misaligned with the dust lane of this galaxy. Overall it would seem  that the search for galactic wind signatures in other polar ring galaxies or early-type spirals with  kinematically decoupled cores and gas-rich disks may yield important clues to the origin of both phenomena.  Finally, we note that if the inner polar ring structure described in Section~\\ref{unsharp} is real, its origin is likely to require an additional event independent of the formation of the main galaxy. This would most likely have involved the accretion of a low-mass, gas-rich satellite, and thus serves as another possible source for triggering the nuclear star formation and wind in UGC~10043. For an orbiting inner ring, collisions between gas clouds within the ring and gaseous material in the disk are expected to produce compression to drive molecular gas formation and therefore drive star formation (e.g., Tenorio-Tagle 1981; Galletta 1991). Hot (shock-heated) gas may also be produced in this process (Galletta 1991)."
    },
    "0403/astro-ph0403178_arXiv.txt": {
        "abstract": "{  We present results from the {\\xmm} observations of the dwarf nova YZ Cnc in a quiescent state. We have performed a detailed time series analysis of the resulting light curves. Unusually, we do not detect any orbital modulation in the UV, with only marginal evidence for X-ray modulation on this period. Although there are peaks in the X-ray periodograms at periods less than 5000 s, we attribute them to red noise effects and assign significance to them using a novel approach. The variability in the UV and optical bands can also be modelled as a result of aperiodic variability  (red noise) in the system. There is evidence that  the UV and X-ray fluxes are anti-correlated with a time delay of about 100 s, with the UV lagging behind the X-ray emission. This anti-correlation is intriguing, but is only present on two occasions lasting several 1000 s each. The X-ray spectrum shows  similar emission features to other dwarf novae and is well fitted using a multi-temperature emission model. We measure a relatively high X-ray luminosity of $\\sim1.4\\times 10^{32}$ ergs/s, although this is consistent with a low binary inclination. Finally, we find evidence for a possible -1200km/s blue shift  in the fitted Fe K line energies, possibly indicating the presence of an outflow in this low inclination system. ",
        "introduction": "Dwarf novae are interacting stellar binary systems in which material gets transfered from a red dwarf secondary star onto a white dwarf via Roche lobe overflow. In the absence of a strong magnetic field this material forms an accretion disk around the white dwarf. They show outbursts which occur on week to month timescales. They have been well studied in many energy bands, including the optical and X-ray bands. In X-rays, the emission is suppressed during an outburst while it only recovers at the end of the optical outburst (eg Wheatley et al 1996). This is due to the gas being optically thick during outburst but optically thin in quiescence (eg Popham \\& Narayan 1995).  YZ Cnc is a dwarf nova with an orbital period of 125.17 min (van Paradijs et al 1994). It is also a member of the SU UMa class of dwarf novae - systems which show normal dwarf novae outbursts and also longer and brighter outbursts. UV observations show strong lines with P Cygni profiles which vary in shape over the orbital period - this was interpreted as an asymmetry in the wind flow pattern (Drew \\& Verbunt 1988, Woods et al 1992). White light observations by Pezzuto, Bernacca \\& Stagni (1992) show some evidence for dwarf nova oscillations on a period of 26 s, although this was only seen on one night.  \\begin{figure*} \\begin{center} \\setlength{\\unitlength}{1cm} \\begin{picture}(16,10) \\put(2,9){\\special{psfile=nofold-new.ps hscale=60 vscale=60 hoffset=420 voffset=-270 angle=90}} \\end{picture} \\end{center} \\caption{The light curves of YZ Cnc, from the top; the UV data; 0.15-1.0keV, 1--4keV and 4--10keV. The date is HJD-2452000.0. The bin size is 60 s. The times of anti-correlation between the X-ray and UV bands are shown with horizontal lines in the top panel.} \\label{nofold} \\end{figure*}  The most prominent period seen in dwarf novae in the optical band is the binary orbital period. Other observed periods include the superhump period (which is due to disc precession), dwarf novae oscillations (on timescales of several 10's of s) and quasi-periodic oscillations which have also been seen in X-rays and have modulation periods of $\\sim$100-1000 s (eg Patterson, Robinson \\& Nather 1977, Woudt \\& Warner 2002 and references therein).  Unlike the magnetic systems, direct evidence of the spin period of the white dwarf has been hard to come by.  Recently, observations using {\\xmm} of the dwarf nova OY Car showed the X-rays varied on a timescale of 2240 s (Ramsay et al 2001a): the amplitude of this variation was stronger at soft X-rays. A further investigation also showed a period near 3500 s (Hakala \\& Ramsay 2003). They interpreted the 2240 s period as the spin period of the white dwarf and the 3500 s period as a beat between the spin and the binary orbital period.  This led us to investigate whether other dwarf novae, if they are observed for a sufficiently long time, and with X-ray detectors of sufficiently large area, show similar evidence for variations which could be attributed to the spin period of the white dwarf. We have obtained observations of YZ Cnc using {\\xmm} which cover $\\sim$4.5 orbital cycles. Our prime goal was to search for modulations in the X-ray and UV light curves. Here we present the results of our time series analysis. We also model the X-ray spectra and compare it with previous X-ray observations of dwarf novae. ",
        "conclusions": "Our $\\xmm$ observations of YZ Cnc do not reveal any periodicities in the system, even if clear short term variability is present in both X-ray and UV bands. The cross-correlation analysis reveals that during major UV flares the X-rays and UV are anticorrelated, with UV lagging behind the X-rays by $\\sim$ 100 s. The measured X-ray luminosity, $\\sim 1.4\\times 10^{32}$ \\ergss, is relatively high for a dwarf nova in quiescent state. A multitemperature plasma model (possibly with a partially covering absorber) is required to model the X-ray spectrum. Finally we find that the iron line complex near 6.7 keV shows some evidence for blue shifted emission at about -1200 km/s. However, we cannot entirely rule out instrumental effects."
    },
    "0403/astro-ph0403387_arXiv.txt": {
        "abstract": "From the 389 2002~OGLE-III observations of Galactic Bulge microlensing events we select 321 that are well described by a point-source point-lens lightcurve model. From this sample we identify $n=1$ event, 2002-BLG-055, which we regard as a strong planetary lensing candidate, and another, 2002-BLG-140, which is a possible candidate. If each of the 321 lens stars has 1 planet with a mass ratio $q=m/M=10^{-3}$ and orbit radius $a=R_{\\rm E}$, the Einstein ring radius, analysis of detection efficiencies indicates that 14 planets should have been detectable with $\\Delta\\chi^2 > 25$. Assuming our candidate is due to planetary lensing, then the abundance of planets with $q=10^{-3}$ and $a=R_{\\rm E}$ is $n_p \\approx n/14$ = 7\\%. Conversion to physical units (M$_{\\rm Jup}$, and AU) gives the abundance of `cool Jupiters' ($m\\approx{\\rm M}_{\\rm Jup}$,$a\\approx4$~AU) per lens star as $n_p \\approx n/5.5$ = 18\\%. The detection probability scales roughly with $q$ and $(\\Delta\\chi^2)^{-1/2}$, and drops off from a peak at $a\\approx4$~AU like a Gaussian with a dispersion of 0.4~dex. ",
        "introduction": "Gravitational lensing was first proposed by \\citet{E} -- a consequence of his theory of general relativity, it is the effect of a massive body warping space-time and bending light around it, acting as a lens. This produces two or more images of background objects distorted around the `Einstein ring', which can be thought of as a scale radius ($R_{\\rm E}$) of the lens and depends on its mass, $M$.  The angular radius of the ring is given by: \\begin{equation} {\\theta}_{\\rm E} = \\frac{R_{\\rm E}}{D_L} = \\sqrt{ \\frac{D_{S}-D_L}{D_{S}D_{L}} \\frac{4GM}{c^{2}}}, \\end{equation} where $D_L$ and $D_S$ are the observer -- lens and observer -- source distances respectively.  For objects of stellar mass, this radius is below the resolution limit of our telescopes, and we observe the phenomenon of microlensing.  In this case, we cannot resolve multiple images of the source, but detect an increase in brightness as the individual amplifications combine to produce a brighter stellar image.  As we observe a foreground (lens) star passing in front of a background (source) star, we see an increase in brightness up to a maximum amplitude $A_0$ at the time $t_0$ of closest approach, followed by a decrease.  This gives a symmetrical light curve around $t_0$, providing the timescale for the event is short enough for parallax effects to be negligible.  The timescale depends on the lens transverse velocity, relative to the observer-source line of sight, $v$, and is given by the length of time it takes the source to cross the Einstein ring of the lens: \\begin{equation} t_{\\rm E} = \\frac{2R_{\\rm E} }{v}. \\end{equation}  This depends on both mass and distances, through the dependence on $R_{\\rm E}$.  $v$ is also indirectly influenced by distance, as the positions within the galaxy of the source and lens stars will give them average motions based on the orbital velocity at their galactic radius. For a survey looking toward the galactic bulge (i.e. $(D_S-D_L)/D_SD_L \\sim$ kpc), a microlensing event due to a stellar mass ($M={\\rm M}_{\\odot}$) lens will last a few weeks. A Jupiter mass object will have $t_{\\rm E} \\approx$ 1 day, while an Earth would produce an effect which would be observable for only a few hours.  If a stellar lens has a planetary companion, this will appear as a brief deviation from the stellar microlensing light curve, where the amplifications due to the star and planet combine \\citep{MP,GL,BR}. These deviations should be observable, even for low mass objects, and therefore provide a method to detect extra-solar planets which is potentially more sensitive to small planets than radial velocity searches, given a programme of intensive monitoring with data taken many times per hour.  \\subsection{ The OGLE project }\\label{OGLE}  The main difficulty with using gravitational microlensing to search for exoplanets is the low probability that any given source star will be lensed. This can be overcome by a monitoring programme which studies dense star fields and can detect a flux increase in any of the stars in the field.  The Optical Gravitational Lensing Experiment (OGLE) \\citep{U} monitors $\\sim$ 150 million stars in the galactic bulge, and alerts the astronomical community via Internet and e-mail alerts as soon as a lensing event is seen to begin.  OGLE began in its current form, OGLE-III, in 2001. This version uses Difference Image Analysis (DIA) photometry \\citep{W}, instead of the profile fitting (PSF) method used for the previous three seasons. This has substantially increased the detection rate of stellar microlensing events by OGLE -- there were 167 events discovered during the three years of \\mbox{OGLE-II}, compared to 389 recorded during the 2002 season with \\mbox{OGLE-III}.  We use OGLE 2002 data to put firm upper limits on the population of planets in the bulge of our galaxy. Our method follows that of \\citet{THK}, and is  described in Section \\ref{lcf}, before Section \\ref{2002data} looks at the results found for the 2002 data. The abundance of planets is discussed in Section \\ref{ptot}. ",
        "conclusions": ""
    },
    "0403/astro-ph0403664_arXiv.txt": {
        "abstract": "Understanding of the nature of dark energy, which appears to drive the expansion of the universe, is one of the central problems of physical cosmology today. In an earlier paper [Daly \\& Djorgovski (2003)] we proposed a novel method to determine the expansion rate $E(z)$ and the deceleration parameter $q(z)$ in a largely model-independent way, directly from the data on coordinate distances $y(z)$.  Here we expand this methodology to include measurements of the pressure of dark energy $p(z)$, its normalized energy density fraction $f(z)$, and the equation of state parameter $w(z)$.  We then apply this methodology to a new, combined data set of distances to supernovae and radio galaxies. In evaluating $E(z)$ and $q(z)$, we make only the assumptions that the FRW metric applies, and that the universe is spatially flat (an assumption strongly supported by modern CMBR measurements). The determinations of $E(z)$ and $q(z)$ are independent of any theory of gravity.  For evaluations of $p(z)$, $f(z)$ and $w(z)$, a theory of gravity must be adopted, and General Relativity is assumed here.  No a priori assumptions regarding the properties or redshift evolution of the dark energy are needed. We obtain trends for $y(z)$ and $E(z)$ which are fully consistent with the standard Friedmann-Lemaitre concordance cosmology with $\\Omega_0 = 0.3$ and $\\Lambda_0 = 0.7$. The measured trend for $q(z)$ deviates systematically from the predictions of this model on a $\\sim 1-2 ~\\sigma$ level, but may be consistent for smaller values of $\\Lambda_0$. We confirm our previous result that the universe transitions from acceleration to deceleration at a redshift $z_T \\approx 0.4$. The trends for $p(z)$, $f(z)$, and $w(z)$ are consistent with being constant at least out to $z \\sim 0.3 - 0.5$, and broadly consistent with being constant out to higher redshits, but with large uncertainties. For the present values of these parameters we obtain: $E_0 = 0.97 \\pm 0.03$, $q_0 = -0.35 \\pm 0.15$, $p_0 = -0.6 \\pm 0.15$, $f_0 = -0.62 - (\\Omega_{0} - 0.3) \\pm 0.05$, and $w_0 = -0.9 - \\epsilon (\\Omega_{0} - 0.3) \\pm 0.1$, where $\\Omega_{0}$ is the density parameter for nonrelativistic matter, and $\\epsilon \\approx 1.5 \\pm 0.1$. We note that in the standard Friedmann-Lemaitre models $p_0 = - \\Lambda_0$, and thus we can measure the value of the cosmological constant directly, and obtain results in agreement with other contemporary results. ",
        "introduction": "Observations of supernovae (Riess et al. 1998; Perlmutter et al. 1999; Tonry et al. 2003; Knop et al. 2003; Barris et al. 2004; Riess et al. 2004) indicate that the universe is accelerating in its expansion. Precision measurements of cosmological parameters from CMBR experiments confirm this remarkable finding (e.g., Bennett et al. 2003, Spergel et al. 2003, and references therein). Results similar to those obtained using supernovae are also obtained using radio galaxies (Guerra \\& Daly 1998; Guerra, Daly, \\& Wan 2000; Daly \\& Guerra 2002; Podariu et al. 2003). The acceleration of the univesere at the present epoch is one of the key results of modern cosmology, with potentially significant implications for fundamental physics as well. The nature of the ``dark energy,'' which apparently drives the cosmic acceleration, is unknown and it is crucially important to extract information about it from the data in a manner that is as direct and model-independent as possible.  In Daly \\& Djorgovski (2003; hereafter paper I), we showed how the data could be used to study the dimensionless expansion rate of the universe, $E(z)$, and the deceleration parameter of the universe, $q(z)$ directly from combinations of the first and second derivatives of the coordinate distance.  These determinations only depend upon the Friedmann-Robertson-Walker metric, and an assumption of spatially flat geometry, which is now very well established by the CMBR experiments.  The evaluations do not require the specification of anything else, including  a theory of gravity, and thus are direct and model-independent.  The use of model-independent methods to derive information about the dark energy are also discussed, for example, by Huterer \\& Turner (1999, 2001), Saini et al. (2000), Tegmark (2002), Sahni et al. (2003), Huterer \\& Starkman (2003), Wang and Freese (2004), Wang \\& Tegmark (2004), Wang et al. (2004), and Daly \\& Djorgovski (2004). The work of Huterer \\& Turner focuses on determinations on $w(z)$, as does that of Huterer \\& Starkman (2003). Wang \\& Freese (2004) focus on the determination of the energy density of the dark energy, and use an approach that is complementary to that used here, by integrating over shells in redshift space to obtain the energy density as a function of redshift, while we differentiate the data to obtain this function.  The approach taken by most authors to extract the redshift behavior of the dark energy is to integrate over an assumed functional form of the redshift evolution of the dark energy, having first adopted a theory of gravity (e.g. Starobinsky 1998; Huterer \\& Turner 1999, 2001; Saini et al. 2000; Chiba \\& Nakamura 2000; Maor, Brustein, \\& Steinhardt 2001; Golaith et al. 2001; Wang \\& Garnavich 2001; Astier 2001; Gerke \\& Efstathiou 2002; Weller \\& Albrecht 2002; Padmanabhan \\& Choudhury 2002; Tegmark 2002; Huterer \\& Starkman 2003;  Sahni et al. 2003; Alam et al. 2003; Wang \\& Freese 2004; Wang et al. 2004; Wang \\& Tegmark 2004; Nessier \\& Perivolaropoulos 2004; Gong 2004; Zhu, Fujimoto, \\& He 2004; Elgaroy \\& Multamaki 2004; Huterer \\& Cooray 2004; Alam, Sahni, \\& Starobinsky 2004). However, it can be difficult to extract information about the redshift behavior of the dark energy using these ``integral'' approaches (Maor, Brustein, \\& Steinhardt 2001; and Barger \\& Marfatia 2001). Thus, we continue to follow the complementary approach of differentiating the data, as described in paper I.  Here, the approach presented in paper I is taken a step further, to obtain the pressure, energy density, and equation of state of the dark energy directly from combinations of the first and second derivatives of the coordinate distance with respect to redshift.  This approach is complementary to the standard approach of assuming a theory of gravity, assuming a parameterization for the dark energy and its redshift evolution, and obtaining the best fit model parameters.  We apply this methodology on an improved set of distances to supernovae (SNe) from Riess et al. (2004), supplemented with the data on high-redshift radio galaxies (RGs) from paper I. ",
        "conclusions": "We expanded and used the method developed in paper I on a new sample of coordinate distances to SNe and RGs, to evaluate the trends of the expansion rate $E(z)$, deceleration parameter $q(z)$, pressure of the dark energy $p(z)$, its fractional energy density $f(z)$, and its equation of state parameter $w(z)$, as functions of redshift. We make an assumption that the FRW metric is valid, and the observationally supported assumption of the spatially flat universe. This enables us to derive the trends for $E(z)$ and $q(z)$ which are otherwise model-independent, and thus can help discriminate at least some proposed models of the dark energy.  By assuming that the standard GR is the correct theory of gravity on cosmological scales, we can also produce trends of $p(z)$, $f(z)$ and $w(z)$, without any additional assumptions about the nature of dark energy. These trends may be also used to discriminate between different physical models of the dark energy.  We find that the data are generally, but perhaps not entirely consistent with the standard Friedmann-Lemaitre concordance cosmology with $w=-1$, $\\Omega_0 = 0.3$, and $\\Lambda_0 = 0.7$, although somewhat lower values of $\\Lambda_0$ may be preferred.  We confirm the result paper I and that of Riess et al. (2004), that there is a clear increase $q(z)$ with redshift, with the present value $q_0 = 0.35 \\pm 0.1$, and the transition from decelerating to accelerating universe at $z_T \\approx 0.4$.  Functions $p(z)$, $f(z)$, and $w(z)$ are consistent with being constant at least out to $z \\sim 0.5$, and possibly beyond; the existing data are inadequate to constrain their evolution beyond $z \\sim 0.5$, but there are some hints of increase with redshift for $f(z)$ and $w(z)$.  At lower redshifts, the data are cosistent with cosmological constant models.  We obtain for the present values $w_0 = -0.9 \\pm 0.1$ and $p_0 = -0.6 \\pm 0.15$ ($= - \\Lambda_0$ for the Friedmann-Lemaitre models).  Even with the currently available data, these results represent new observational constraints for models of the dark energy. We believe that this methodology will prove increasingly useful in determining the nature and evolution of the dark energy as better and more extensive data sets become available. Clearly, this method works best when redshift space is densely sampled.  Our current results suggest that redshift space is sufficiently sampled at redshifts less than about 0.4.  More accurate results could be obtained by increasing the sampling of data points with redshifts greater than 0.4, particularly in the redshift range from 0.4 to 1.5."
    },
    "0403/astro-ph0403514_arXiv.txt": {
        "abstract": "High-angular resolution CO observations of small-area molecular structures (SAMS) are presented.  The feature-less structures seen in the single-dish measurements break up into several smaller clumps in the interferometer map. At an adopted distance of 100pc their sizes are of order a few hundred AU, some of which are still unresolved at an angular resolution of about $3''$. The clumps have a fractal structure with a fractal index between 1.7 and 2.0. Their kinetic temperature is between 7~K and 18~K. Adopting standard conversion factors masses are about 1/10 $M_{Jupiter}$\\ for individual clumps and densities are higher than 20000~cm$^{-3}$.  The clumps are highly overpressured and it is unknown what creates or maintains such structures. ",
        "introduction": "Measurements of the 3K microwave background in connection with big-bang nucleosynthesis have impressively shown that most of the matter in the universe is in some unknown non-baryonic form and only 16\\% is baryonic in nature \\citep{spergel2003}, of which only a small part has been detected so far \\citep{turner1999}. According to gravitational lensing experiments less than 25\\% of the unseen baryonic matter in our Galaxy can be in form of massive compact halo objects \\citep{afonso2003}. An interesting alternative is molecular hydrogen \\citep{pfenniger1994,gerhard1996,walker1998}, because at most temperatures in the interstellar medium it cannot be observed directly, but only through secondary tracers such as carbon monoxide, CO. \\citet{pfenniger1994} have proposed that most of the dark matter in the outskirts of our Milky Way could be in form of cold molecular gas with a fractal structure. Basic building blocks in their model are so-called clumpuscules with sizes of about 100AU and Jupiter mass ($10^{-3}$~M$_\\sun$).  Clumpuscules should form in the fragmentation process of an interstellar cloud initiated by efficient cooling. Thus they should exist throughout the Galaxy. In the outer region of the Galaxy they are hard to detect, because due to their low metallicity, thus low CO abundance, and low kinetic temperature they emit only very low intensities, only slightly above the 3K cosmological background. In the inner Galaxy conditions are more favourable to detect them in CO lines, but here it is impossible to isolate individual clumpuscules due to overcrowding in normal molecular gas.  The best location to detect individual clumpuscules is at high galactic latitudes far away from known molecular clouds.  Good candidates for such structures are the small-area molecular structures (SAMS) which were recently detected in a region with very low extinction in the CO $J=2\\to1$ and $1\\to0$ lines \\citep{heithausen2002}.  In the low-angular resolution data they appeared as unresolved clouds with full widths at half maximum below 1 arcmin. Adopting normal conditions one can show that such structures cannot survive very long in the interstellar radiation field.  In this {\\it Letter} I will present high-angular resolution observations of the SAMS, which show that they are composed of tiny molecular clumps, which with respect to their structural properties resemble the molecular clumpuscules proposed by \\citet{pfenniger1994} as bayonic dark matter candidate, however are warmer and probably less massive. ",
        "conclusions": "In recent years molecular hydrogen has been detected in the diffuse interstellar medium directly via absorption-line measurements towards many distant quasars \\citep{jenkins2003,richter2003}. Such observations, which trace mainly warm gas with low column densities, show that H$_2$ is wide spread in the Galaxy even outside star-forming regions. They however provide only little information on the spatial structure of the clouds. The emission-line observations presented here for the first time disclose that molecular clouds in the diffuse interstellar medium are fractally structured. The building blocks of such clouds are low mass clumps or clumpuscules with sizes of order hundred AU.  Due to the high uncertainty in mass and density it is impossible to judge whether these clumps are stabilized by their own gravity. The mass determination is based on conditions as for an average molecular cloud in the galactic plane, which may not hold at such small scales. Further excitation studies of the clumps are required to solve this issue.  It is interesting to note that the clumps have the same radial velocity as the surrounding atomic gas \\citep{heithausen2002} and are within the range of velocities for the molecular gas in the Ursa Major cirrus clouds \\citep{devries1987} which are a few degrees away from the clumps. This means that tiny molecular clumps are possibly a natural constituent of the diffuse interstellar medium, however not recognized as such so far. They can form or survive even in regions with low column densities. From $^{12}$CO and $^{13}$CO observations of molecular gas in translucent clouds the existence of small scale structure down to a few hundred AU has been inferred \\citep{falgarone1998}; the observation presented here indicate that such structures are possibly intrinsically linked to the formation process of molecular clouds.   Whether or not fractally structured clouds similar to the ones described here could account for all the missing baryonic dark matter in our Galaxy is hard to determine. To be consistent with the $\\gamma-$ray background denser and more massive clumps than described in this {\\it Letter} are predicted by models, which explain the missing baryonic mass with molecular clumpuscules \\citep{depaolis1999,kalberla1999}. Due to their small sizes and their narrow line widths both, the hypothetic dark matter clumpuscules and the clumps described here, are notoriously hard to detect unless they form small clusters and thus fill at least some part of the beam of a larger telescope. With small telescopes used to survey the Galaxy they are undetectable due to the low beam filling; high angular resolution and high sensitivity observations are needed for detection. Such clouds thus provide an ideal means to hide matter from the observer."
    },
    "0403/astro-ph0403208_arXiv.txt": {
        "abstract": "We report extensive photometry of frequently outbursting dwarf nova IX Draconis. During five months of observations the star went into three superoutbursts and seven ordinary outbursts. This allowed us to determine its supercycle and cycle lengths as equal to $54\\pm1$ and $3.1\\pm0.1$ days, respectively. During the September 2003 superoutburst, which had the best observational coverage, IX Dra displayed clear superhumps with a period of $P_{sh}=0.066968(17)$ days ($96.43\\pm0.02$ min). This period was constant during the whole superoutburst. Another period, which was clearly present in the light curve of IX Dra in superoutburst, had a value of 0.06646(6) days ($95.70\\pm0.09$ min) and we interpret it as the orbital period of the binary. Thus IX Dra is the first SU UMa star showing orbital modulation during the entire superoutburst. The beat between these two periods is the main cause of an unusual phase reversal of superhumps - a phenomenon which was previously observed in ER UMa. If our interpretation of the second periodicity is correct, IX Dra has an extremely low period excess $\\epsilon$ equal to only $0.76\\% \\pm 0.03$\\%. This implies very low mass ratio $q=0.035\\pm0.003$, which strongly suggests that the system contains a brown dwarf-like degenerate secondary of mass $\\sim0.03 ~\\cal M_\\odot$ and that IX Dra is the most evolved dwarf nova known.  Such a very low mass ratio results in the outer edge of the accretion disk reaching 80\\% of the distance between the components of the system. In turn, this allows the disk particles to enter a 2:1 resonance and leads to the appearance of the orbital period in the light curve of the entire superoutburst.  The high level of activity and brightness of IX Dra indicate that very old cataclysmic variables go through episodes of increased activity leading to loss of angular momentum through mass loss from the system.  Modulations with the orbital period are also detectable during normal outbursts and in quiescence.  \\noindent {\\bf Key words:} Stars: individual: IX Dra -- binaries: close -- novae, cataclysmic variables ",
        "introduction": "Dwarf novae are non-magnetic cataclysmic variables, which are close binary systems containing white dwarf primary and Roche lobe filling secondary. The secondary is typically a low mass main sequence star, which loses its material through the inner Lagrangian point. This material forms an accretion disc around a white dwarf (Warner 1995).  One of the most intriguing classes of dwarf novae are SU UMa stars which have short orbital periods (less than 2.5 hours) and show two types of outbursts: normal outbursts and superoutbursts. Superoutbursts are typically about one magnitude brighter than normal outbursts, occur about ten times less frequently and display characteristic tooth-shape light modulations with a period a few percent longer than the orbital period of the binary.  The behavior of SU UMa stars in now well understood within the frame of the thermal-tidal instability model (see Osaki 1996 for review). Superhumps occur at a period slightly longer than the orbital period of the binary system. They are most probably the result of accretion disc precession caused by gravitational perturbations from the secondary. These perturbations are most effective when disc particles moving in eccentric orbits enter the 3:1 resonance. Then the superhump period is simply the beat period between orbital and precession rate periods.  Over ten years ago, the diversity of behavior observed in cataclysmic variables seemed to be quite small. These variables were classified into few distinct groups which were clearly visible in the $P_{orb}-\\dot M$ diagram (see for example Fig. 3 of Osaki 1996). There is a group of SU UMa stars located below the period gap. At longer orbital periods one can find three other groups: U Gem, Z Cam and nova-like stars, listed according to increasing $\\dot M$.  At the beginning of the 1990s, the situation started to be more complicated. The systems showing superhumps were divided into four subgroups:  \\begin{itemize}  \\item ordinary SU UMa stars,  \\item permanent superhumpers - high accretion rate systems permanently in superoutbursts (Skillman and Patterson 1993),  \\item WZ Sge stars, characterized by an extremely long quiescent state, going into superoutburst every $\\sim$10 years and showing no or infrequent ordinary outbursts,  \\item ER UMa stars - systems characterized by an extremely short supercycle (20-60 days), a short interval between normal outbursts (3-4 days) and small amplitude (3 mag) of superoutbursts (Kato and Kunjaya 1995, Robertson et al. 1995)  \\end{itemize}  There are only five known ER UMa stars: ER UMa itself (Kato and Kunjaya 1995, Robertson et al. 1995), V1159 Ori (Robertson et al. 1995, Patterson et al. 1995), RZ LMi (Robertson et al. 1995, Nogami et al. 1995), DI UMa (Kato et al. 1996) and IX Dra (Ishioka et al. 2001). The most poorly observed object in this group is IX Dra, thus we decided to include it in the list of stars observed within the Curious Variables Experiment (Olech et al. 2003a,b). Here we report the results of a five month observational campaign performed in 2003. ",
        "conclusions": "We presented the results of the 2003 observational campaign of frequently outbursting dwarf nova IX Draconis. The main conclusions of our work are:  \\begin{enumerate}  \\item Analysis of the global light curve spanning five months of observations shows that IX Dra is a very active object and goes into superoutburst every 54 days and into a normal outburst every 3.1 days. According to the model by Osaki (1995), this is due to a very high mass rate of mass transfer from the secondary.  \\item During two best observed superoutburts, we detected the clear superhumps with a period of \\linebreak 0.066968(17) days ($96.43\\pm0.02$ min). The superhump period seems to be constant during entire superoutburst.  \\item During September 2003 superoutburst, a second periodicity was clearly visible in the light curve of the star. Its value was 0.06646(6) days ($95.70\\pm0.09$ min). A wave with the same period was detected during normal outbursts and in quiescence. We suggest that this is the orbital period of the binary system.  \\item The beat between superhump and orbital period is the main cause of an unusual phase reversal of maxima of superhumps - a phenomenon which was previously observed also in ER UMa itself (Kato et al. 2003a). We suggest that careful reanalysis of observations of ER UMa from superoutburst obtained by Kato et al. (2003a) should result in discovery of the orbital period.  \\item The period difference observed in IX Dra corresponds to a mass ratio $q = 0.035\\pm0.003$. Assuming a typical white dwarf mass of $0.6 - 0.8 {\\cal M}_\\odot$, the secondary component must have a mass lower than $0.03 {\\cal M}_\\odot$, which makes it the best candidate among dwarf novae (alongside EG Cnc) for a brown dwarf.  \\item IX Draconis is the first SU UMa star showing orbital modulations during the entire superoutburst. This is possibly due to the extremely low mass ratio which allows the edge of the disc to reach 80\\% of the separation of the binary. In the outer region of the disc, its particles can enter into a 2:1 orbital resonance. This generates a two-armed spiral pattern of tidal dissipation and is responsible for the appearance of the light modulations with orbital period (Osaki and Meyer 2002).  \\item The position of IX Dra in the period excess vs. orbital period plane suggests that the star is the most evolved cataclysmic variable star in the Galaxy, which reached its period minimum long time ago and now evolves towards longer periods.  \\item High activity of IX Dra is in clear contrast with the behavior of its neighbours in the $\\epsilon - \\log P_{orb}$ plane which are mostly faint and inactive objects. Thus we suggest that very old dwarf novae, which most of the time are quiescent and behave like WZ Sge stars, show occasionally greatly increased activity with a high rate of accretion. DI UMa and IX Dra are currently in this state. This high rate of accretion may cause mass loss from the system, supplementing the emission of gravitational waves as the cause of angular momentum loss.  \\end{enumerate}  \\bigskip \\noindent {\\bf Acknowledgments.} ~We acknowledge generous allocation of  the Warsaw Observatory 0.6-m telescope time.  We would like to thank Prof. J\\'ozef Smak for reading and commenting on the manuscript.  {\\small"
    },
    "0403/hep-ph0403167_arXiv.txt": {
        "abstract": "We assess how well next generation neutrinoless double beta decay and normal neutrino beta decay experiments can answer four fundamental questions. 1) If neutrinoless double beta decay searches do not detect a signal, and if the spectrum is known to be inverted hierarchy, can we conclude that neutrinos are Dirac particles? 2) If neutrinoless double beta decay searches are negative and a next generation ordinary beta decay experiment detects the neutrino mass scale, can we conclude that neutrinos are Dirac particles? 3) If neutrinoless double beta decay is observed with a large neutrino mass element, what is the total mass in neutrinos? 4) If neutrinoless double beta decay is observed but next generation beta decay searches for a neutrino mass only set a mass upper limit, can we establish whether the mass hierarchy is normal or inverted? We base our answers on the expected performance of next generation neutrinoless double beta decay experiments and on simulations of the accuracy of calculations of nuclear matrix elements. ",
        "introduction": "\\label{sec:introduction}     A new generation of double beta decay experiments will be undertaken with unprecedented accuracy. In approximately the same time frame,  it will become possible to make much more precise measurements of, or set constraints on, the mass of neutrinos emitted in ordinary beta decay. The results of these next generation experiments will be important for understanding the physics of weak interactions.  If neutrinoless double beta decay is observed, then one can conclude~\\cite{majoranacharacter} immediately that neutrinos are Majorana particles without messing around with detailed calculations and qualifications of the kind discussed in this paper. (We will not consider alternative interpretations such as $R$-parity violation \\cite{Hall:1983id,Mohapatra:su,Mohapatra:1981pm}, which can probably be verified or excluded at high-energy colliders.  The violation of lepton number is clear in either case.)  The community of physicists can and will celebrate if double beta decay is observed.  In this paper, we provide quantitative estimates of how well we can answer four other fundamental questions about neutrinos using the assumed results of the next generation of neutrinoless double beta decay experiments and normal beta decay experiments.  Our principal results are summarized in Table~\\ref{tab:summary}.  \\begin{table}[t] \\caption{ Answers to some questions about the potential of neutrinoless double beta decay experiments. Answers refer to a CL of 99.73 \\%CL for the assumed probability distributions. We adopt a sensitivity $s$ equal to what is projected for the Majorana experiment \\protect\\cite{majorana} (if another reference sensitivity $s'$ is assumed, the required number of experiments should be scaled by $N'_{exp}=N_{exp} s/s'$). If the answer for an inverted neutrino mass hierarchy is different from the answer for a normal mass hierarchy (see Fig.~\\ref{fig:hierarchy}), we show in parentheses the answer for a normal mass hierarchy. \\label{tab:summary}} \\begin{tabular}{cccc} Section & Assumptions & Question  &$N_{\\rm exp}$ at 99.73 \\%CL\\\\ \\noalign{\\smallskip}\\hline \\noalign{\\medskip} II & No  detected neutrinoless double $\\beta$-decay  &  Dirac ? & 230 $(\\infty$) \\\\ \\noalign{\\smallskip}\\hline \\noalign{\\medskip} III & lightest mass scale (1 $\\pm$ 0.05 eV), ~No  neutrinoless double $\\beta$-decay &  Dirac ? &  1 \\\\ \\noalign{\\medskip} III & lightest mass scale (0.35 $\\pm$ 0.07 eV), ~No  neutrinoless double $\\beta$-decay  &  Dirac ? &  5 (6) \\\\ \\noalign{\\medskip} III & lightest mass scale (0.3 $\\pm$ 0.1 eV), ~No  neutrinoless double $\\beta$-decay &  Dirac ? & 16 ($\\infty$) \\\\ \\noalign{\\smallskip}\\hline \\noalign{\\medskip} IV &  Neutrinoless double $\\beta$-decay: ~$T_{\\frac{1}{2}} (^{76}{\\rm Ge})= (3.2 \\pm 0.2) \\times 10^{25}$ yr & Total mass ? &  [0.46,9.56] ([0.48,9.58])\\\\ \\noalign{\\medskip} IV & Neutrinoless double $\\beta$-decay: ~$T_{\\frac{1}{2}} (^{76}{\\rm Ge})= (1.\\pm 0.1) \\times 10^{26}$ yr & Total mass ? &  [0.24,8.34] ([0.28,8.40])\\\\ \\noalign{\\medskip} IV &  Neutrinoless double $\\beta$-decay: ~$T_{\\frac{1}{2}} (^{76}{\\rm Ge})= (3.2 \\pm 0.5) \\times 10^{26}$ yr & Total mass ? &  [0.08,5.68] ([0.16,6.06])\\\\ \\noalign{\\smallskip}\\hline \\noalign{\\medskip} V & Detected neutrinoless double $\\beta$-decay &  Hierarchy ? &  No \\\\ \\noalign{\\medskip} V & Detected neutrinoless double $\\beta$-decay, ~{\\rm ~ private~communication:}~ m~=~0&  Hierarchy ? &  Yes \\\\ \\end{tabular} \\end{table}  \\subsection{How can we estimate the uncertainties in calculated nuclear matrix elements?} \\label{subsec:introuncertainties}  The uncertainty in the calculated nuclear matrix elements for neutrinoless double beta decay will constitute the principal obstacle to answering some basic questions about neutrinos.  The essential problem is that the correct theory of nuclei is QCD, a notoriously difficult theory to do calculations with for nuclei with several nucleons. For neutrinoless double beta decay, the situation is even more severe since double beta candidates involve systems with $A \\sim 50$ to $A \\sim 100$ and even larger. Very attractive next generation experiments have been proposed for a number of different isotopes, including $^{48}$Ca~\\cite{KIS01}, $^{76}$Ge~\\cite{majorana,ZDE01,KLA01b}, $^{100}$Mo~\\cite{NEMO3,EJI00}, $^{116}$Cd~\\cite{BEL01}, $^{130}$Te~\\cite{ZUB01,AVI}, $^{136}$Xe\\cite{XMASS,CAC01,Gratta}, $^{150}$Nd~\\cite{ISH00}, and $^{160}$Gd~\\cite{DAN01,WANGS01}.  In the foreseeable future, it does not seem possible to derive in a  direct and controlled manner from QCD nuclear matrix elements for large $A$. Thus there is no way of quantifying with absolute confidence the range of uncertainties in nuclear matrix elements calculated with different theoretical models or approximations.  In the absence of being able to derive the errors directly from QCD, we assume that the published range of calculated matrix elements defines a plausible approximation to the uncertainty in our knowledge of the matrix elements. We do not, for example, favor a particular calculation because it happens to give better agreement with the inferred matrix element for two-neutrino double beta-decay (in the rare cases where this decay has been observed). We have no way of knowing for sure what the improved agreement for the two-neutrino case implies for the neutrinoless double-beta decay matrix element and whether, indeed, the agreement in a special case is accidental or not\\footnote{Fukugita and Yanagita~\\cite{fukugitabook} note that the nuclear levels that are important for neutrinoless double beta-decay are typically at excitation energies of order 10 MeV, while for two neutrino double beta-decay the characteristic excitation energies are lower, a few MeV. Thus even if the lower excitation states are correctly described, there is no guarantee that the higher excitation states are also correctly described.}.  We recognize that different individuals may regard the calculated range of nuclear matrix elements as either too narrow or too broad to reflect the actual uncertainty. However, we do not know of any way to settle objectively and conclusively whether our estimate of the uncertainty is pessimistic or optimistic in any particular case.  \\subsection{Some definitions} \\label{subsec:somedefinitions}  The neutrino mass matrix element that appears in neutrinoless double beta decay \\cite{Vogel,ev,rev} is given by \\begin{equation} |\\langle m^\\nu_{ee} \\rangle| ~=~ m_e\\frac{1}{\\sqrt{T_{\\frac{1}{2}}F_N}} ~=~m_e \\sqrt{{\\frac{\\lambda}{\\ln2 F_N}}} \\, , \\label{eq:matrixelementdefn} \\end{equation} where $m_e$ is the electron mass, $T_{\\frac{1}{2}} (\\lambda)$ is the half life (exponential decay constant) of the double beta decay process, and the nuclear structure parameter $F_N$ is given by  \\begin{equation} F_N =  G^{0\\nu} \\left|M_F^{0\\nu} - \\left(\\frac{g_A}{g_V}\\right)^2 M_{GT}^{0\\nu}\\right|^2 \\, . \\end{equation}    For specificity, we consider a neutrinoless double beta decay experiment with sensitivity to $T_{\\frac{1}{2}}F_N$ that is exemplified by what is expected for the Majorana experiment \\cite{majorana} (see also, compilation in Ref.~\\cite{ev}). We will consider that a number $N_{exp}$ of neutrinoless double beta decay experiments are performed with the expected Majorana sensitivity $s$.  Our results are, however, general. If the experiments that actually are or could be performed have a different sensitivity, then our results should be rescaled by  \\begin{equation} N'_{exp}~=~ N_{exp} s/s'\\, . \\label{eq:rescale} \\end{equation} If a specific neutrinoless double beta-decay experiment successfully detects a signal, then a greatly increased exposure with the same detector will not improve much the confidence with which one can answer the questions raised in this paper. For a single detector, the uncertainty will be dominated by the nuclear factor of that nucleus. Measurements with different nuclei will be required to improve the statistical significance of the answers to questions about the nature and properties of neutrinos. On the other hand, suppose the search for neutrinoless double beta decay is negative with a given detector. Then an increase in the exposure time by a factor $N_{\\rm exposure}$ is equivalent to performing $N_{\\rm exposure}$ new experiment that have the identical sensitivity.  \\begin{figure}[!t] \\centerline{\\includegraphics[width=3.75in]{fn.ps}} \\caption{Distribution of $^{76}$Ge nuclear factor results. A compilation of 20 different calculations spans in the range $2.7 \\times 10^{-15} - 2.9 \\times 10^{-13}$ yr$^{-1}$ \\protect\\cite{majorana} }. \\label{fig:fn} \\end{figure}   The neutrino mass element $|\\langle m^\\nu_{ee} \\rangle|$ is related to the fundamental neutrino parameters by the expression \\begin{equation} |\\langle m^\\nu_{ee} \\rangle| ~=~ \\left| m_1 |U^2_{e1}| e^{i\\phi_1} + m_2 |U^2_{e2}| e^{i\\phi_2} +  m_3 |U^2_{e3}| \\right| \\, ,\\label{eq:massIH} \\end{equation} where $m_i$ are the mass eigenvalues of the Majorana neutrinos, $U$ is the lepton mixing matrix, and $\\phi_i$ are relative Majorana phases. Normal (Inverted) hierarchy corresponds to the ratio between mass eigenvalues (labelled in increasing mass eigenvalue $m_1<m_2<m_3$) given by $m_3/m_2>m_2/m_1$ ($m_3/m_2<m_2/m_1$). If hierarchies are indistinguishable, what happens when $\\Delta m^2_{ij} << m^2_1$, then the mass scheme is called degenerate.   \\subsection{The dispersion in calculated nuclear matrix elements} \\label{subsec:disperioncalculated}     The dispersion of the calculated nuclear matrix elements obtained by different theoretical methods is large. For example,  a compilation of 20 different calculations \\cite{majorana,Tretyak:1995rj,Civitarese:2002tu} for $^{76}$Ge spans the range $2.7 \\times 10^{-15} - 2.9 \\times 10^{-13}$ yr$^{-1}$.  Figure~\\ref{fig:fn} shows the distribution of $^{76}$Ge nuclear factors binned in a logarithmic scale. In our analyses of how much we can learn about different fundamental neutrino questions, we will also consider $F_N$ as a random variable in linear and logarithmic scales of the constant and the Gaussian probability distributions. For the Gaussian distribution, we will adopt  the central value of the $F_N$ interval as the mean and one third of the radius of the interval covered by calculated values of $F_N$ as the standard deviation. The lowest nuclear factor $F_N$ shown in Fig.~\\ref{fig:fn} corresponds to a recent calculation~\\cite{Bobyk:2000dw} that used a self-consistent renormalized quasi-particle random phase approximation. We do not know of any rigorous argument that would exclude this recent calculation while including the other calculations shown in the figure.  For the numerical calculations given in this paper,  we used the distribution of calculated nuclear factor for  $^{76}$Ge because this nucleus is the one for which we found the largest number of published calculations of $F_N$. We performed Kolmogorov-Smirnov tests to test if the distributions of $F_N$ that were calculated for other double beta-decay candidates ( $^{82}$Se, $^{130}$Te, $^{136}$Xe) are consistent with the distribution shown in Fig.~\\ref{fig:fn} of $^{76}$Ge. Table 2 of Ref.~\\cite{ev} compiles a list of six calculations \\cite{calc} for these nuclei. The Kolmogorov-Smirnov tests show that we can not reject at 95\\%CL, for any of the nuclei $^{82}$Se, $^{130}$Te, or $^{136}$Xe, the hypothesis that the distribution of calculations of $F_N$ given in Table 2 of Ref.~\\cite{ev} is the same distribution as shown in Fig.~\\ref{fig:fn} for $^{76}$Ge. We also checked that the distribution of the six calculations listed in Table 2 of Ref.~\\cite{ev} for $^{76}$Ge is consistent with the distribution of  20 calculations of $F_N$ used  in the present work.  The fact that the uncertainty in the nuclear matrix element plays a major role in our ability to resolve fundamental questions in neutrinoless double beta decay experiments is well known (see for example the famous reviews in Ref.~\\cite{rev}). Reference~\\cite{Bilenky:2004wn} is the most recent example with which we are familiar of a systematic analysis that assumes a small uncertainty in the nuclear matrix elements for neutrinoless double beta-decay experiments (for the nuclear physics discussion see Ref.~\\cite{Rodin:2003eb}). The discussion in Ref.~\\cite{Bilenky:2004wn} assumes the correctness of the renormalized quasi-particle random phase approximation (RQRPA) that leads to the lowest nuclear factor in Fig.~\\ref{fig:fn}. Readers who are optimistic regarding the validity of current calculational methods for calculating nuclear matrix elements in neutrinoless double beta-decay may prefer the conclusions of Ref.~\\cite{Bilenky:2004wn} instead of the more conservative conclusions of the present paper.  The position adopted in this paper is that the RQRPA could be accurate, or some other calculational scheme could be more accurate, but we will not know for sure how precise any approximation is until calculations can be done in a controlled manner using QCD.  Our attitude is consistent with the point of view expressed in the recent discussion of the RQRPA and QRPA approximations in Ref.~\\cite{Rodin:2003eb}.  These authors summarized their analysis with the statement~\\cite{Rodin:2003eb}: ``Even though we cannot guarantee this basic method [RQRPA] is trustworthy, we have eliminated, or at least greatly reduced, the arbitrariness commonly present in published calculations.'' In other words, the recommended prescription results in a small dispersion in calculated nuclear matrix elements, which may or may not be close to the true value.  The reader will chose what to believe based upon the reader's convictions about the accuracy of the calculations of nuclear matrix elements. We believe that the burden of proof is upon the person drawing conclusions that depend upon the size of the nuclear matrix elements. The conclusions must be supported by a proof that the matrix elements are equal to the QCD values within the stated errors.    \\begin{figure}[!t] \\begin{center} \\includegraphics[width=3.75in]{hierarchy.ps} \\end{center} \\caption{\\label{fig:hierarchy} Normal and inverted neutrino mass hierarchy. The larger splitting is $\\Delta m^2_{\\rm atm} \\simeq 2 \\times 10^{-3}$~eV$^2$;  the smaller splitting is $\\Delta m^2_{\\rm solar} \\simeq 7 \\times 10^{-5}$~eV$^2$. The hierarchy is referred as \"Degenerate\" if the square of the smallest mass is much larger than either $\\Delta m^2_{\\rm atm}$ or $\\Delta m^2_{\\rm solar}$.} \\end{figure}   Our goal is to provide, for the reader's consideration, an alternative viewpoint to the one that is usually adopted in discussing neutrinoless double-beta decay experiments. As far as we know, there is no previous systematic, quantitative study to evaluate the impact of the uncertainty in the nuclear matrix element  for different assumed probability distributions. Recently, it has been demonstrated that it is not practical to detect in neutrinoless double beta-decay experiments neutrino CP violation arising from Majorana phases \\cite{barger,Pascoli:2002qm}.   \\subsection{How do we determine how many experiments are required?} \\label{subsec:howmany}     For each question about neutrino properties that we address, we make specific assumptions about what is, or is not, observed experimentally. Depending upon the particular question we are addressing, we will assume that the neutrino masses satisfy a normal or an inverted hierarchy, as illustrated in Fig.~\\ref{fig:hierarchy}. We will also make assumptions regarding the observation, or non-observation, of a neutrino mass in ordinary (tritium) beta-decay.  Figure~\\ref{fig:mee_m1} shows the relationship between the neutrinoless double beta decay mass element $|\\langle m^\\nu_{ee} \\rangle|$ and the smallest neutrino mass $m$ \\cite{recent2,recent3,recent4}. This figure plays a key role in our discussion; we will return to Fig.~\\ref{fig:mee_m1} in Sec.~\\ref{sec:question1}, Sec.~\\ref{sec:question3}, and Sec.~\\ref{sec:question4}.  For a given set of assumptions as described above, we compute the different probability distributions that are implied by the assumed experimental constraints. In the final step of our analysis, we combine the computed probability distributions in order to determine how many experiments are required to answer a stated question at a specific confidence level.  \\subsection{What is the bottom line?} \\label{subsec:bottomline}  Some readers will only care about the bottom line. How many neutrinoless double beta-decay experiments are required in order to determine whether neutrinos are Majorana or Dirac particles?\\footnote{Note that we are referring to the dominant neutrino masses relevant to the currently observed neutrino oscillation.  Even if the dominant masses are Dirac, there may be much smaller Majorana masses not relevant to neutrino oscillation, sometimes called pseudo-Dirac.  We do not distinguish Dirac and pseudo-Dirac neutrinos in this paper.}  What fraction of the closure mass of the universe do neutrinos constitute? Can we establish whether the neutrino masses satisfy a normal or an inverted mass hierarchy?  Table~\\ref{tab:summary} summarizes our numerical results.  We state in column 2 of Table~\\ref{tab:summary} the different assumptions that we have made about future experiments.  In column 3, we give abbreviated names to the questions that we have asked. Finally, in column 4, we present a brief summary of our answers to the different physical questions about neutrinos. The reader interested in the details of how a specific question was answered can look in the section of this paper that is listed in column 1 of Table~\\ref{tab:summary}.     \\begin{figure}[!t] \\centerline{\\includegraphics[width=3.75in]{mee_m1.ps}} \\caption{The neutrinoless double beta decay mass element $|\\langle m^\\nu_{ee} \\rangle|$ versus the lowest neutrino mass $m$. The regions allowed at 90\\% CL by existing neutrino oscillation data are shown for a normal neutrino hierarchy (NH), an inverted hierarchy (IH), and degenerate neutrinos (D) (see Fig.~\\ref{fig:hierarchy} for an explanation of the different hierarchical arrangements of neutrino masses and Sec.~\\ref{sec:question4} for a description of how the allowed regions were computed). The hatched area shows the parameter space that can be excluded by the Katrin experiment\\protect \\cite{KATRIN} if no evidence for a neutrino mass is detected in tritium beta decay. The three dashed lines labelled a), b), and c) refer to three possible positive results for a next generation neutrinoless double beta decay search and are discussed in Sec.~\\ref{sec:question3} and Sec.~\\ref{sec:question4}. The dotted horizontal line near $10^{-2}$ eV illustrates the sensitivity that is expected for the Majorana experiment \\protect \\cite{majorana}. For an original version of this figure, see Ref.~\\cite{recent2}.} \\label{fig:mee_m1} \\end{figure}   \\subsection{Outline of this paper} \\label{subsec:outline}  In Sec.~\\ref{sec:question1}, we show that an impractically large number of neutrinoless double beta decay experiments would be required to show that neutrinos are Dirac particles if next generation experiments do not reveal neutrinoless double beta decay. We show in Sec.~\\ref{sec:question2} that non-observation of neutrinoless double beta decay taken together with a measurement in ordinary beta decay of the lowest neutrino mass that is near the present upper limit (e.g., $\\sim 1$ eV) would be sufficient to show that neutrinos are Dirac particles. However, if the neutrino mass is as low as $0.3$ eV or lower, then many neutrinoless double beta decay experiments would be required to show that neutrinos are Dirac particles. We present in Sec.~\\ref{sec:question3} the allowed ranges in the total mass in neutrinos if neutrinoless double beta decay is detected at different possible half-lives. Finally, we show in Sec.~\\ref{sec:question4} that even if neutrinoless double beta decay is observed in next generation experiments we nevertheless will not be able to decide from beta decay experiments alone whether the mass hierarchy is normal or inverted.  We discuss our principal results in Sec.~\\ref{sec:summary}. ",
        "conclusions": "} \\begin{tabular}{cccc} Section & Assumptions & Question  &$N_{\\rm exp}$ at 99.73 \\%CL\\\\ \\noalign{\\smallskip}\\hline \\noalign{\\medskip} II & No  detected neutrinoless double $\\beta$-decay  &  Dirac ? & 230 $(\\infty$) \\\\ \\noalign{\\smallskip}\\hline \\noalign{\\medskip} III & lightest mass scale (1 $\\pm$ 0.05 eV), ~No  neutrinoless double $\\beta$-decay &  Dirac ? &  1 \\\\ \\noalign{\\medskip} III & lightest mass scale (0.35 $\\pm$ 0.07 eV), ~No  neutrinoless double $\\beta$-decay  &  Dirac ? &  5 (6) \\\\ \\noalign{\\medskip} III & lightest mass scale (0.3 $\\pm$ 0.1 eV), ~No  neutrinoless double $\\beta$-decay &  Dirac ? & 16 ($\\infty$) \\\\ \\noalign{\\smallskip}\\hline \\noalign{\\medskip} IV &  Neutrinoless double $\\beta$-decay: ~$T_{\\frac{1}{2}} (^{76}{\\rm Ge})= (3.2 \\pm 0.2) \\times 10^{25}$ yr & Total mass ? &  [0.46,9.56] ([0.48,9.58])\\\\ \\noalign{\\medskip} IV & Neutrinoless double $\\beta$-decay: ~$T_{\\frac{1}{2}} (^{76}{\\rm Ge})= (1.\\pm 0.1) \\times 10^{26}$ yr & Total mass ? &  [0.24,8.34] ([0.28,8.40])\\\\ \\noalign{\\medskip} IV &  Neutrinoless double $\\beta$-decay: ~$T_{\\frac{1}{2}} (^{76}{\\rm Ge})= (3.2 \\pm 0.5) \\times 10^{26}$ yr & Total mass ? &  [0.08,5.68] ([0.16,6.06])\\\\ \\noalign{\\smallskip}\\hline \\noalign{\\medskip} V & Detected neutrinoless double $\\beta$-decay &  Hierarchy ? &  No \\\\ \\noalign{\\medskip} V & Detected neutrinoless double $\\beta$-decay, ~{\\rm ~ private~communication:}~ m~=~0&  Hierarchy ? &  Yes \\\\ \\end{tabular} \\end{table}  \\subsection{How can we estimate the uncertainties in calculated nuclear matrix elements?} \\label{subsec:introuncertainties}  The uncertainty in the calculated nuclear matrix elements for neutrinoless double beta decay will constitute the principal obstacle to answering some basic questions about neutrinos.  The essential problem is that the correct theory of nuclei is QCD, a notoriously difficult theory to do calculations with for nuclei with several nucleons. For neutrinoless double beta decay, the situation is even more severe since double beta candidates involve systems with $A \\sim 50$ to $A \\sim 100$ and even larger. Very attractive next generation experiments have been proposed for a number of different isotopes, including $^{48}$Ca~\\cite{KIS01}, $^{76}$Ge~\\cite{majorana,ZDE01,KLA01b}, $^{100}$Mo~\\cite{NEMO3,EJI00}, $^{116}$Cd~\\cite{BEL01}, $^{130}$Te~\\cite{ZUB01,AVI}, $^{136}$Xe\\cite{XMASS,CAC01,Gratta}, $^{150}$Nd~\\cite{ISH00}, and $^{160}$Gd~\\cite{DAN01,WANGS01}.  In the foreseeable future, it does not seem possible to derive in a  direct and controlled manner from QCD nuclear matrix elements for large $A$. Thus there is no way of quantifying with absolute confidence the range of uncertainties in nuclear matrix elements calculated with different theoretical models or approximations.  In the absence of being able to derive the errors directly from QCD, we assume that the published range of calculated matrix elements defines a plausible approximation to the uncertainty in our knowledge of the matrix elements. We do not, for example, favor a particular calculation because it happens to give better agreement with the inferred matrix element for two-neutrino double beta-decay (in the rare cases where this decay has been observed). We have no way of knowing for sure what the improved agreement for the two-neutrino case implies for the neutrinoless double-beta decay matrix element and whether, indeed, the agreement in a special case is accidental or not\\footnote{Fukugita and Yanagita~\\cite{fukugitabook} note that the nuclear levels that are important for neutrinoless double beta-decay are typically at excitation energies of order 10 MeV, while for two neutrino double beta-decay the characteristic excitation energies are lower, a few MeV. Thus even if the lower excitation states are correctly described, there is no guarantee that the higher excitation states are also correctly described.}.  We recognize that different individuals may regard the calculated range of nuclear matrix elements as either too narrow or too broad to reflect the actual uncertainty. However, we do not know of any way to settle objectively and conclusively whether our estimate of the uncertainty is pessimistic or optimistic in any particular case.  \\subsection{Some definitions} \\label{subsec:somedefinitions}  The neutrino mass matrix element that appears in neutrinoless double beta decay \\cite{Vogel,ev,rev} is given by \\begin{equation} |\\langle m^\\nu_{ee} \\rangle| ~=~ m_e\\frac{1}{\\sqrt{T_{\\frac{1}{2}}F_N}} ~=~m_e \\sqrt{{\\frac{\\lambda}{\\ln2 F_N}}} \\, , \\label{eq:matrixelementdefn} \\end{equation} where $m_e$ is the electron mass, $T_{\\frac{1}{2}} (\\lambda)$ is the half life (exponential decay constant) of the double beta decay process, and the nuclear structure parameter $F_N$ is given by  \\begin{equation} F_N =  G^{0\\nu} \\left|M_F^{0\\nu} - \\left(\\frac{g_A}{g_V}\\right)^2 M_{GT}^{0\\nu}\\right|^2 \\, . \\end{equation}    For specificity, we consider a neutrinoless double beta decay experiment with sensitivity to $T_{\\frac{1}{2}}F_N$ that is exemplified by what is expected for the Majorana experiment \\cite{majorana} (see also, compilation in Ref.~\\cite{ev}). We will consider that a number $N_{exp}$ of neutrinoless double beta decay experiments are performed with the expected Majorana sensitivity $s$.  Our results are, however, general. If the experiments that actually are or could be performed have a different sensitivity, then our results should be rescaled by  \\begin{equation} N'_{exp}~=~ N_{exp} s/s'\\, . \\label{eq:rescale} \\end{equation} If a specific neutrinoless double beta-decay experiment successfully detects a signal, then a greatly increased exposure with the same detector will not improve much the confidence with which one can answer the questions raised in this paper. For a single detector, the uncertainty will be dominated by the nuclear factor of that nucleus. Measurements with different nuclei will be required to improve the statistical significance of the answers to questions about the nature and properties of neutrinos. On the other hand, suppose the search for neutrinoless double beta decay is negative with a given detector. Then an increase in the exposure time by a factor $N_{\\rm exposure}$ is equivalent to performing $N_{\\rm exposure}$ new experiment that have the identical sensitivity.  \\begin{figure}[!t] \\centerline{\\includegraphics[width=3.75in]{fn.ps}} \\caption{Distribution of $^{76}$Ge nuclear factor results. A compilation of 20 different calculations spans in the range $2.7 \\times 10^{-15} - 2.9 \\times 10^{-13}$ yr$^{-1}$ \\protect\\cite{majorana} }. \\label{fig:fn} \\end{figure}   The neutrino mass element $|\\langle m^\\nu_{ee} \\rangle|$ is related to the fundamental neutrino parameters by the expression \\begin{equation} |\\langle m^\\nu_{ee} \\rangle| ~=~ \\left| m_1 |U^2_{e1}| e^{i\\phi_1} + m_2 |U^2_{e2}| e^{i\\phi_2} +  m_3 |U^2_{e3}| \\right| \\, ,\\label{eq:massIH} \\end{equation} where $m_i$ are the mass eigenvalues of the Majorana neutrinos, $U$ is the lepton mixing matrix, and $\\phi_i$ are relative Majorana phases. Normal (Inverted) hierarchy corresponds to the ratio between mass eigenvalues (labelled in increasing mass eigenvalue $m_1<m_2<m_3$) given by $m_3/m_2>m_2/m_1$ ($m_3/m_2<m_2/m_1$). If hierarchies are indistinguishable, what happens when $\\Delta m^2_{ij} << m^2_1$, then the mass scheme is called degenerate.   \\subsection{The dispersion in calculated nuclear matrix elements} \\label{subsec:disperioncalculated}     The dispersion of the calculated nuclear matrix elements obtained by different theoretical methods is large. For example,  a compilation of 20 different calculations \\cite{majorana,Tretyak:1995rj,Civitarese:2002tu} for $^{76}$Ge spans the range $2.7 \\times 10^{-15} - 2.9 \\times 10^{-13}$ yr$^{-1}$.  Figure~\\ref{fig:fn} shows the distribution of $^{76}$Ge nuclear factors binned in a logarithmic scale. In our analyses of how much we can learn about different fundamental neutrino questions, we will also consider $F_N$ as a random variable in linear and logarithmic scales of the constant and the Gaussian probability distributions. For the Gaussian distribution, we will adopt  the central value of the $F_N$ interval as the mean and one third of the radius of the interval covered by calculated values of $F_N$ as the standard deviation. The lowest nuclear factor $F_N$ shown in Fig.~\\ref{fig:fn} corresponds to a recent calculation~\\cite{Bobyk:2000dw} that used a self-consistent renormalized quasi-particle random phase approximation. We do not know of any rigorous argument that would exclude this recent calculation while including the other calculations shown in the figure.  For the numerical calculations given in this paper,  we used the distribution of calculated nuclear factor for  $^{76}$Ge because this nucleus is the one for which we found the largest number of published calculations of $F_N$. We performed Kolmogorov-Smirnov tests to test if the distributions of $F_N$ that were calculated for other double beta-decay candidates ( $^{82}$Se, $^{130}$Te, $^{136}$Xe) are consistent with the distribution shown in Fig.~\\ref{fig:fn} of $^{76}$Ge. Table 2 of Ref.~\\cite{ev} compiles a list of six calculations \\cite{calc} for these nuclei. The Kolmogorov-Smirnov tests show that we can not reject at 95\\%CL, for any of the nuclei $^{82}$Se, $^{130}$Te, or $^{136}$Xe, the hypothesis that the distribution of calculations of $F_N$ given in Table 2 of Ref.~\\cite{ev} is the same distribution as shown in Fig.~\\ref{fig:fn} for $^{76}$Ge. We also checked that the distribution of the six calculations listed in Table 2 of Ref.~\\cite{ev} for $^{76}$Ge is consistent with the distribution of  20 calculations of $F_N$ used  in the present work.  The fact that the uncertainty in the nuclear matrix element plays a major role in our ability to resolve fundamental questions in neutrinoless double beta decay experiments is well known (see for example the famous reviews in Ref.~\\cite{rev}). Reference~\\cite{Bilenky:2004wn} is the most recent example with which we are familiar of a systematic analysis that assumes a small uncertainty in the nuclear matrix elements for neutrinoless double beta-decay experiments (for the nuclear physics discussion see Ref.~\\cite{Rodin:2003eb}). The discussion in Ref.~\\cite{Bilenky:2004wn} assumes the correctness of the renormalized quasi-particle random phase approximation (RQRPA) that leads to the lowest nuclear factor in Fig.~\\ref{fig:fn}. Readers who are optimistic regarding the validity of current calculational methods for calculating nuclear matrix elements in neutrinoless double beta-decay may prefer the conclusions of Ref.~\\cite{Bilenky:2004wn} instead of the more conservative conclusions of the present paper.  The position adopted in this paper is that the RQRPA could be accurate, or some other calculational scheme could be more accurate, but we will not know for sure how precise any approximation is until calculations can be done in a controlled manner using QCD.  Our attitude is consistent with the point of view expressed in the recent discussion of the RQRPA and QRPA approximations in Ref.~\\cite{Rodin:2003eb}.  These authors summarized their analysis with the statement~\\cite{Rodin:2003eb}: ``Even though we cannot guarantee this basic method [RQRPA] is trustworthy, we have eliminated, or at least greatly reduced, the arbitrariness commonly present in published calculations.'' In other words, the recommended prescription results in a small dispersion in calculated nuclear matrix elements, which may or may not be close to the true value.  The reader will chose what to believe based upon the reader's convictions about the accuracy of the calculations of nuclear matrix elements. We believe that the burden of proof is upon the person drawing conclusions that depend upon the size of the nuclear matrix elements. The conclusions must be supported by a proof that the matrix elements are equal to the QCD values within the stated errors.    \\begin{figure}[!t] \\begin{center} \\includegraphics[width=3.75in]{hierarchy.ps} \\end{center} \\caption{\\label{fig:hierarchy} Normal and inverted neutrino mass hierarchy. The larger splitting is $\\Delta m^2_{\\rm atm} \\simeq 2 \\times 10^{-3}$~eV$^2$;  the smaller splitting is $\\Delta m^2_{\\rm solar} \\simeq 7 \\times 10^{-5}$~eV$^2$. The hierarchy is referred as \"Degenerate\" if the square of the smallest mass is much larger than either $\\Delta m^2_{\\rm atm}$ or $\\Delta m^2_{\\rm solar}$.} \\end{figure}   Our goal is to provide, for the reader's consideration, an alternative viewpoint to the one that is usually adopted in discussing neutrinoless double-beta decay experiments. As far as we know, there is no previous systematic, quantitative study to evaluate the impact of the uncertainty in the nuclear matrix element  for different assumed probability distributions. Recently, it has been demonstrated that it is not practical to detect in neutrinoless double beta-decay experiments neutrino CP violation arising from Majorana phases \\cite{barger,Pascoli:2002qm}.   \\subsection{How do we determine how many experiments are required?} \\label{subsec:howmany}     For each question about neutrino properties that we address, we make specific assumptions about what is, or is not, observed experimentally. Depending upon the particular question we are addressing, we will assume that the neutrino masses satisfy a normal or an inverted hierarchy, as illustrated in Fig.~\\ref{fig:hierarchy}. We will also make assumptions regarding the observation, or non-observation, of a neutrino mass in ordinary (tritium) beta-decay.  Figure~\\ref{fig:mee_m1} shows the relationship between the neutrinoless double beta decay mass element $|\\langle m^\\nu_{ee} \\rangle|$ and the smallest neutrino mass $m$ \\cite{recent2,recent3,recent4}. This figure plays a key role in our discussion; we will return to Fig.~\\ref{fig:mee_m1} in Sec.~\\ref{sec:question1}, Sec.~\\ref{sec:question3}, and Sec.~\\ref{sec:question4}.  For a given set of assumptions as described above, we compute the different probability distributions that are implied by the assumed experimental constraints. In the final step of our analysis, we combine the computed probability distributions in order to determine how many experiments are required to answer a stated question at a specific confidence level.  \\subsection{What is the bottom line?} \\label{subsec:bottomline}  Some readers will only care about the bottom line. How many neutrinoless double beta-decay experiments are required in order to determine whether neutrinos are Majorana or Dirac particles?\\footnote{Note that we are referring to the dominant neutrino masses relevant to the currently observed neutrino oscillation.  Even if the dominant masses are Dirac, there may be much smaller Majorana masses not relevant to neutrino oscillation, sometimes called pseudo-Dirac.  We do not distinguish Dirac and pseudo-Dirac neutrinos in this paper.}  What fraction of the closure mass of the universe do neutrinos constitute? Can we establish whether the neutrino masses satisfy a normal or an inverted mass hierarchy?  Table~\\ref{tab:summary} summarizes our numerical results.  We state in column 2 of Table~\\ref{tab:summary} the different assumptions that we have made about future experiments.  In column 3, we give abbreviated names to the questions that we have asked. Finally, in column 4, we present a brief summary of our answers to the different physical questions about neutrinos. The reader interested in the details of how a specific question was answered can look in the section of this paper that is listed in column 1 of Table~\\ref{tab:summary}.     \\begin{figure}[!t] \\centerline{\\includegraphics[width=3.75in]{mee_m1.ps}} \\caption{The neutrinoless double beta decay mass element $|\\langle m^\\nu_{ee} \\rangle|$ versus the lowest neutrino mass $m$. The regions allowed at 90\\% CL by existing neutrino oscillation data are shown for a normal neutrino hierarchy (NH), an inverted hierarchy (IH), and degenerate neutrinos (D) (see Fig.~\\ref{fig:hierarchy} for an explanation of the different hierarchical arrangements of neutrino masses and Sec.~\\ref{sec:question4} for a description of how the allowed regions were computed). The hatched area shows the parameter space that can be excluded by the Katrin experiment\\protect \\cite{KATRIN} if no evidence for a neutrino mass is detected in tritium beta decay. The three dashed lines labelled a), b), and c) refer to three possible positive results for a next generation neutrinoless double beta decay search and are discussed in Sec.~\\ref{sec:question3} and Sec.~\\ref{sec:question4}. The dotted horizontal line near $10^{-2}$ eV illustrates the sensitivity that is expected for the Majorana experiment \\protect \\cite{majorana}. For an original version of this figure, see Ref.~\\cite{recent2}.} \\label{fig:mee_m1} \\end{figure}   \\subsection{Outline of this paper} \\label{subsec:outline}  In Sec.~\\ref{sec:question1}, we show that an impractically large number of neutrinoless double beta decay experiments would be required to show that neutrinos are Dirac particles if next generation experiments do not reveal neutrinoless double beta decay. We show in Sec.~\\ref{sec:question2} that non-observation of neutrinoless double beta decay taken together with a measurement in ordinary beta decay of the lowest neutrino mass that is near the present upper limit (e.g., $\\sim 1$ eV) would be sufficient to show that neutrinos are Dirac particles. However, if the neutrino mass is as low as $0.3$ eV or lower, then many neutrinoless double beta decay experiments would be required to show that neutrinos are Dirac particles. We present in Sec.~\\ref{sec:question3} the allowed ranges in the total mass in neutrinos if neutrinoless double beta decay is detected at different possible half-lives. Finally, we show in Sec.~\\ref{sec:question4} that even if neutrinoless double beta decay is observed in next generation experiments we nevertheless will not be able to decide from beta decay experiments alone whether the mass hierarchy is normal or inverted.  We discuss our principal results in Sec.~\\ref{sec:summary}.      \\label{sec:summary}  Next generation neutrinoless double beta decay experiments offer the promise of a fundamental discovery, namely, that neutrinos are their own anti-particles. No other feasible experimental technique could establish this profound result. If a single experiment conclusively detects zero neutrino double beta decay, then weak interaction theory will be both profoundly simplified and greatly clarified.  Even if neutrinoless double beta decay is not observed in next generation experiments, we may still be able to conclusively determine the particle/anti-particle nature of neutrinos. If an ordinary beta-decay experiment detects a neutrino mass near 1 eV, then we will be able to conclude in this case that neutrinos are Dirac not Majorana particles.  In all other cases, the situation will be much less favorable, as can be seen readily from the summary given in Table~\\ref{tab:summary}. If ordinary beta-decay reveals a neutrino mass scale of less than 0.3 eV, then we will not be able to conclude that neutrinos are Dirac particles from the non-observation of neutrinoless double beta decay in currently envisioned experiments. The particle/anti-particle nature of neutrinos will remain ambiguous.  The observation of neutrinoless double beta decay will determine a large allowed range of the total mass in the form of neutrinos, a range that permits an uncertainty in the total mass of between one and two orders of magnitude.  This range translates into a total cosmic neutrino mass density ~ (cf. \\cite{WMAP}) $\\Omega_{\\nu} =0.009-0.20$, $\\Omega_{\\nu}=0.005-0.17$, or $\\Omega_{\\nu} =0.0016-0.12$ at 3$\\sigma$ for the three assumed lifetimes listed in Table~\\ref{tab:summary} and discussed in Sec.~\\ref{sec:question3}.  Finally, we note that we will not be able to decide whether the neutrino mass hierarchy is normal or inverted even if neutrinoless double beta decay is detected. In order to decide this important question, information from other types of experiments like long baseline oscillation studies will be necessary."
    },
    "0403/astro-ph0403500_arXiv.txt": {
        "abstract": "We perform three-dimensional, Newtonian hydrodynamic simulations with a nuclear equation of state to investigate the accretion dynamics in neutron star black hole systems. We find as a general result that non-spinning donor stars yield larger circularization radii than corotating donors. Therefore, the matter from a neutron star without spin will more likely settle into an accretion disk outside the Schwarzschild radius. With the used stiff equation of state we find it hard to form an accretion disk that is promising to launch a gamma-ray burst. In all relevant cases the core of the neutron star survives and keeps orbiting the black hole as a mini neutron star for the rest of the simulation time (up to several hundred dynamical neutron star times scales). The existence of this mini neutron star leaves a clear imprint on the gravitational wave signal which thus can be used to probe the physics at supra-nuclear densities. ",
        "introduction": "Neutron star black hole binaries are thought to be excellent candidates for ground-based gravitational wave (GW) detection (e.g. Belczynski et al. 2002 and references therein). If their true coalescence rate is close to the upper end of the estimated values,  neutron star black hole systems may be among the first sources detected. They are further -together with double neutron star systems- the favoured model for the subclass of short Gamma-ray bursts (GRBs; Paczynski 1992, Narayan et al. 1992); a coincident detection of a gravitational wave ``chirp'' signal together with a GRB could be the ultimate proof for the compact binary nature of the short GRB progenitors. Moreover, the cold decompression of nuclear matter by the tidal disruption of a neutron star has been considered as a possible production site of rapid neutron capture elements (Lattimer and Schramm, 1974, 1976; Lattimer et al. 1977 ; Symbalisty and Schramm 1982, Eichler et al. 1989, Rosswog et al. 1999, Freiburghaus et al. 1999).\\\\ Full general relativistic simulations of this event are very challenging as apart from the complicated neutron star physics also the dynamic evolution of the space-time has to be followed since the systems of interest have comparable masses for both binary components and therefore a fixed space-time cannot be assumed. But the dynamics of the event is demanding enough even if treated in the Newtonian approximation. Simulations of neutron star black hole encounters have been described in a series of papers by Lee/Lee and Kluzniak (Lee and Kluzniak 1999a, 1999b, Lee 2000, Lee 2001). They used 3D smoothed particle hydrodynamics, Newtonian self-gravity plus backreaction forces from gravitational wave emission and polytropic equations of state (EOS) of varying stiffness. Janka et al. (1999) performed neutron star black hole merger calculations using a 3D PPM-code, again using Newtonian self-gravity  plus backreaction forces, but with the nuclear equation of state of Lattimer and Swesty (1991) and a detailed neutrino treatment (Ruffert et al. 1996).\\\\ The focus of our investigation in this paper is the complex accretion dynamics in a neutron star black hole system. Many of the (analytical) approximations that are successfully used in other accreting systems are of only limited value in this context. For example, the Roche-potential of point mass binaries is only a rough approximation to the neutron star with its rather flat density profile; neutron star black hole binaries are most likely not corotating like other close binaries (Bildsten and Cutler 1992) and, due to their tiny pressure scale heights in the neutron star surface, 'standard' estimates of the distance where mass transfer sets in can only give a very rough guideline.\\\\ We present in Sect. 2 analytic estimates that explain why irrotational donor stars do generally form accretion disks easier than corotating donors. We explain in Sect. 3 our numerical methods and initial conditions, results concerning the accretion dynamics and the corresponding gravitational wave emission are presented in Sect. 4. In Sect. 5 we summarize our results. ",
        "conclusions": "We have performed simulations of the accretion dynamics in neutron star black hole systems. Our simulations used a 3D smoothed particle hydrodynamics code together with the relativistic mean field nuclear equation of state of Shen et al. (1998a,1998b) and a detailed neutrino scheme. Gravity was approximated with Newtonian forces plus back-reaction from the emission of gravitational waves.\\\\ Generally we find a rather complex dynamics determined by the interplay between mass transfer (trying to increase the orbital separation), gravitational wave backreaction (acting to reduce the separation) and the reaction of the neutron star to mass loss. This reaction includes both the change of the neutron star radius and (in some cases) very violent fluid motions inside the star with amplitudes comparable to the stellar radius. In some of the cases the neutron star even seems to be close to fissioning into two fragments. Thus, a certain fraction of systems may result in a black hole orbited by two low-mass neutron stars. We find that many of the analytic estimates that are useful in other accreting systems are of limited use for neutron star black hole systems. The overall evolution is very sensitive to the neutron star black hole mass ratio.\\\\ We show generally that -if not superimposed by other effects such as a different tidal deformation or a different tidal lag angle- irrotational donor stars form easier accretion disks than corotating ones. This result is not restricted to neutron star black hole systems, but holds generally. The reason is that in the irrotational case mass overflowing the $L_1$ Lagrange point has more specific angular momentum and therefore a larger circularization radius. In the neutron star black hole case more material can avoid being fed directly into the Schwarzschild radius if the donor star is not spinning.\\\\ In the context of GRB progenitors neutron star black hole systems are usually regarded only as a slight modification of the binary neutron star case. It is often stated that they will yield the ``standard GRB central engine'', a black hole with a hot and massive accretion disk. Here, however, we find it very difficult to form such a disk. Although somewhat easier in the irrotational case, we find for all but extreme cases rather low temperature disks ($\\sim 3$ MeV) of moderate densities, which are not promising to either launch a GRB via neutrino annihilation or via energy extraction from the black hole rotation with suitable magnetic field configurations. Some of our investigated systems do not form any resolvable disk at all. At the end of the simulation we have in these cases a binary system of a black hole and a low mass neutron star engulfed in a common envelope of low density material that has been tidally removed from the neutron star.\\\\ We find in all but extreme cases that the core of the initial star survives all mass transfer episodes as a ``mini neutron star''. Test runs with polytropic equations of state ($\\Gamma=2$ and $\\Gamma=3$) show that in the case of a stiff equation of state a mini-neutron star forms while a soft EOS leads to a complete tidal disruption of the neutron star. This result is consistent with the work of Lee/Lee and Kluzniak. Janka et al. (1999) also find in some cases that the core of the neutron star survives the initial mass transfer episodes but is later on completely disrupted. This is consistent with their EOS being somewhat softer than the one we use (for a comparison of both EOSs see Figure 5 of Rosswog and Davies (2002)).\\\\ The survival of a mini neutron star and the difficulty to form a massive accretion disk are not unrelated; the surviving stellar core acts as some kind of storage that prevents mass from flowing towards the hole. Both effects are a result of the stiffness of the neutron star equation of state as can be seen from an inspection of the test runs with a stiff (neutron star core, low mass disk) and a soft polytropic equation of state (complete disruption, more massive disk).  Janka et al. (1999) et al. found values for the neutrino emission and annihilation that are rather optimistic for GRBs. In both sets of simulations different numerical methods are used, but we regard the differences resulting from different techniques as minor. We consider the uncertainty in the basic physics at supranuclear densities and the sensitivity of the system dynamics to this part of the EOS to be the decisive factors for the differences in the results. There are not many equations of state for the purpose of such simulations (temperature dependent, large range of densities, no $\\beta$-equilibrium assumed) available, the used examples of the Lattimer-Swesty EOS (simulations of Janka et al. 1999) and the Shen-EOS (our simulations) may therefore indicate the possible range of results as a function of the poorly known input physics.\\\\ In most cases huge tidal tails of decompressed neutron star material form. During the decompression phase the adiabatic exponent drops from values above 3 to values close to 4/3. This goes along with a recombination of nucleons into nuclei and therefore the release of several MeV per nucleon. These effects drive a rapid lateral expansion of the tidal tails that will therefore end up as large, structured disks around the hole.\\\\ The gravitational wave luminosity reaches peak values close to $10^{56}$ erg/s at the first encounter. The survival of the mini neutron star leaves a clear imprint on the gravitational wave signature. All the subsequent periastron approaches yield a luminosity peak roughly two orders of magnitude below the value of the first close approach. In our extreme case with mass ratio close to one both amplitudes and luminosity die off abruptly once the mini neutron star is finally disrupted. Thus a softening of the equation of state at supranuclear densities due to, for example, hyperon condensates would be directly observable in the gravitational wave signature of a neutron star black hole system.   \\vspace*{1cm} {\\bf Acknowledgements}\\\\ \\vspace*{0.1cm}\\\\ It is a pleasure to thank the Leicester supercomputer team Stuart Poulton, Chris Rudge and Richard West for their excellent support. Most of the computations reported here were performed using the UK Astrophysical Fluids Facility (UKAFF). Part of this work has been performed using the University of Leicester Mathematical Modelling Centre's supercomputer which was purchased through the EPSRC strategic equipment initiative.\\\\ This work was supported by a PPARC Rolling Grant for Theoretical Astrophysics. SR gratefully acknowledges the support via a PPARC Advanced Fellowship. RS whishes to thank the UK Astrophysical Fluid Facility (UKAFF) in Leicester for the kind hospitality during a visit, where parts of this work were accomplished, and he would like to acknowledge the funding of this visit by the EU FP5 programme.   \\begin{appendix}"
    },
    "0403/astro-ph0403670_arXiv.txt": {
        "abstract": "The anisotropies in the galaxy two-point correlation function measured from  redshift surveys exhibits deviations from the predictions of the linear theory of redshift space distortion on scales as large as $20\\, h^{-1} \\, {\\rm Mpc}$  where we expect linear theory to hold in real space.  Any attempt at analyzing the anisotropies in the redshift correlation function and determining the linear  distortion parameter $\\beta$ requires these deviations  to  be correctly modeled and taken into account. These deviations are usually attributed to galaxy random motions and these are incorporated in the analysis through a phenomenological model where the linear redshift correlation is convolved with the random pairwise  velocity distribution  function along the line of sight. We show that a substantial part of the deviations  arise from non-linear effects in the mapping from real to redshift space caused  by the coherent flows. Models which incorporate this effect provide an equally good  fit to N-body results as compared to the phenomenological model which has only the effect of random motions. We find that the pairwise velocity dispersion  predicted by all the models that we have considered are in excess of the values determined directly from the N-body simulations. This indicates a shortcoming in our understanding of the statistical properties of peculiar velocities and their relation to redshift distortion. ",
        "introduction": "Galaxy redshifts are not perfectly described by pure Hubble's law. Density fluctuations  induce peculiar velocities relative to the general Hubble expansion. The peculiar velocities perturb galaxy redshifts which in turn affects  their inferred distances and this leads to a systematic distortion in the clustering pattern of galaxies in redshift space. The peculiar velocities cause the two-point correlation function in redshift space $ \\xi^s( \\s)$   to be anisotropic {\\it i.e.} it depends separately on  the component of the pair separation  $\\s$  parallel $ (s_{\\|}) $ and perpendicular $(s_\\bot )$ to the observer's line of sight $\\n$. There are two characteristic effects of peculiar velocities . On large scales structures are compressed along the line of sight due to coherent flows into over dense regions and out of under dense regions, thereby amplifying $ \\xi^s ( s_{\\|}, s_\\bot )$. On small scales $  \\xi^s ( s_{\\|}, s_\\bot)$ is suppressed due to the structures being elongated along the line of sight by random motions in virialized clusters.  \\citet{kais} first quantified the correlation anisotropy that results from large-scale peculiar flows in terms of the power spectrum of galaxy clustering.  Using linear theory and the plane parallel approximation he showed that the power spectrum in redshift space $P_s(\\k)$ and it's real space counterpart $P_r(k) $ are related as \\begin{equation} P_s(\\k) = ( 1+ \\beta \\mu ^2_k )^2 \\ P_r (k) \\end{equation} where $ \\mu_k $ is the cosine of the angle between $\\k$ and the line of sight $\\n$,  and $ \\beta \\simeq \\Omega_m^{0.6}/b$ is the linear distortion parameter. Here $\\Omega_m$ is the cosmic mass density parameter and b is the linear bias parameter which differs from  unity if the galaxies represent a biased sample of the underlying dark matter distribution. It may be noted that the factor $\\Omega_m^{0.6}$ relates peculiar velocities to density  density fluctuations \\citep{peeb}. This is slightly modified in the presence of a cosmological constant \\citep{lahav1} and it is more accurate to use $\\beta=f(\\Omega_m)/b$ where $ f(\\Omega_m)=\\Omega_m^{0.6}+\\frac{1}{70}[1-\\frac{1}{2}\\Omega_m(1+\\Omega_m)] \\,. $ The important point is that the anisotropies observed in $P_s(\\k)$ can be used to determine the value of $\\beta$, and thereby place interesting constraints on the density parameter $\\Omega_m$  and the bias $b$. This has been the single most important motivation for a substantial amount of the research which has been carried out in trying to understand and quantify the nature of redshift space distortions.  \\citet{ham} translated Kaiser's linear formula from Fourier to real space. He showed that it is most convenient to parameterize the anisotropy of $ \\xi ^s ( s_\\parallel , s_\\perp)$ in terms of spherical harmonics as \\begin{equation} \\xi^s (s_\\parallel , s_\\perp) = \\sum_{l=0}^{\\infty} \\xi_l (s) P_l (\\mu ) \\end{equation} where $s=\\sqrt{s^2_\\parallel+s^2_\\perp}$, $\\mu=s_\\parallel/s$, $P_l(\\mu)$ are the Legendre polynomials and $\\xi_l(s)$ are the different angular moments of the redshift space two-point correlation function. Only the first three even angular moments, namely the monopole $\\xi_0(s)$, the quadrupole $\\xi_2(s)$ and the hexadecapole $\\xi_4(s)$ are non-zero and these can be expressed in terms of the real space galaxy two-point correlation $\\xi(r)$ and its moments which are defined as  \\begin{equation} \\xib_n(s)=\\frac{n+1}{s^{n+1}} \\int_0^s \\xi(y) y^n dy\\,. \\label{eq:a2} \\end{equation} through \\begin{equation} \\xi_0(s)=(1+\\frac{2}{3} \\beta  + \\frac{1}{5} \\beta^2) \\,  \\xi(s) \\label{eq:a3} \\end{equation} \\begin{equation} \\xi_2(s)= (\\frac{4}{3} \\beta + \\frac{4}{7} \\beta^2) \\, \\left[\\xi(s)-\\xib_2(s) \\right] \\label{eq:a4} \\end{equation} \\begin{equation} \\xi_4(s)=\\frac{8}{35} \\beta^2 \\, [\\xi(s) + \\frac{5}{2} \\xib_2(s) -\\frac{7}{2} \\xib_4(s) ] \\label{eq:a5} \\end{equation} The linear analysis predicts a negative quadrupole ($i.e. \\, \\xi_2(s) <0$) arising from the squashing of large scale structures along the line of sight.  Hamilton proposed that the observed redshift space correlation function be decomposed into spherical harmonics, and the  ratio  \\begin{equation} Q(s)=\\frac{\\xi_2(s)}{\\frac{3}{s^3}\\int_0^s \\xi_0(s^{'}) d s^{'}- \\xi_0(s)} = \\left[ \\frac{\\frac{4}{3} \\beta + \\frac{4}{7} \\beta^2} {1+\\frac{2}{3} \\beta  + \\frac{1}{5} \\beta^2}\\right] \\label{eq:a6} \\end{equation} which is expected to have  a constant value (shown in $[...]$ in eq. \\ref{eq:a6}) be used to determine the value of $\\beta$. Alternatively, if the real space correlation function has a power law behaviour $\\xi(r)\\propto r^{-\\gamma}$, the ratio $\\xi_2(s)/\\xi_0(s)$ is also expected to be a constant, and this can be used to determine the value of $\\beta$.   Investigations using N-body simulations to study the redshift space two-point correlation (eg. \\citealt{suto},  \\citealt{fish94}, \\citealt{brain},  \\citealt{brom}) find deviations from the linear predictions out to scales as large as $20 \\, h^{-1} \\, {\\rm Mpc}$ and even larger where linear theory is known to be valid in real space. Such deviations are also seen in the redshift space two-point correlations determined from different redshift surveys (eg. \\citealt{landy} (LCRS), \\citealt{pcok} (2dFGRS), \\citealt{haw} (2dFGRS)).  In addition to the squashing predicted by the linear analysis,  the two-point correlation  function exhibits an elongation along the line of sight at scales as large as  $ 20 h^{-1} $ Mpc. This causes the quadrupole moment to remain positive even at scales where one would expect linear theory to hold in real space. The values of $Q(s)$ which are expected to be constant (eq. \\ref{eq:a6}) do not flatten out to scales as large as $20 \\, h^{-1} \\, {\\rm Mpc}$ in N-body simulations, nor is the flattening observed at these scales in the redshift surveys. All this indicates that there are non-linear effects which are important in the mapping from real space to redshift space at length-scales where linear theory is known to be valid in real space.  The most popular approach is to attribute the deviations from the linear predictions to the effects of the random peculiar velocities of galaxies located in virialized clusters and other highly non-linear regions. This effect is incorporated through a phenomenological model (eg. \\citealt{dav83}, \\citealt{fish94}, \\citealt{pcokanddodds},  \\citealt{hevtayl}, \\citealt{marj}, \\citealt{bali}, \\citealt{tadefsta}, \\citealt{brom}, \\citealt{rat}) which assumes that at large scales the deviations from linear theory can be incorporated by convolving the linear redshift space correlation function $ \\xi ^s _L $  with the line of sight component of the random, isotropic pairwise velocity distribution function  ${\\it f}(\\upsilon)$. The resulting non-linear redshift space two-point correlation function is given by \\begin{equation} \\xi^s ( s_\\parallel ,s_\\perp) = \\int \\xi^s _L ( s_\\parallel +\\upsilon , s_\\perp ) {\\it f}(\\upsilon) d \\upsilon \\end{equation} where the distribution function ${\\it f }(\\upsilon)$  is normalized to $\\int^{\\infty}_{-\\infty} { \\it f}(\\upsilon) d \\upsilon = 1 $. The authors who have invoked this model have generally adopted either a Gaussian or else an exponential pairwise velocity distribution function. In both cases, the distribution function has only one unknown quantity $\\sigma_R^2$ which is the velocity dispersion of the random component of the pairwise peculiar velocity of the galaxies. In this model, the observations of the anisotropies in $\\xi^s$ can be used to jointly determine the value of $\\beta$ and $\\sigma_R$. This has recently been accomplished for the 2dFGRS where they find $\\beta=0.49 \\pm 0.09$ and $\\sigma_R=506 \\pm 52 \\, {\\rm km s^{-1}}$ \\citep{haw}.  An alternative approach is to attribute the deviations from the linear predictions in $\\xi^s$ to non-linear effects arising from the coherent flows.  \\citet{tayl}, \\citet{fish96} and \\citet{hatton} have used the Zel'dovich approximation to analytically study  the behaviour of the redshift-space power spectrum in the translinear regime. They find that the results from the Zel'dovich approximation are in reasonable agreement with the predictions of N-body simulations,  indicating that the coherent flows may be making a significant contributions to the non-linear effects observed in the redshift space two-point correlation function.  In a different approach to studying the deviations in $\\xi^s$ from the linear predictions at scales where linear theory is known to be valid in real space, \\citet{bharad}  has considered the non-linear effects introduced by the mapping from real space to redshift space. Under the assumption that linear theory is valid in real space and that the density fluctuations are a Gaussian random filed,   $\\xi^s$ has been calculated taking into account all the non-linear effects that arise due to the mapping from real to redshift space. It may be noted that the original calculation of \\citet{kais} and \\citet{ham} treats the mapping from real to redshift space  to linear order only.  In summary, at large scales where linear theory is known to be valid in real space, the commonly used phenomenological model for $\\xi^s$ attributes all the deviations from the linear predictions to the effects of random motions on the mapping from real to redshift space. On the other hand, \\citet{bharad} calculated $\\xi^s$ incorporating all the non-linear effects which arise in the mapping from real to redshift space assuming that they are caused only by the coherent flows. In all probability, the deviations from the linear predictions found in $\\xi^s$  in the N-body simulations and in actual redshift surveys is a consequence of non-linear effects in the mapping from real space to redshift space arising from both these effects namely, random motions and  coherent flows. In this paper we consider models for the redshift space distortions which combine both these effects, We compare the predictions of these models with the commonly used phenomenological model which has only the non-linear effects from random motions. We also compare all  these models with N-body simulations and investigate which model best  fits the N-body results. The different models are presented in Section 2 and the results of the comparison with N-body simulations are presented in Section 3. The galaxy pairwise  velocity dispersion is a quantity which crops up in any discussion of the effects of redshift space distortions on the two-point correlation function. This quantity is very interesting in its own right and it has received a considerable amount of attention (\\citealt{dav83}, \\citealt{bin},\\citealt{mojingbo}, \\citealt{brain}, \\citealt{somr},\\citealt{bharad97},\\citealt{mojing}, \\citealt{landy}, \\citealt{rat}, \\citealt{striker}, \\citealt{jingbo},\\citealt{jingbor}, \\citealt{bharad}, \\citealt{sheth2}, \\citealt{pop} ). This quantity has been observationally determined for different redshift surveys(eg. \\citealt{jingmobo}, LCRS ; \\citealt{jevi}, SDSS ; \\citealt{haw} ,2dFGRS ).  The models we use for the redshift space distortion also make definite  predictions for the pair-wise velocity dispersion.   In Section 4 we calculate the pair-wise velocity dispersion predicted by the different models and compare these with the pair-wise velocity dispersion determined directly from the N-body simulations.  In Section 5 we discuss our results and present conclusions.  We would also like to point out that the models which we have considered for $\\xi^s$  are very similar in spirit to those considered by  \\citet{mats}, \\citet{regos} and \\citet{fish95}. ",
        "conclusions": "The galaxy two-point correlation function determined from redshift surveys shows significant deviations from the predictions of the linear theory of redshift space distortion even on scales as large as $20 - 30 \\, h^{-1} {\\rm Mpc}$ where linear theory is expected to be valid on real  space. Any attempt to determine $\\beta$ from redshift surveys requires that these deviations  be properly modeled and taken into account. Modeling redshift space distortions basically requires a joint model for galaxy peculiar velocities and their correlations with the galaxy clustering pattern.  Such models test our understanding of the gravitational instability process by which the large scale structures are believed to have formed.   We have considered four different models (details in Section 2) for the  redshift space two-point correlation function $\\xi^s$. All the models are based on the assumption that galaxy peculiar velocities may be decomposed into two parts, one arising from coherent flows and another from random motions. It is also assumed that in real space the coherent flows are well described by the linear theory of density perturbation.  Deviations from the predictions of the linear theory of redshift space distortion arise from  two distinct causes  which affect the mapping from real to redshift space (a.) non-linear effects  due to the coherent flows  (b.) the random motions. Among the four models, Model A does not incorporate the non-linear effects due to the coherent flows. It combines the predictions of the linear theory of redshift space distortion (\\citealt{kais}, \\citealt{ham}) with the effect of the random motions which is modeled through an exponential distribution function for the pairwise velocity.  This is the popular phenomenological model which has been widely applied to the analysis of galaxy redshift surveys (eg. \\citealt{haw}). Models B, C and D all take into account non-linear effects arising from the coherent flows \\citep{bharad}, and they differ from one another in the choice of the distribution function for the random part of the pairwise velocity.   All the models have only one free parameter, $\\sigma_R$ which is the one dimensional random pairwise velocity dispersion. For each  model we have    determined the value of $\\sigma_R$  for which the  model predictions best fit the quadrupole anisotropy of $\\xi^s$ determined from N-body simulations. We find that Model C gives the lowest value of the best fit $\\chi^2/\\nu$   over the range of length-scales $8\\le s \\le 40 \\, h^{-1} {\\rm Mpc}$ where we expect linear theory to be valid in real space. In this  model the distribution function for the random part of a galaxy's peculiar velocity is modeled as an exponential function.  It may be noted that the other three models also give acceptable fits to the N-body results.   We find that three of the models (A, B and C) also give acceptable fits over length-scales $5 \\le s \\le 40 \\, h^{-1} {\\rm Mpc}$ which includes a small region where perturbations are expected to be mildly non-linear. Model D where the distribution function $f(v)$ for the random part of  pairwise velocity  is a Gaussian fails to give an acceptable fit. Model B where  $f(v)$ is an exponential gives the lowest value of best fit $\\chi^2/\\nu$. It may be noted that though  the best fit value of $\\chi^2/\\nu$  for model  A, the  commonly used phenomenological model, is around three times larger than that for Model B,  it is not possible to draw a strong statistical conclusion as to which model is superior. This is because $\\chi^2/\\nu<1$  for Models A,B and C  and they all provide acceptable fits. The present work is limited by the large statistical error-bars in the quantities determined at large scales from N-body simulations. These errors arise from the limited volume of the simulations (cosmic variance). It should be possible to achieve lower $1-\\sigma$ error-bars using larger simulations whereby we could  distinguish between these models at a higher level of statistical significance. We propose to carry this out in the future.    Interestingly, the best fit value of $\\sigma_R$ shows substantial variations across the models. The best fit value of $\\sigma_R$ is substantially smaller in the  models which incorporate the non-linear effects of the coherent flows (B, C and D) as compared to Model A which does not include these effects. This indicates that there are significant nonlinear effects in the mapping from real to redshift space arising from the  coherent flows.  The commonly used phenomenological model does not incorporate these effects and in this model all deviations from the linear predictions are attributed to random motions. This leads to the pairwise velocity dispersion of the random motion ($\\sigma_R$) being  overestimated. For eaxample, \\citet{haw} have used Model A to fit the redshift space two-point correlation function  of the 2dFGRS to obtain the best fit value $\\sigma_R=506 \\pm 52 \\, {\\rm km/s}$.  The findings of this paper show that Models B and C would be equally successful in fitting the same observation,  and the best fit value of $\\sigma_R$ would be different for each of these models.  This raises questions as to the interpretation of $\\sigma_R$ determined by this method.   Although the models are all reasonably successful in fitting the quadrupole anisotropies of $\\xi^s$, the model predictions for the pairwise velocity dispersion  are much  larger than the values determined directly from N-body simulations. Surprisingly, the predictions of linear theory which has a contribution from only the coherent flows and not the random motions are much closer to the N-body results as compared to the model predictions. At large scales the predictions of linear theory, all the models and the N-body results are all very similar. The linear theory predictions are slightly below the N-body results, and one would expect that it would be possible to recover the N-body results by combining the linear theory predictions  with the contribution from random motions. Unfortunately, all the models appear to be overestimating the contribution from random motions and the model predictions are significantly in excess of the N-body results. Also, the predictions of Model A fare the worst in comparison to the other models. A possible explanation why equations (\\ref{eq:c3}) and (\\ref{eq:c4}) overpredicts  the pairwise velocity dispersion is that the linear component of the peculiar velocity also makes a contribution to the random motions. It is possible that  this is already present in $\\sigma_R$, and it  contributes more than its due share to the pairwise velocity dispersion.  In the linear theory of  redshift space distortions  the hexadecapole anisotropy arises from the line of sight component of the pairwise velocity dispersion. The fact that none of the four models considered here  give a very good fit to the hexadecapole is probably related to the fact that the models also do not predict the correct   pairwise velocity dispersion.  We note that the assumption that galaxy peculiar velocities can be decomposed into two parts,  one coherent and another random is consistent with the halo model. The random part may be attributed to motions inside the halo and the coherent part to the overall motion of the halo. \\citet{seljak}, \\citet{white} and \\citet{kang} have calculated the galaxy power spectrum in redshift space using the halo model. \\citet{sheth1} have calculated the pairwise velocity dispersion using the halo model.  In conclusion we note that the nonlinear effects in the mapping from real to redshift space introduced by the coherent flows are important. Models which incorporate these effects provide an equally good   to the quadrupole anisotropies  of $\\xi^s$ as compared to models which are based on the linear theory of redshift distortion. Unfortunately, none of these models make correct predictions for the pairwise velocity dispersion. This indicates that there is a gap in our understanding of the statistical properties of the peculiar velocities and their effect on the redshift space two-point correlation function."
    },
    "0403/astro-ph0403393_arXiv.txt": {
        "abstract": "Recently, deuterium has been the focus of a high level of experimental and theoretical activity that was sparked by a disagreement on the experimental value of the maximum compression along the Hugoniot.  The behavior of deuterium at Mbar pressures is not well understood.  It is of great interest to understand how the current uncertainty on the hydrogen/deuterium EOS affects the inferred structures of Jupiter and Saturn.  In particular, the mass of a core of heavy elements (other than H and He) and the total mass of those heavy elements in these two planets are quite sensitive to the EOS of hydrogen and constitute important clues to their formation process.  We present a study of the range of structures allowed for Jupiter and Saturn by the current uncertainty in the hydrogen EOS and astrophysical observations of the two planets.  An improved experimental understanding of hydrogen at Mbar pressures and better determinations of the gravitational moments of both planets are necessary to put tight bounds on their internal structure. ",
        "introduction": "Dense hydrogen has been the subject of numerous theoretical and experimental studies over the last half century, in part because it has the simplest electronic structure of all the elements and because of its prevalence in the universe.  This simplicity is only superficial however, as multiple experiments using static and dynamic compression to Mbar pressures have revealed increasingly complex properties and a rich phase diagram.  For example, there are three known phases for the molecular solid \\citep{mao94}, the molecular fluid becomes conducting at $P\\wig> 1.4\\,$Mbar and $T \\sim 4000\\,$K (Weir, Mitchell \\& Nellis 1996), the molecular nature of hydrogen persists to pressures above 3$\\,$Mbar (Loubeyre, Occelli \\& LeToullec 2002), and superfluid and superconducting phases have been suggested \\citep{ashcroft03}.  Theoretical and experimental research on hydrogen is usually motivated by interest in warm dense matter problems in condensed matter physics, and applications to inertial confinement fusion research and to the interiors of giant planets.  Indeed, Jupiter and Saturn are the largest reservoirs of fluid metallic hydrogen in the solar system and our understanding of these two planets depends rather critically on the properties of hydrogen, especially its equation of state (EOS).  Jupiter and Saturn are composed of about 50\\% to 70\\% hydrogen by mass.  At pressures of $P \\sim 1 - 3\\,$Mbar (corresponding to $\\sim 80$\\% and $\\sim 60$\\% of the planet's radius for Jupiter and Saturn, respectively), hydrogen goes from an insulating molecular fluid to an atomic metallic fluid.  This transition has been the subject of much theoretical and experimental work but remains poorly understood.  Recent shock compression experiments of deuterium to Mbar pressures disagree significantly \\citep{collins98,knudson04,belov02}.  The internal structure of jovian planets is inferred indirectly from their global properties.  Their rapid rotation results in a noticeable deformation and a non-spherical gravitational field that can be expressed as an expansion in Legendre polynomials: \\begin{equation} V(r,\\theta)=-{GM \\over r}\\Bigg[1 - \\sum^\\infty_{n=1} \\bigg({R_{\\rm eq} \\over r}\\bigg)^n J_n P_n(\\cos \\theta)\\Bigg], \\end{equation} where $G$ is the gravitational constant, $M$ the mass of the planet, $R_{\\rm eq}$ its equatorial radius, and the coefficients $J_n$ are given by \\begin{equation} J_n=-{1 \\over MR_{\\rm eq}^n}\\int_V {r^\\prime}^n P_n(\\cos \\theta) \\rho(r^\\prime,\\theta)\\,d^3r^\\prime \\end{equation} and are known as the gravitational moments.  The integral is performed over the volume of the planet.  To a very high degree of accuracy, the planet is in hydrostatic equilibrium and the symmetry between the northern and southern hemispheres implies that $J_n=0$ for odd $n$. The first three non-vanishing moments, $J_2$, $J_4$, and $J_6$ have been measured during spacecraft flybys of both planets.  Combined with the total mass, the radius, and the rotation period of the planet, these provide integral constraints on the density profile of the planet $\\rho(r,\\theta)$.  These constraints cannot be inverted to obtain the density profile, however.  Instead, a simple model is assumed for the structure and composition of the rotating planet which is then subjected to the condition of hydrostatic equilibrium, including the rotational potential perturbation.  Model parameters are adjusted to fit the observed constraints.  The EOS provides the $P(\\rho)$ relation needed to close the system of equations.  The structure thus inferred is sensitive to the choice of hydrogen EOS used in the models.  The total amount of heavy elements and their distribution inside these two planets bear directly on their formation process by accretion of both gaseous and solid material from the protoplanetary nebula.  It is generally thought that Jupiter and Saturn formed according to the core accretion model.  In the regions where they formed, refractory compounds (most significantly, water) had condensed and formed small solid bodies.  Collisional processes between these planetesimals led to the build up of cores of 10 -- 20$\\,M_\\oplus\\,$ \\citep{pollack96} of heavy elements that subsequently accreted the surrounding gas on a dynamical time scale to build the large planets that we see today. \\footnote{The total masses are 317.83$\\,M_\\oplus$ and 95.147$\\,M_\\oplus$ for Jupiter and Saturn, respectively.}  Formation by a disk instability in the gaseous protoplanetary disk, which does not require a seed core of heavy elements, has also been suggested as a formation mechanism \\citep{boss00}. Whether formation occurred through core accretion or disk instability (or perhaps another mechanism) depends very much on the internal structure and distribution of heavy elements that we see today in these planets.  Since the last determination of the gravitational moments of Jupiter and Saturn two decades ago \\citep{campbell85, campbell89}, modeling of the interiors of Jupiter and Saturn has been driven primarily by the availability of improved equations of state for hydrogen.  Recently, measurements of the reshock temperature (Holmes, Ross \\& Nellis 1995), the development of laser-induced shock compression \\citep{dasilva97} and pulsed-power shock compression \\citep{knudson01} have provided a wealth of new data and fueled intense theoretical debate on the physics of hydrogen at Mbar pressures and temperatures of $\\sim 10^4\\,$K.  For the first time, there are experimental data that allow us to constrain the EOS of hydrogen inside jovian planets at pressures of $\\wig> 1\\,$Mbar, where the EOS remains most uncertain.  Experimental results along the principal Hugoniot\\footnote{A Hugoniot is a curve in $(P,V,T)$ space that represents all the possible shocked states that can be achieved from a given initial state $(P_0,V_0,T_0)$.} of deuterium do not agree \\citep{collins98, knudson01} in this pressure range.  Laser compression data give a maximum compression of $\\sim 6\\,$ \\citep{dasilva97, collins98} while the pulsed-power compression experiments (\\cite{knudson01}) and convergent spherical shock wave experiments \\citep{belov02, boriskov03} find a value of $\\sim 4\\,$.  Further study will reveal which experimental result is correct.  However, it is generally agreed that the actual Hugoniot does not lie outside the range allowed by the current experiments.  It is now possible to assign realistic error bars on the principal Hugoniot of deuterium, and hence on the EOS at pressures of a few Mbar.  Until now, estimating the uncertainties on the EOS of hydrogen at these pressures had been little more than guess work.  In view of these new developments, a new study of the interiors of Jupiter and Saturn is warranted.  In this paper, we examine how the uncertainty on the EOS of hydrogen maps unto uncertainties on the interior structure of Jupiter and Saturn.  The resulting range of models provide firm bounds on the amount and distribution of heavy elements in these two planets.  We first generated a number of EOS, supplemented with existing tabular EOS to represent subsets of $(P,V,T)$ data along the principal and reshock Hugoniots of deuterium. Together, these EOS give Hugoniots that reproduce the extremes found in the data as well as intermediate behavior.  The EOS are described in section 2.  The computation of interior models of Jupiter and Saturn, described in section 3, includes many improvements over previous work.  We present the new models and the effect of the hydrogen EOS on the cooling time scale of Jupiter in sections 4 and 5, respectively. The astrophysical implications are discussed in section 6. ",
        "conclusions": "We have computed a new generation of interior models for Jupiter and Saturn, with an emphasis on improving the treatment of the EOS of hydrogen and of elements heavier than H and He. In particular, we have used 7 different EOS of hydrogen that were chosen to reproduce the range of possibilities indicated by first and second shock Hugoniot data.  This allows, for the first time, a determination of the effects of the present uncertainty on the EOS of hydrogen on the interior structure of jovian planets. The parameters of interest, which characterize the total mass of heavy elements and their radial distribution in both Jupiter and Saturn are now constrained reliably by this study.  It is interesting to compare the results to previous models \\citep{ggh97, guillot99} for which a third parameter was added to the optimization: a discontinuity in the abundance of heavy elements accross the helium-poor/helium region (located at the transition between molecular and metallic hydrogen in the 3-layer model). Our results for Jupiter are very similar, in particular concerning the low central core mass and highly uncertain total mass of heavy elements. Differences arise in the case of Saturn, for which some models \\citep{guillot99} were found to fit the planet's gravitational field with no central dense core. However, these previous models consisted in solutions with a high abundance discontinuity at the molecular/metallic transition (a large abundance of heavy elements in the helium-rich region mimicking the effect of a central dense core). We can thus be relatively confident that even with a simple three-layer model with two free parameters ($M_{\\rm core}$ and $M_{\\sss Z}$) one can constrain the global structures of Jupiter and Saturn. On the other hand, we should keep in mind that the real structures of these two planets maybe more complex so that the actual interpretation of these parameters is not precisely defined (e.g. the central core could be either diluted or present as a well-defined structure).  Quantitatively, we confirm that both Jupiter and Saturn are enriched in heavy elements compared to the Sun, by factors of 1.5 -- 6 and 6 -- 14, respectively.  Maximum compression ratios of $\\sim 4$ along the principal Hugoniot are supported by three independent experiments \\citep{knudson01, belov02, boriskov03, knudson03, knudson04}, and by most ab initio EOS simulations (Lenosky, Kress \\& Collins 1997; Militzer \\& Ceperley 2000; Galli et al. 2000; Desjarlais 2003). This type of Hugoniot response is represented by the LM-SOCP and the SESAME-p EOS. Interestingly, both EOS lead to very small core masses ($M_{\\rm core} \\wig< 3\\,M_\\oplus$) in Jupiter, and a substantial amount of heavy elements in the envelope.  The same EOS predict a more massive core in Saturn (10 -- 20$\\,M_\\oplus$).  Changing the He-SCVH helium EOS for the He-SESAME-p EOS has a modest effect on the structure of Jupiter and increases the total amount of heavy elements by a few Earth masses, split more or less evenly between the core and the envelope.  This does not change the qualitative picture that emerges, however.  Taken at face value, these results imply that Jupiter formed by disk instability in the protoplanetary disk while Saturn formed by core accretion.  It is rather unlikely that the two planets formed by such different mechanisms, however.  We speculate that the only way to reconcile the formation processes of Jupiter and Saturn is that both of them formed by core accretion and that for Jupiter, the subsequent accretion of the gaseous H/He envelope resulted in partial or complete mixing of the core with the gas, increasing $M_{\\sss Z}$ at the expense of $M_{\\rm core}$.  The larger accretion rate in the proto-Jupiter may have caused larger mixing than in the proto-Saturn \\citep{gshs03}.  On the other hand, Jupiter models computed with the LM-A EOS have core masses that are (barely) consistent with the mass required for formation with core accretion formation ($\\sim 10\\,M_\\oplus$).  In this case, mixing of the core of proto-Jupiter with the accreting envelope would not be required and both planets would form by the same process.  The LM-A EOS was constructed to fit the NOVA $(P,V,T)$ Hugoniot as well as the gas gun reshock temperatures.  It represents the softest and coolest EOS allowed by the experiments.  The reshock temperatures measurements may be too low, however, and a reevaluation of these measurements may well rule out the LM-A EOS.  This study has been conducted with the point of view of learning about the interiors of Jupiter and Saturn from our current knowledge of the EOS of hydrogen and the associated uncertainties.  The opposite approach can also be considered: Can astrophysical knowledge contribute to the debate on the high-pressure EOS of hydrogen? Because much astrophysical knowledge is not amenable to direct observation or experimentation, our knowledge of processes that are hidden from view or that are no longer taking place is very sketchy. It is usually difficult to draw strong conclusions about the underlying microphysics when only the global properties of a complex, natural object are known.  The ability of a given EOS to give acceptable models depends somewhat on the assumptions for the model structure.  In general, more elaborate models (with more parameters) can accommodate a wider range of EOS. Nevertheless, we venture to comment on interesting patterns that emerge in the more extreme cases that we have encountered in this study.  We could not obtain satisfactory models of Jupiter with the original SESAME EOS because it is relatively stiff between 0.1 and 3$\\,$Mbar, and relatively soft at higher pressures along Jupiter's adiabat.  Both of these effects combine to decrease the core mass \\citep{guillot99} but even with $M_{\\rm core}=0$ models that fit the gravitational moments could not be obtained with this EOS.  It is only after it was patched at low pressures and that a 2\\% uncertainty in the EOS was introduced that acceptable models could be found (SESAME-p EOS). Similarly, we could not obtain satisfactory models of Jupiter with the LM-B EOS because it is too stiff along the adiabat above 5\\,Mbar. Taken together, these constrain the $(P,\\rho)$ relation along the adiabat a little more than the experimental results alone (see Fig.~\\ref{fig:adiab_Prho}).  We find that most EOS predict relatively small core masses for Jupiter. While it may be possible to accommodate such a situation astrophysically, it would require a substantial revision of the generally accepted formation process of the planet.  Cooling calculations provide additional information.  We find that EOS that fit the gas gun reshock temperatures (e.g. LM-A) give cooling times that are uncomfortably short for Jupiter, suggesting that those temperatures are systematically too low.  Models of Saturn do not provide any useful information on the EOS of hydrogen.  Further progress in determining the interior structures of Jupiter and Saturn can be accomplished with a two-pronged approach.  Proposed space missions to Jupiter would lead to a tenfold reduction of the uncertainty on its gravitational moments, greatly reducing the range of acceptable models ({\\it i.e.} smaller boxes in Fig.~\\ref{fig:boites-jup}).  Such improved measurements would most likely lead to rather distinct solutions for different choices of hydrogen EOS.  The Cassini spacecraft will enter Saturn orbit in July 2004 for a 4-year mission and could reduce the error bars on the gravitational moments significantly (in the context of an extended mission with accurate Saturn gravity measurements).  This would greatly reduce the range of acceptable models (Fig.~\\ref{fig:boites-sat}).  Improved astrophysical data will not be sufficient, however.  It is as important to reduce the uncertainties surrounding the EOS of hydrogen in the 1 to 30\\,Mbar range.  This can be achieved indirectly by reducing the experimental uncertainty along the principal Hugoniot, or by directly reproducing the conditions found inside giant planets with reshock and reverberation experiments \\citep{knudson03, boehly04}. From the perspective of planetary science, two additional EOS problems deserve much attention. The EOS of helium and especially that of hydrogen/helium mixtures, with the possiblity of a demixing phase transition, need to be investigated both experimentally and theoretically. The phase diagram of H/He mixtures is well-constrained by the astrophysics of Jupiter and Saturn, making this an especially interesting problem to study with the most advanced ab initio methods available.  This work was supported in part by NASA Planetary Geology \\& Geophysics grant NAG5-8906, by the United States Department of Energy under contract W-7405-ENG-36, and by the {\\it Programme National de Plan\\'etologie} (France). We thank Bill Hubbard for sharing results and codes on rotating polytropic models, and Julie Castillo for thourough discussions on expected gravity measurements with Cassini.  \\clearpage"
    },
    "0403/astro-ph0403320_arXiv.txt": {
        "abstract": "We report the results of a 15~GHz (2~cm) multi-epoch VLBA program, begun in 1994 to study the outflow in radio jets ejected from quasars and active galaxies. The observed flow of 208 distinct features measured in 110 quasars, active galaxies, and BL~Lac objects shows highly collimated relativistic motion with apparent transverse velocities typically between zero and about $15c$, with a tail extending up to about $34c$. Within individual jets, different features appear to move with a similar characteristic velocity which may represent an underlying continuous jet flow, but we also see some stationary and even apparently inward moving features which co-exist with the main features. Comparison of our observations with published data at other wavelengths suggests that there is a systematic decrease in apparent velocity with increasing wavelength, probably because the observations at different wavelengths sample different parts of the jet structure.  The observed distribution of linear velocities is not consistent with any simple ballistic model. Either there is a rather broad range of Lorentz factors, a significant difference between the velocity of the bulk relativistic flow and the pattern speed of underlying shocks, or a combination of these options. Assuming a ballistic flow, comparison of observed apparent velocities and Doppler factors computed from the time scale of flux density variations is consistent with a steep power law distribution of intrinsic Lorentz factors, an isotropic distribution of orientations of the parent population, and intrinsic brightness temperatures about an order of magnitude below the canonical inverse Compton limit. It appears that the parent population of radio jets is not dominated by highly relativistic flows, and contrary to the assumption of simple unified models, not all sources have intrinsic speeds close to $c$.  Usually, the observed jet flow is in the general direction of an established jet. However, many jets show significant bends and twists, where the observed motions are non-radial, but are alingned with the local jet direction suggesting that the jet flow occurs along pre-existing bent channels.  In a few cases we have observed a clear change in the direction of a feature as it flows along the jet.  Radio jets which are also strong gamma-ray sources detected by EGRET appear to have significantly faster speeds than the non EGRET sources, consistent with the idea that gamma ray sources have larger Doppler factors than non gamma-ray sources. Sources at high redshift have systematically lower angular speeds than low redshift jets, consistent with standard cosmologies. ",
        "introduction": "\\label{intro}  The discovery in the mid 1960s of rapid variability in extragalactic radio sources \\citep{S65,D65,PK66} appeared difficult to explain in terms of conventional synchrotron radiation theory. These early observations showed changes by as much as 25 percent over a few weeks as well as significant day-to-day variations \\citep{PK66}. It was quickly realized (e.g., \\citealt{HBS66}), that the observed rapid variability implied such small linear dimensions, that the relativistic electron population would be rapidly extinguished by inverse Compton scattering. Later, \\cite{KP69} put these arguments on a quantitative observational basis showing that due to inverse Compton cooling, the maximum sustainable brightness temperature for an incoherent source of electron synchrotron radiation is less than $10^{12}$~K.  However, \\cite{W66} and \\cite{R66,R67} pointed out that if there is relativistic bulk motion of the emitting material, since the radiation is then beamed along the direction of motion, the apparent luminosity in that direction is enhanced, while at the same time the cross-section for inverse Compton scattering is reduced. Additionally, since the radiating source nearly catches up with its radiation, for a favorably positioned observer, the apparent transverse motion can exceed the speed of light. Early VLBI observations \\citep{W71, CCP71} showed evidence for the predicted high-velocity outflow; however, the arguments were indirect, and one had to have faith in the interpretation of the very limited radio interferometric data which did not fully sample the source structure (e.g., \\citealt{D72}). Subsequent higher quality VLBI observations confirmed the existence of superluminal motion in the well-collimated radio jets, found in the nuclei of many quasars as well as in nearby active galaxies \\citep{CLM77}.  VLBI observations provide a direct method to investigate aspects of the formation, acceleration, and propagation of extragalactic radio jets. Early studies discussing apparent jet speeds were statistically unreliable because they were largely based on the analysis of published observations of only a few tens of sources made at different times, often at only two or three epochs, by different groups, using different antenna array configurations, and different observing/data reduction strategies. The dynamic range of the images was often inadequate to identify individual features, especially for the more complex sources, and spacings in time between successive observations were frequently too long to uniquely identify and track moving jet features from epoch to epoch. Moreover, the source selection criteria for many previously published studies of superluminal motion were not well defined. A systematic study by \\cite{V95} of 81 sources from the flux density limited 6~cm Caltech-Jodrell Bank Survey indicated smaller typical speeds than in earlier reports.  This suggested that earlier studies preferentially tended to observe and/or report only sources in which rapid motion was detected or suspected.  In 1994 we began a systematic 15~GHz VLBA survey of relativistic outflows in a sample of over one hundred quasar and active galaxy radio jets. Our motivation was to study the distributions of velocities, bending, pattern motions, accelerations, and other complexities of the jet kinematics that may exist, as well as changes in the strength and morphology of features as they propagate along the jet. These observations have provided the homogeneity, resolution and dynamic range to reliably distinguish and identify individual components between different epochs. The central VLBA antennas provide short interferometric spacings that have allowed us to a) track features moving for some distance down the jet where they become more diffuse, and b) observe the continuous jet rather than just the bright features which are often referred to as ``components'' or ``blobs''.  We chose 15~GHz as a compromise between achieving the best angular resolution, and the better sensitivity and immunity from weather conditions found at lower frequencies. At this observing frequency the resolution of the VLBA is approximately 0.5~mas in right ascension and between 0.5 and 1~mas in declination. We therefore have sufficient angular resolution in many cases to resolve the 2-dimensional jet transverse to its flow, which allows for tests of some theoretical models. Also, individual features can often be recognized at significantly smaller separations from the origin. On the other hand, our observations are less sensitive to structure and motions of the more diffuse features located far out along the jet or within 0.5~mas of the core.  Our 15~GHz data represent a significant improvement over previous studies of superluminal motion in AGN. First, our sample of sources is large, and the sample membership was based on criteria other than observed superluminal motion. Second, the speeds are better determined because a) our data are well-sampled and span a longer time period, and b) our higher image resolution results in less blending of features.  In Paper I of this series \\citep{K98A}, we described the parsec-scale structure of 132 of the strongest known radio jets based on a single epoch of observation. In Paper II \\citep{Z02} we discussed the structure of an additional 39 sources. Contour maps of all of our multi-epoch observations are available on our web site \\footnote{http://www.nrao.edu/2cmsurvey/} along with kinematic and other data on each source. In this paper we discuss our multi-epoch observations made through 15~March 2001 and the kinematics derived from these observations.  Some preliminary results of our program have already been published (\\citealt{KVZ99, K00, K02, RPL02, CRH03, Kel03, K03, Z03}). We have also discussed the kinematics of several individual sources including PKS\\,1345+125 \\citep{LKV03}, NGC\\,1052 \\citep{VRK03} and 3C\\,279 \\citep{HLK03}.  In \\S~\\ref{sampledef} we describe our sample and source selection criteria, and in \\S~\\ref{observations} we discuss the details of our observational program. In \\S~\\ref{beaming} we summarize the predictions of relativistic beaming models while in \\S~\\ref{kinematics} we present the observed jet kinematics. In \\S~\\ref{beamingmodels} and \\S~\\ref{comparison} we discuss the implications of our observations for the nature of the relativistic flow. In \\S~\\ref{muz}, we discuss the angular velocity -- redshift relation, and in \\S~\\ref{summary} we summarize our results and describe our on-going observing program.  Throughout this paper we use the following cosmological parameters: $H_0 = 70$~km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_\\mathrm{m}=0.3$, and $\\Omega_\\Lambda=0.7$. ",
        "conclusions": "} \\subsection{\\it Superluminal Motion and Relativistic Beaming \\label{beaming}}  We interpret the structural changes seen in the AGN of our sample in terms of the twin relativistic jet model of \\cite{BR74}. In this framework, the bulk velocity of the relativistic flow, in units of the speed of light, $c$, is usually denoted as $\\beta_\\mathrm{b}$, and $\\gamma_\\mathrm{b}$ is the corresponding Lorentz factor defined by $\\gamma_\\mathrm{b} = (1 - \\beta_\\mathrm{b}^{2})^{-1/2}$. For a jet flow at an angle $\\theta$ to the observer's line of sight, the Doppler factor, $\\delta$, is given by $\\delta= \\gamma_\\mathrm{b}^{-1}(1 - \\beta_\\mathrm{b}\\cos{\\theta})^{-1}$.  Three important observational consequences of relativistic motion in synchrotron sources are:  a) A Doppler frequency shift of the radiation, where the ratio of observed frequency, $\\nu$, to the emitted frequency $\\nu_{e}$ is given by \\begin{equation} \\nu/\\nu_\\mathrm{e} = \\delta. \\end{equation}  b) A change in the observed transverse velocity due to the apparent time compression. The apparent transverse velocity, in units of $c$, is given by \\begin{equation} \\beta_\\mathrm{app} = \\beta_\\mathrm{p} \\sin{\\theta} /(1 - \\beta_\\mathrm{p} \\cos{\\theta}), \\end{equation} where $\\beta_\\mathrm{p}= v_\\mathrm{p}/c$ is the pattern velocity. We make this distinction between pattern and bulk velocities because jet features can move with significantly different velocity than the bulk jet flow (see, e.g., \\citealt{LB85, ZCU95, VC94}). This could be the case, for example, if the pattern motion is due to the propagation of shocks. By studying 25 core-dominated quasars selected from the literature, \\cite{VC94} showed that the simplest model, which has only one value of $\\gamma$ that is the same in all sources and also has $\\beta_\\mathrm{b}=\\beta_\\mathrm{p}$ in all sources, is not tenable. They showed that either there must be a distribution of $\\gamma$ among the sources; or that $\\beta_\\mathrm{b} \\neq \\beta_\\mathrm{p}$.  Most likely, there is both some distribution of $\\gamma$, and also a difference between $\\gamma_\\mathrm{p}$ and $\\gamma_\\mathrm{b}$  The apparent transverse velocity of an approaching component with a pattern Lorentz factor $\\gamma_\\mathrm{p}$, reaches a maximum apparent speed $\\beta_\\mathrm{p}\\gamma_\\mathrm{p}$ when the component moves at an angle $\\sin(\\theta) = 1/\\gamma_\\mathrm{p}$ to the line of sight. For the corresponding component in the receding jet, the observed velocity is $\\beta_\\mathrm{p}/2$.  c) A change in the apparent flux density, $S$, of a moving component over its stationary value, $S_0$, by a factor \\begin{equation} S/S_0 = \\delta^{x-\\alpha}, \\label{eq_ss0} \\end{equation} where $\\alpha$ is the spectral index and $x$ has a value of 2 for a continuous jet or 3 for discrete components (see e.g., \\citealt{UP95}).  The Doppler factor is sharply peaked along the direction of motion, so sources with highly relativistic jets that happen to be pointed close to the line of sight will appear strongly boosted, and hence likely to be selected in a flux-limited sample. Although the jets may be intrinsically two-sided, unless they are very close to the plane of the sky, they will appear highly asymmetric (i.e., one-sided) since the radiation from the receding jet is highly beamed away from the observer.  Equation (\\ref{eq_ss0}) has the important consequence that the strongest, most compact radio sources we observe tend to have highly relativistic jets that are aligned close to the line of sight, and they will therefore likely display superluminal motion. The most probable angle for sources selected on the basis of beamed flux density alone is close to $(2\\gamma)^{-1}$, where $\\beta_\\mathrm{app} \\sim \\gamma/2$ \\citep{VC94}. The distribution of observed speeds and flux densities is dependent on the distributions of $\\gamma$, redshift, and the intrinsic luminosity function.  This is illustrated in Figure~\\ref{rene} where we show the predicted distribution of apparent velocities for three different models.  If the effect of Doppler boosting is ignored; for example, if the observed motions are due to the propagation of shocks rather than actual bulk motion, then most jets will lie close to the plane of the sky and have apparent speeds, $\\beta_\\mathrm{app}$ close to $\\beta_\\mathrm{p}$. If we take into account the effect of Doppler boosting, then, in a flux limited sample, with a single Lorentz factor, most sources are found to lie close to the line of sight, and have an apparent velocity, $\\beta_\\mathrm{app} \\sim \\beta_\\mathrm{p}\\gamma$. In the typical situation which we consider, where $\\beta_\\mathrm{p} \\sim 1$, then in the absence of Doppler boosting, case most sources appear to have an apparent velocity close to $c$, whereas in the Doppler boosted case, most sources have an apparent velocity close to $\\gamma c$. Note that in the absence of Doppler boosting, there are virtually no sources with apparent speeds less than $\\beta c$ since this would require rare end-on orientations, The randomly oriented model also drops much more sharply with increasing $\\beta$ than the model with a range of Lorentz factors.   \\begin{figure}[t] \\begin{center} \\resizebox{1.0\\hsize}{!}{\\includegraphics[trim=0cm 0cm 0cm -0.2cm,clip]{f3.eps}} \\end{center} \\caption{\\label{rene} Predicted distributions of $\\beta_\\mathrm{app}$ when the sample has randomly oriented jets all with the same Lorentz factor $\\gamma$ (dotted line), when the sample has sources drawn from such a jet population, but based on flux density which can be enhanced by Doppler boosting (solid line), and for a Doppler boosting with a range of Lorentz factors which favors low values of the Lorentz factor (dashed line).  } \\end{figure}   \\cite{BK79} suggested that one of the observed jet components might be the stationary feature of the approaching jet at the point where it becomes optically thin, so that the flux density of this ``core'' component as well as the moving component is Doppler boosted. The observed overall spectrum is typically flat, which may due to the superposition of different parts of the jet, which each have different Doppler factors, causing the synchrotron cutoff to appear at different frequencies. Consideration of differential Doppler boosting led to the concept of unified models (e.g., \\citealt{OB82,B89,UP95}) which is usually invoked to understand the observed differences in the properties of quasars, active galaxies, and BL~Lac objects as due to the orientation of the relativistic beam and an obscuring torus with respect to the line of sight. Thus comparison of apparent jet speeds with optical classification is an effective test of these unified models and also serves to refine their parameters.   \\subsection{\\it Jet Kinematics  \\label{kinematics}}  The motion statistics collected for our sample address many physical questions related to relativistic jets. The process of their formation, that is, the initial acceleration and collimation of the flow, may be constrained by studying the speeds as a function of distance from the beginning of the visible jet (often referred to as the ``core''), and also by comparing the times when flares occur with the back-extrapolated epochs at which moving features appear to originate. Long-term multi-epoch observations can show whether there are accelerations or decelerations farther down the jets, and whether radio features follow straight or curved trajectories.  In order to discriminate between various jet models, and then to refine the relevant ones, it is important to establish whether jets exist as predefined channels, along which multiple moving features can be seen, or whether instead successive components follow different trajectories, at either the same or at different speeds. Viable models of jet formation will also need to be able to reproduce the observed Lorentz factor distribution. The moving features may in fact be patterns, caused by the propagation of shocks rather than the flow of matter.  This can be studied by comparing the Lorentz factors inferred from the motions to the Lorentz factors estimated by other means, such as variability, brightness temperature, and (possibly) the presence of X-ray and gamma-ray emission. By probing for correlations between the apparent velocities and other quantities such as the radio or X-ray luminosity, more can be learned about the physical parameters relevant for jet formation. Finally, studying the distribution of velocities as a function of optical host type is relevant to constraining unification models.  Most of the jets in our study are well collimated and are unresolved transverse to their flow, although there often is significant curvature. In some radio galaxies with two sided structure, the appearance of the source is very frequency dependent, suggesting free-free absorption in a disk or torus, probably associated with the accretion disk surrounding the central engine.   We found several jet features to have apparent negative velocities; that is, they appear to be approaching rather than receding from the core. Most of these apparently negative velocities are consistent, within the errors, with no significant motion. Observations extending over a longer time frame are needed to determine if these inward motions are real. Apparent inward motion may be produced if there is a newly emerging jet feature that is ejected from the core, and the combination is not resolved by our beam. This would cause a shift in the measured position of the centroid and a corresponding decrease with time in the apparent separation of the core and other jet features. It is also possible that the true core is not seen, possibly due to absorption, and that the only observed features are parts of a jet which are moving with different velocities. If the furthest component is moving with a slower velocity than the one closest to the obscured core, then they will appear to be approaching each other. The apparent decrease in component separation from the core could also be due to component motion away from the core along a highly curved jet which bends back toward the line of sight, so that the projected separation from the core appears to decrease with time. Finally, the moving features could be only patterns in the flow, some of which might even be moving backward. \\cite{IP96} have discussed an electron-positron jet model where an observer located close to the direction of motion will see backward-moving knots. We note that any model involving patterns must account for the the very large numbers of observed outward motions, as opposed to inward ones. For this reason, many of the simplest ``moving marquee'' models for superluminal motion have already been ruled out. In the case of 0735+178, the most distant feature appears to be moving inward, but this is probably an artifact of the complex brightness distribution whose centroid shifts when the intensity distribution changes.  \\subsubsection{\\it Velocity Distributions}  Figure~\\ref{p3histall} shows the distribution of the observed values of $\\beta_\\mathrm{app}=1.58\\times 10^{-2}\\,\\mu\\,D_\\mathrm{A}\\,(1+z)$, where $D_\\mathrm{A}$ is the angular size distance to the radio source in megaparsecs and the angular velocity, $\\mu$, is in mas per year. Figure~\\ref{p3histmoj} shows the same distributions for those jets found in the representative sub-sample described in \\S~\\ref{subsample}.  We include in these figures only those jet features which we have been able to measure with a quality code of `E' or `G' and which have measured redshifts.   \\begin{figure}[t] \\begin{center} \\resizebox{1.0\\hsize}{!}{\\includegraphics[trim=0cm 0cm 0cm -0.2cm,clip]{f4.eps}} \\end{center} \\caption{\\label{p3histall} Distribution of the apparent linear velocity in all sources with a quality code `E' or `G' and which have measured redshifts. The left hand column displays the distribution for all individual features which we have observed. Distributions in the center and right columns show only one feature per source, the brightest or the fastest respectively. Sources are divided by optical class in the second, third and fourth rows of the figure. } \\end{figure}  \\begin{figure}[t] \\begin{center} \\resizebox{1.0\\hsize}{!}{\\includegraphics[trim=0cm 0cm 0cm -0.2cm,clip]{f5.eps}} \\end{center} \\caption{\\label{p3histmoj} Same distributions as in Figure~\\ref{p3histall} except only those sources contained in our representative flux density-limited sub-sample are included. } \\end{figure}   The central and right-hand columns of Figures~\\ref{p3histall} and \\ref{p3histmoj} show the motions of only one feature per source, the brightest and the fastest respectively. The brightest feature in each jet is defined as the one with the largest peak flux density at the epoch for which the source image was presented in Papers I or II, although in nearly all cases this does not change over the period covered by these observations. These features generally have well-determined motions; however, in a few cases, the brightest jet feature does not have an `E' or `G' quality code, usually because it has a large diffuse structure. In these cases, we have substituted the brightest component which did have an `E' or `G' quality code. For the right-hand panel of Figure~\\ref{p3histall}, the fastest feature within a source is simply defined as the one with the largest (absolute) linear speed with a quality code of either `E' or `G'.  K--S tests indicate that at a 95 percent confidence level, there is no difference in the distribution of the jet speeds, independent of whether we consider the brightest, fastest, or even all the jet features within a source. We consider the brightest feature of each source to be the most representative for our analysis, since in the case of sources with weak secondary features, observations having a limited sensitivity or dynamic range might only detect the brightest jet feature.  The distributions shown in Figures~\\ref{p3histall} and \\ref{p3histmoj} are also sub-divided according to optical class. We find that the velocity distributions for the quasars, radio galaxies, and BL~Lacs are mostly concentrated in the same range ($0<\\beta_\\mathrm{app} < 15$).  However, the quasars have a tail ranging up to $\\beta_\\mathrm{app} \\sim 34$, while jets associated with active galaxies have a narrow range of velocities with $\\beta_\\mathrm{app}\\leq 6$. A K--S test confirms at the 98 percent confidence level that the quasars have a different speed distribution than the galaxies and BL~Lac objects, while the distribution of speeds for the galaxies and BL Lacs are statistically indistinguishable. Examination of Figure~\\ref{p3histall} and Figure~\\ref{p3histmoj} suggests that the distribution of quasar velocities may be bi-modal with a minimum near $\\beta_\\mathrm{app}\\sim 5$. We have compared our observed distributions with one and two Gaussian distributions and find in each case a significantly lower value of reduced $\\chi^2$ for the two component distribution for all of the quasar distributions.  However, we are reluctant to quantify this further, as the observed distributions are clearly more complex than can be represented by two Gaussians.  We have also compared the dispersion in the speed of different features within each jet, with the dispersion of the average jet speeds.  We find that the dispersion of the average jet speeds among the 50 sources with two or more features with quality factors E or G is 6.3$c$, whereas in all but four jets (92 percent) the dispersion in the speed of individual jet features is smaller than this value.  This suggests that there is an underlying flow, characteristic of each jet, that determines the speed of individual features with the jet.   \\subsubsection{\\it Apparent Speed and Apparent Luminosity}  Figure~\\ref{p3betalum} shows the relation between $\\beta_\\mathrm{app}$ on apparent VLBA core luminosity at 15~GHz. There is a distinct upper envelope to the distribution, which is very similar to that found at 5~GHz in the Caltech Jodrell Flat-spectrum survey by \\cite{V95}. In particular, the low-luminosity sources all have slow apparent speeds. If we divide the sources at the median luminosity of $L_{15}<10^{27.4}$~W\\,Hz$^{-1}$, a K--S test yields less than a 1\\% probability that speed distributions for the brightest components of the high and low luminosity sources have the same parent distribution. We note that if the apparent jet speeds that we are measuring are simply random patterns, then no envelope would be expected in Figure~\\ref{p3betalum}.  As discussed by \\cite{LM97}, the presence of this envelope does not necessarily imply that {\\it intrinsically} faint jets have low {\\it intrinsic} speeds. As discussed in \\S~\\ref{pgamma}, in practice there is likely to be a spread in Lorentz factors.  The lowest-luminosity sources in a flux-limited sample will tend to lie at low redshifts, where the co-moving volume element is small. Ignoring evolutionary effects and assuming a constant co-moving space density and a steep luminosity function, these sources will therefore be representative of the most common jets in the parent population. The probability of any of them having {\\it both} a high Lorentz factor and small viewing angle will be very small, especially if the parent population is dominated by low-Lorentz factor jets (i.e., $n(\\gamma) \\propto \\gamma^{-1.5}$ as found by \\citealt*{LM97}). This may therefore explain why none of the low-luminosity sources in Figure~\\ref{p3betalum} are highly superluminal.   \\begin{figure}[thb] \\begin{center} \\resizebox{1.0\\hsize}{!}{\\includegraphics[trim=0cm 0.5cm 0cm 0.9cm,clip]{f6.eps}} \\end{center} \\caption{\\label{p3betalum} Apparent speed of the fastest feature found in 110 jets as a function of the radio luminosity. } \\end{figure}    \\subsubsection{\\it Non-Radial and Bent Trajectories \\label{non-radial-bent}}  For most of the sources in our sample we have determined only the radial speed and position angle of the motion with respect to the assumed core component.  For those sources with sufficiently high-quality data, we also computed the full two-dimensional vector proper motion on the sky.  This was accomplished by fitting separately for the proper motion in both right ascension and declination.  These results were combined to form a vector velocity, determined by the speed, $\\mu$, and the direction, $\\phi$, on the sky relative to the radio core.  We are particularly interested in comparing the direction of motion, $\\phi$, to the mean structural position angle of the jet, $<$$\\vartheta$$>$.  Table~\\ref{non-radial} compiles the vector proper motions for all jet components in our sample which have at least five epochs of observation and for which the vector velocity is of at least $5\\sigma$ significance. These criteria guarantee that only the highest quality vector motions are used in our comparison of velocity direction to structural position angle. Of the 60 component motions that meet these criteria, we find that 20 are significantly ``non-radial'', meaning the velocity direction, $\\phi$, differs by at least $3\\sigma$ from the mean structural position angle, $<$$\\vartheta$$>$. These non-radial motions are high-lighted in bold in Table~\\ref{non-radial}.  Non-radial motions do not follow a direction which extrapolates back to the jet origin and are by definition considered non-ballistic. The occurrence of such non-ballistic motions in approximately a third of the highest quality motions we examined is striking. Assuming the observed velocities trace the underlying jet flow, this result indicates that bends in jet direction and/or jet collimation are common. The actual bends may be small, only a few degrees, but the observed bends can be large, of the order of $90\\arcdeg$, since they appear amplified by projection.  In Figure~\\ref{bent}, we show tracks in the RA -- Dec plane of several sources that have clearly defined non-radial or bent trajectories. In general, the trajectories are bent toward the next structure down the jet, whether seen directly in our maps, or in published larger scale images. This suggests that jet features may trace out a flow in a pre-existing curved channel. 1226$+$023 (3C\\,273), 1219$+$285, 1532$+$016, and 2200$+$420 (BL~Lac) are excellent examples of this trend. In 1548$+$056, the component motion appears to be transverse to the main jet direction which is known from lower resolution images, and is hinted at in our images, to lie to the north. However, the motion we observe is directed toward the small extension toward the east, and looking back toward the core, the jet ridge-line appears to curve into the component from the opposite direction.   \\begin{figure*} \\begin{center} \\resizebox{0.85\\hsize}{!}{\\includegraphics[trim=0cm 0cm 0cm 0cm,clip]{f7.eps}} \\end{center} \\caption{\\label{bent} Selected images of sources with jets that show non-ballistic component motion. Measured component positions for each epoch are superimposed on the images along with the vector motion (solid lines) in the RA -- Dec plane. Dashed lines represent the mean structural position angles, $<$$\\vartheta$$>$, for each component. } \\end{figure*}   In 3C\\,390.3 (1845$+$797) the two most prominent features appear to be moving with a similar speed of about $2.5c$, but along slightly different trajectories of $27\\arcdeg$ and $30\\arcdeg$. Neither trajectory is aligned with the direction of the core or with the narrow $60\\arcsec$ jet that points toward one of the distant hot spots in position angle $35\\arcdeg$ \\citep{AWK96}. This is a clear exception to the trend identified above.  In general we do not have sufficient data to robustly detect changes in velocity, although the high incidence of non-ballistic motion implies that such changes must occur. However, in BL~Lac, there is evidence that the moving feature traces the curved ridge-line of the jet. Also, we have previously reported an abrupt change in both speed and direction in 3C\\,279 that occurred in 1998 \\citep{HLK03}.   \\subsubsection{\\it Two-sided Sources \\label{symmetric}}  Seven sources in Table~\\ref{motions} display two-sided structure in which the jet flow appears bi-directional away from a central core. We have previously reported our results on the two-sided source NGC\\,1052 \\citep{VRK03}, where the component motion is close to the plane of the sky and multiple components move in opposite directions from the core with a velocity, $\\beta\\sim 0.26$, which is only mildly relativistic. Plots of component position vs time for all of the two sided jets are shown in Figure~\\ref{cso}. For most of these sources we were able to identify the central component at each epoch and so determine the motion of individual features away from the center. In a few cases, however, such as 1404+286 (epochs 1998.83 and 1999.55) the central core was not detected at one or more epochs, either because it was weak, or because the dynamic range was limited for that observation. In these cases we have interpolated from the positions of the outer features, to find the ``virtual'' center, and we have used that to determine the motion of the individual outer features. For the source 1323+321 even that procedure was not possible since the structure appears to be two-sided, but we do not see the core at any epoch. Moreover, our imaging of this source is not completely satisfactory due to its large angular size and the limited interferometer spacings sampled by our data.   \\subsubsection{\\it Peaked Spectrum Sources \\label{gps}}  Gigahertz Peaked Spectrum (GPS) sources are characterized by their sharp low frequency cutoff and general absence of prominent large scale structure. There is no consensus as to whether the observed spectral cutoff in these sources is due to synchrotron self absorption or to free-free absorption \\citep{SKC99}. Previous VLBI observations have suggested that GPS sources have a simple double structure morphology with little or no significant motions (e.g., \\citealt{OD98}). Thirteen of our sources are classified as GPS sources of which we, so far, have observations of eight at three or more epochs. Of these, only 1345+125 shows superluminal motion. The mean velocity observed for the GPS sources in our sample is only ($0.5\\pm1.5$)$c$. None of the GPS sources have $\\beta_\\mathrm{app}>1.5$.  The existence of a sharp peak in the spectra of these sources implies that the individual components each peak up at about the same frequency. This means that the different components probably have comparable Doppler shifts, and if the spectral cutoff is due to synchrotron self absorption they should have comparable brightness temperatures. Some GPS sources, such as 0552+398 and 0642+449 contain a single unresolved strong component, and so an interpretation in terms of synchrotron self absorption is required for any reasonable magnetic field strength. However, for other GPS sources we note the very different appearance of the core and jet components; nevertheless, they have a common cutoff frequency. Moreover, many of the peaked spectrum sources included in our study have sharply bent jet structure suggesting very different amounts of Doppler boosting and frequency shift of the self absorption peak. Free-free absorption probably plays an important role in determining the spectra of these sources. Since many of the core-jet GPS sources show little or no motion, we suggest that possibly, the jet flow in these sources is non-relativistic, or that these sources are seen at very large viewing angles, rather than that the measured velocity is the advance speed of a young double lobe radio galaxy as suggested by \\cite{OC98}.  Either of these scenarios would imply a small Doppler shift.  GPS sources do not have any extended double lobe structure which may be a consequence of the slow jets which do not carry sufficient energy to form extended radio lobes.  We note that in some cases, such as CTA\\,102 which had a peaked spectrum in the past, there are large variations in the total flux density \\citep{KKNB02} which probably reflects significant relativistic boosting. \\cite{L03} and \\cite{K04} have argued that these sources are not bona fide GPS sources but ``masquerading blazars'' which often contain bright transient jet features that dominate their radio spectrum.   \\subsubsection{\\it Sources With Extended Double Structure \\label{lobedom}}  As a result of our selection criteria, most of the sources included in our study are dominated by their flat spectrum compact structure; they are mostly identified with quasars. Some quasars, however, also have extended structure with relatively steep radio spectra, in addition to their flat spectrum compact core. We have included a few of these lobe-dominated sources in our study, although they did not meet our selection criteria. Extended structure is more common among the active galaxies. The following sources in Table~\\ref{motions} have significant double-lobe steep-spectrum extended structure: NGC\\,315 (0055+300), NGC\\,1052 (0238$-$084), 3C\\,111 (0415+379), 3C\\,120 (0430+052), M87 (1228+126), 3C\\,390.3 (1845+797), and Cygnus~A (1945+405). All are identified with galaxies containing a relatively strong AGN. All but NGC\\,1052 and 3C\\,120 are dominated by the extended steep spectrum double lobe structure even at 15~GHz. In the framework of standard relativistic beaming models, radio galaxies and lobe-dominated sources are presumed to lie close to the plane of the sky, and thus should show values of $\\beta_\\mathrm{app}\\sim 1$. However, 3C\\,111, 3C\\,120, and 3C\\,390.3 each show rather large apparent motions, and it is not clear how they fit into standard unification models.  These objects are all broad-line radio galaxies (BLRG), although 3C\\,120 has also been classified as a narrow-line radio galaxy (NLRG). The NLRG are thought to lie at rather large angles to the line of sight, and polarimetry shows that in some cases they contain a central quasar and a broad line region that are hidden by a dusty torus. The NLRG do not show strong superluminal motion. 3C\\,111, 3C\\,120, and 3C\\,390.3 do show superluminal motion and hence must be at small angles to the line of sight (less than about $11\\arcdeg$ for 3C 111). The BLRG are not well-understood, and they may have a wide range of intrinsic luminosity and orientation\\citep{DBB00}.  On the other hand, the jets of other lobe-dominated sources appear to be sub-relativistic. The powerful radio galaxy Cygnus~A contains twin jets pointed toward the distant hot spots. The approaching and receding jets appear to propagate with velocities of $0.7c$ and $0.2c$ respectively, and we find relatively slow speeds of ($0.16\\pm0.03$)$c$ in the parsec scale jet of NGC\\,315. In the case of the radio galaxy NGC\\,1052, we have reported a two-sided flow with sub-relativistic jet speeds of only $0.26c$ \\citep{VRK03} on both sides of the core. At least in these radio galaxies, the intrinsic jet flow close to the core appears to have only moderate speeds with $\\beta_\\mathrm{app}\\ll 1$.  Thus, it appears that the observed distribution of speeds and/or degree of jet asymmetry cannot be interpreted simply in terms of the orientation of a twin relativistic jet with $\\beta\\sim 1$, unless the observed speeds do not reflect that actual jet flow speed, or the radiation from a high speed inner spine is not observed due to a more slowly moving outer sheath that dominates the emission at large angles from the jet direction.  As we show in \\S~\\ref{flux} highly relativistic jets are relatively rare among the general population of radio jets.   \\subsubsection{\\it Comparison With Other Velocity Studies \\label{othervels}}  Other recent VLBA monitoring observations (e.g., \\citealt{JMMW01,JMMA01,BVT01,HOW01,VBT03}) made at other wavelengths with different resolution and with different sampling intervals are complementary to our observations and may be used to extend the range of size and time scales over which jet kinematics may be studied.  \\cite{JMMW01,JMMA01} have used the VLBA to study the motions in a sample of strong sources at 7~mm and 13~mm. These observations probe the source structure and motions on a scale about three times smaller than our 2~cm observations, but due to the decreased surface brightness at 7~mm, they are generally not able to trace the motions beyond a few milliarcseconds from the core. From the comparison of the observed motions for the sources in common to our two studies, it is possible to trace the motions of individual components over a wider range of scales than is possible from either set of observations alone.  We have compared our velocities determined at 15~GHz with those of \\cite{JMMW01,JMMA01}.  There appears to be little agreement in the individual source speeds found by the two studies. On average, the \\cite{JMMW01,JMMA01} speeds appear systematically higher, possibly because their observations were carried out at shorter wavelengths and probed jet regions closer to the core.  \\cite{BVT99,BVT01} and \\cite{VBT03} have discussed jet motions in a large sample of sources at 6~cm, taken from the Caltech-Jodrell CJ surveys covering declinations greater than $+35\\arcdeg$. These observations have lower angular resolution than ours, and are thus more sensitive to the lower surface brightness structure located downstream. \\cite{VBT03} quote a mean velocity for quasar and galaxy jets of $2.9c$ and $0.9c$ respectively. This appears be less than the values of ($7.3\\pm0.8$)$c$ and ($1.7\\pm0.8$)$c$ which we measure at 2~cm for the brightest features in each source. On the other hand, \\cite{JMMW01}, working primarily at 0.7 and 1.3~cm, find systematically faster velocities in those sources where our samples overlap. These results suggest that there is a systematic decrease in $\\beta_\\mathrm{app}$ with increasing wavelength, probably because the observations at different wavelengths, sample different parts of jet structure. In a separate paper (R.~C.\\ Vermeulen et al., in preparation) we will discuss, in more detail, the motions of those sources in common to the 6 and 2~cm samples.   \\subsection{\\it Implications for Relativistic Beaming Models \\label{beamingmodels}} \\subsubsection{\\it Ballistic Models \\label{pgamma}}  The simplest model to consider is a pure ballistic model with a common flow velocity for all sources, with $\\beta_\\mathrm{p}=\\beta_\\mathrm{b}\\sim 1$; and a flow which is intrinsically symmetric. In this simple case, observations of the apparent velocities, and the ratio of flux densities of approaching and receding components, can, in principle, be used to solve uniquely for $\\gamma$ and $\\theta$ and thus provide a test of the hypothesis that $\\beta_\\mathrm{p}=\\beta_\\mathrm{b}$. However even early VLBI data indicated that this simple model is not tenable (\\citealt{VC94,LM97}).  In a flux density-limited sample, the combined effect of available solid angle and Doppler boosting leads to a distribution of $\\beta_\\mathrm{app}$ with many values close to $\\gamma$ \\citep{VC94}. On the other hand, if Doppler boosting is not an important selection mechanism, then in most cases $\\beta_\\mathrm{app}\\sim$ 1. Figure~\\ref{rene} shows the expected distribution of apparent velocities for the case where $\\gamma_\\mathrm{p} = \\gamma_\\mathrm{b}$ (see \\S~\\ref{beaming}) along with the corresponding distribution for a randomly oriented sample. None of the distributions shown in Figure~\\ref{p3histall} or Figure~\\ref{p3histmoj} are consistent with either of these simple ballistic models as there is neither evidence for the peak at $\\beta_\\mathrm{app}\\sim\\gamma$ characteristic of the simple Doppler boosted models, nor for the sharp low speed cutoff at $\\beta_\\mathrm{app}\\sim 1$ characteristic of models which ignore Doppler boosting.  How do we reconcile the difference between the observed and predicted distributions? Either there must be spread in intrinsic velocity (e.g., \\citealt{LM97}) or a difference between the bulk flow velocity and the pattern velocity, so that there is less Doppler bias in favor of observing beams which are oriented close to the critical angle \\citep{VC94}. In Figure~\\ref{rene}, we plot a model with a spread of intrinsic velocity such that $n(\\gamma) \\propto \\gamma^{-1.5}$. This model has equal pattern and bulk velocities, and reproduces the general characteristics of the observed distribution of superluminal speeds. A more detailed analysis will be discussed by M.~L.\\ Lister et al.~(in preparation).  \\subsubsection{\\it Randomly Oriented Samples}  Based on earlier observations of apparent velocity distributions \\cite{EL90} suggested that the simple model with no Doppler bias, provided an adequate fit to the observed distributions of $\\beta_\\mathrm{app}$ with only a slight adjustment needed which was satisfied by introducing an obscuring torus. However, comparison of our data with models which do not include Doppler boosting shows poor agreement, since, as shown in Figure~\\ref{rene}, in the absence of Doppler boosting, most sources are expected to lie close to the plane of the sky where $\\beta_\\mathrm{app}\\sim 1$. This appears to be inconsistent with the tail of the velocity distribution which we find extending toward larger apparent velocities.  Following the discussion of \\cite{C90} a detailed analysis of the observed velocities shows that the probability of the sources having been picked at random from a parent population that is isotropically distributed is less than $10^{-5}$. We therefore conclude that the sources are preferentially aligned along the line of sight, as would be expected if Doppler boosting is in fact important.  \\subsection{\\it Comparison with Other Relativistic Velocity Indicators} \\label{comparison}  While the actual observation of superluminal motion in radio source jets remains the most direct way of establishing the existence of relativistic motion, comparison of the observed values of $\\beta_\\mathrm{app}$ with other velocity indicators also provides important tests of beaming models and specifically the relation between pattern and bulk flow speeds. This includes flux density variability (\\citealt{HAA92, LV99, LVW99}), maximum brightness temperature (\\citealt{KP69, R94, GD96}), the ratio of core-to-extended radio luminosity (\\citealt{OB82}), the gamma-ray luminosity (\\citealt{VM95, HBB99, MSM97, MHR01}), or observations of inverse Compton X-rays (\\citealt{GPC93}). Relativistic boosting also affects the radio source counts and luminosity functions (\\citealt{PU92, WJ97}) which provide a consistency check on beaming models.  \\subsubsection{\\it Relation Between Observed Velocity and Flux Density Variations \\label{flux}}  Flux density changes are commonly seen in superluminal sources, and their short time scale is generally taken to imply high brightness temperatures. The variability time scale, and the time scale for apparent transverse motion, are both compressed due to the forward motion, and hence we might expect to see a relation between $\\beta_\\mathrm{app}$ and the flux density variability provided that $\\beta_\\mathrm{b}$ is related to $\\beta_\\mathrm{p}$. \\cite{LV99} used variability data at 1.3~cm and 8~mm from the Mets\\\"ahovi Observatory to calculate a variability Doppler factor, $\\delta_\\mathrm{var}$, assuming an intrinsic brightness temperature characteristic of a self-absorbed synchrotron source in which the particle and magnetic energy are in equilibrium. We have recalculated their values using the cosmology given in \\S~\\ref{intro} and for different values of intrinsic brightness temperature. Since $\\delta_\\mathrm{var}$ varies inversely as the cube root of the assumed intrinsic temperature in the synchrotron plasma, $T_\\mathrm{b}^\\mathrm{int}$, $\\delta_\\mathrm{var}$ is relatively insensitive to the assumed value of $T_\\mathrm{b}^\\mathrm{int}$.   \\begin{figure}[t] \\begin{center} \\resizebox{1.0\\hsize}{!}{\\includegraphics[trim=3cm 4.5cm 4.0cm 6.25cm,clip]{f8.eps}} \\end{center} \\caption{\\label{cohen} Apparent velocity, $\\beta_\\mathrm{app}$, plotted against variability Doppler factor, $\\delta_\\mathrm{var}$, for the fastest component found in 49 sources calculated using the method of \\cite{LV99} with an intrinsic brightness temperature $T_\\mathrm{b}^\\mathrm{int}=2\\times 10^{10}$~K. The solid line shows the expected locus of points for Lorentz factor values of 25. The dotted line represents the $1/\\gamma$ cone where $\\delta_\\mathrm{var}=\\beta_\\mathrm{app}$. } \\end{figure}   In Figure~\\ref{cohen}, we plot $\\beta_\\mathrm{app}$ against $\\delta_\\mathrm{var}$ for the 49 sources in common to the Mets{\\\"a}hovi and VLBA samples. We calculate $\\delta_\\mathrm{var} =[T_\\mathrm{b}^\\mathrm{var}/T_\\mathrm{b}^\\mathrm{int}]^{1/3}$ assuming a spectral index of zero and where $T_\\mathrm{b}^\\mathrm{var}$ is the apparent brightness temperature calculated from the variability time scale, by assuming that it is limited by the size of the source divided by the speed of light. Five sources which only have components located at a bend in the jet were excluded, as they probably reflect a standing shock wave, or perhaps a stationary location in a helical jet where the flow is closest to the line of sight, and hence is boosted most strongly. In either case the measured velocity is a poor indicator of the flow velocity, and not relevant in discussing relativistic effects. The sample we use contains 5 active galaxies, 14 BL Lacs, and 30 quasars. Most points lie inside the ``$1/\\gamma$ cone'' ($\\beta_\\mathrm{app}=\\delta_\\mathrm{var}$) as they should for a flux desnity limited sample \\citep{VC94}.  \\cite{CRH03} have compared Figure~\\ref{cohen} with the results of a simulation generated by randomly picking a flux-limited sample from an isotropically distributed population with power law distributions of luminosity and $\\gamma$ and for several values of $T_\\mathrm{b}^\\mathrm{int}$ over the range $4\\times 10^{9}$ to $1 \\times 10^{11}$~K. $T_\\mathrm{b}^\\mathrm{int}\\sim 2\\times10^{10}$~K gives the best fit between the measured and simulated data. However, although there does appear to be an upper limit to $\\beta_\\mathrm{app}$ that is close to the expected locus for components with $\\gamma=25$, the detailed distribution is not well-matched to that expected from the simulation \\citep{LM97}.  Calculations of the variability Doppler factor using values of $T_\\mathrm{b}^\\mathrm{int}$ closer to the inverse Compton limit, $\\sim 10^{12}$~K, lead to distributions on the $\\beta_\\mathrm{app}\\,$--$\\,\\delta_\\mathrm{var}$ plane that are very different from those expected from the simulations. We conclude that $T_\\mathrm{b}^\\mathrm{int}$ is perhaps an order of magnitude below the inverse Compton limit.  For this application, we have used the \\textsl{fastest} feature for each source, on the grounds that these velocities should be more representative of the true flow velocities. Slower-moving components, especially those located at a bend in the jet, may be dominated by standing shock waves. Forward shock waves might also exist, and trying to understand their role is a goal of our study. Other geometries have been suggested for the jet, including a fast ``spine'' which we would preferentially see, surrounded by a slower shell. In this case, the spine would probably also control the flux variations, so that using the fastest (spine) velocity for $\\beta_\\mathrm{app}$ is appropriate.  We have also examined values of $\\delta_\\mathrm{var}$ calculated from the \\mbox{UMRAO} data base at our wavelength of 2~cm, and noted a large dispersion between Doppler factors deduced from these data and the shorter wavelength Mets{\\\"a}hovi data. Thus, the robustness of the Doppler factors calculated in this way appears to be uncertain. The Michigan data cover a longer time span, but are at longer wavelengths where individual outbursts appear to overlap in time.   \\begin{figure}[t] \\begin{center} \\resizebox{1.0\\hsize}{!}{\\includegraphics[trim=1.5cm 1.5cm 5.5cm 10.85cm,clip]{f9.eps}} \\end{center} \\caption{\\label{cohenhist} Histogram of Lorentz factors for 49 sources, calculated from $\\beta_\\mathrm{app}$ and $\\delta_\\mathrm{var}$ based on $T_\\mathrm{b}^\\mathrm{int}=2\\times 10^{10}$~K. } \\end{figure}   The Lorentz factor, $\\gamma$, in a superluminal jet is important, because it is intrinsic to the jet, whereas $\\beta_\\mathrm{app}$ and $\\delta_\\mathrm{var}$ are the observables that depend on the jet orientation. The distribution of $\\gamma$ may give information on the physics of the collimation region. From the values of $\\beta_\\mathrm{app}$ and $\\delta$ shown in Figure~\\ref{cohen} we calculate the actual Lorentz factors, shown as a histogram in Figure~\\ref{cohenhist}. We believe that the apparent deficit in the first bin of Figure~\\ref{cohenhist} for the quasars, may be a selection effect or just due to small number statistics. To be fully consistent with the analysis above, and the choice $T_\\mathrm{b}^\\mathrm{int}=2\\times 10^{10}$~K, the Lorentz factors should have a power-law distribution. However, the number of objects is too small, and the errors are too uncertain, to make a meaningful comparison. A small Lorentz factor implies a low velocity and small flux density variations, and the latter especially is less likely to be measurable. However, the calculated Lorentz factors are useful in showing that there must be a wide range of $\\gamma$ in the superluminal sources. The galaxies all have rather low Lorentz factors, while the quasars have a broad distribution up to $\\gamma\\approx 30$.  \\subsubsection{\\it Gamma-Ray Sources \\label{gamma}}  Many of the sources included in our study have been cataloged as strong gamma-ray sources according to measurements made by the EGRET detector on board the {\\it Compton Gamma-Ray Observatory} \\citep{VM95, HBB99, MSM97, MHR01,SRM03}.  It is generally thought that the gamma-ray emission occurs deep within the relativistic jet. Arguments based on size limits deduced from time variability and the cross section for pair production suggest that the gamma-ray emission, like the radio emission, is Doppler boosted \\citep{DS94}. In fact, the gamma-rays may be even more strongly beamed than the radio emission since the former generally have a steeper spectral index, $\\alpha$, and the flux density boosting varies as $\\delta^{2-\\alpha}$ for continuous jets. Also, \\cite{Dermer95} has shown that if the bulk of the gamma-rays are produced by external Compton scattering off photons associated with the accretion disk, the resulting gamma-ray emission will be boosted by an additional factor of $\\delta^{1-\\alpha}$.  If gamma-ray loud AGN do indeed have systematically high Doppler factors, then we might expect to see a different apparent speed distribution for them than for AGN that have not been detected in gamma-rays. The situation is complicated by the possibility that the gamma-ray loud jets may be seen inside the critical angle for maximizing superluminal motion, ($1/\\gamma$ radians), and therefore might have slow apparent projected speeds. However, Monte Carlo simulations based on a simple linear relationship between radio and gamma-ray luminosity (e.g., \\citealt{LM99,L99}) confirm that in a flux density limited radio sample, AGN detected by EGRET should have typically higher speeds than those that were not detected in gamma-rays.  \\cite{JMMW01} recently measured the apparent speeds of 33 EGRET-detected AGN. They found a mean value of $16c$ for the fastest component in each source and concluded that the gamma-ray sources have larger Lorentz factors than the general population of radio sources. However, \\cite{JMMW01} did not have a non-gamma-ray control sample observed in the same way and at the same wavelength with which to compare their results.  We have classified our sample into EGRET and non-EGRET sources based on the list of \\cite{MHR01} and recently modified by \\cite{SRM03}. These authors classify gamma-ray sources from the third EGRET catalog \\citep{HBB99} as ``highly probable'' and ``plausible'' AGN identifications, based on Bayesian statistics and their proximity to bright flat-spectrum radio sources. In Figure~\\ref{egret} we show the distribution of measured speeds for the EGRET detected and non-detected sources. For the brightest jet component found in the 18 EGRET sources in our representative sample we found a median speed of ($8.0\\pm1.6$)$c$ compared with a value of ($3.9\\pm1.1$)$c$ for the 53 sources with no EGRET detections. A K--S test suggests that the difference appears significant at the 90 percent level. For the purpose of this analysis we have included the two ``plausible'' EGRET sources, 2230+114 and 1156+295 as detections. Re-classifying them as non-detections sources had no effect on our results, and neither did excluding the sources with negative velocities from our analysis.   \\begin{figure}[t] \\begin{center} \\resizebox{1.0\\hsize}{!}{\\includegraphics[trim=0cm 0cm 0cm -0.1cm, angle=0,clip]{f10.eps}} \\end{center} \\caption{\\label{egret} Histogram of the brightest component speed in EGRET and Non-{\\it{}EGRET} detected sources for our representative flux density limited sample.} \\end{figure}   These results are consistent with the idea that the radio emission from gamma-ray quasars is indeed more strongly beamed than for the whole radio quasar population. However, our samples are incomplete and may therefore be biased.    \\subsection{\\it The Angular Velocity Redshift Relation \\label{muz}}  Figure~\\ref{muzfig} shows the measured values of angular velocity, $\\mu$, vs redshift, for the fastest `E' or `G' rated component found in the sources in Table~\\ref{motions}. The line represents $\\mu_\\mathrm{max}$, the fastest proper motion a source can display if it has $\\gamma = 25$. The variables are the observables, uncontaminated by modeling, and hence are of value in showing directly that (a) at all redshifts, the observed velocities are not clustered near the maximum value as expected from the simple ballistic models with a single Lorentz factor for all sources, (b) low values of $\\mu$ are seen at all $z$, and (c) high values of $\\mu$ are seen only at low $z$. This is true for galaxies, BL~Lacs, and quasars and appears inconsistent with non-cosmological interpretation of quasar redshifts as proposed by \\cite{GRB04}.   \\begin{figure}[t] \\begin{center} \\resizebox{1.0\\hsize}{!}{\\includegraphics[trim=2cm 6.5cm 4.0cm 2.85cm,clip]{f11.eps}} \\end{center} \\caption{\\label{muzfig} $\\mu$--$z$ diagram showing the distribution of angular velocity for the fastest component of the 110 sources which has a quality factor of `G' or `E'. The solid line is the maximum value of $\\beta_\\mathrm{app}$ for $\\gamma = 25$.} \\end{figure}   An early version of this plot \\citep{CBP88} was used to show that the standard paradigm for superluminal motion provided a crude upper bound to the points in Figure~\\ref{muzfig}, and that the maximum value of $\\beta_\\mathrm{app}$ in 32 sources was about 13 ($H_0=70$, $q_0=0.5$). The $\\mu_\\mathrm{max}$ line in Figure~\\ref{muzfig} is, similarly, a crude upper bound to the measured points, and shows that $\\beta_\\mathrm{app,max}\\sim 25$. Note that the cosmology used in this paper makes $\\beta_\\mathrm{app}$ somewhat larger than for a cosmology with $\\Omega_\\mathrm{m}=1$ and $\\Omega_\\Lambda=0$.  The data in \\cite{CBP88} were compiled at assorted frequencies, mostly below 15~GHz, and corresponded to measurements made at different distances from the core which may explain the difference in $\\beta_\\mathrm{app,max}$.  As we have discussed in \\S~\\ref{othervels} there is some evidence that observations at high frequencies give higher values of $\\mu$ than lower frequencies.  \\cite{VC94} and \\cite{LM97} have shown that, if Doppler boosting is important, then even rather small samples of superluminal sources will show $\\beta_\\mathrm{app,max}$ near the maximum value of the distribution of $\\gamma$, $\\gamma_\\mathrm{max}$. Hence we expect that $\\gamma_\\mathrm{max}\\sim 25$ for the sample shown in Figure~\\ref{muzfig}. This is marginally consistent with Figure~\\ref{cohenhist}, reflecting that Lorentz factors calculated according to the method in \\S~\\ref{flux} have rather large uncertainties."
    },
    "0403/astro-ph0403116_arXiv.txt": {
        "abstract": "{A new method for a simultaneous search for clusters of galaxies in X-ray photon maps and optical galaxy maps is described. The merging of X-ray and optical data improves the source identification so that a large amount of telescope time for spectroscopic follow-up can be saved. The method appears thus ideally suited for the analysis of the recently proposed wide-angle X-ray missions like DUO and ROSITA.  As a first application, clusters are extracted from the 3rd version of the ROSAT All-Sky Survey and the Early Date Release of the Sloan Digital Sky Survey (SDSS).  The time-consuming computations are performed within the German Astrophysical Virtual Observatory (GAVO). On a test area of 140 square degrees, 75 X-ray clusters are detected down to an X-ray flux limit of $3-5\\times 10^{-13}\\,{\\rm erg}\\,{\\rm s}^{-1}\\,{\\rm cm}^{-2}$ in the ROSAT energy band 0.1-2.4\\,keV. The clusters have redshifts $z\\le 0.5$. The survey thus fills the gap between traditional large-area X-ray surveys and serendipitous X-ray cluster searches based on pointed observations, and has the potential to yield about 4,000 X-ray clusters after completion of SDSS. ",
        "introduction": "Clusters of galaxies trace the peaks of the matter distribution and their dynamics on cosmic scales (see review in Borgani \\& Guzzo 2001). The corresponding cluster number density fluctuations can be measured up to Gigaparsec scales where they display a simple Gaussian random field with a high statistical significance (Schuecker et al. 2002, 2003a). Therefore, the related mean spatial abundance, the two-point correlation function, and the power spectrum give a useful summary of these fluctuations. They are well fit by a flat, low density Cold Dark Matter model without any significant contribution of dark energy others than the cosmological constant (Schuecker et al. 2003b). The data further support a simple scale-invariant biasing scheme although the details are still discussed controversal (e.g., Bahcall 1988, Collins et al. 2000, Schuecker et al. 2001, Zandivarez, Abadi \\& Lambas 2001).  The usuability of rich clusters for cosmological investigations is mainly related to their simple statistical properties and to their comparatively simple physical structure (e.g. Kaiser 1986, Sarazin 1988, Neumann \\& Arnaud 2001).  This is especially the case for X-ray clusters of galaxies where the X-ray flux limit provides a clean measure of the cluster selection function. A crucial point is the reliability of cluster mass estimates (Schindler \\& M\\\"uller 1993, Evrard, Metzler \\& Navarro 1996) where evolutionary effects could further complicate the tests (Vikhlinin et al. 2002). However, for not too distant systems the effects are found to be small (e.g. Gioia et al. 2001, Rosati, Borgani \\& Norman 2001) so that a proper calibration with detailed X-ray observations appears to be possible and is already in progress (H. B\\\"ohringer et al., in preparation).  Of fundamental importance for every precise cosmological test is the construction of large, homogeneously selected and thus representative samples.  Unfortunately, the time-consuming spectroscopic follow-up observations of the cluster candidates selected from X-ray data hampered the construction of X-ray cluster samples with sizes of the order of $10^3$ or more (see review in Edge 2003). A similar problem is expected for the recently proposed large-area X-ray missions like DUO and ROSITA. From the construction of the Northern ROSAT All-Sky (NORAS, B\\\"ohringer et al. 2000) and ROSAT ESO Flux-Limited X-ray (REFLEX, B\\\"ohringer et al. 2003) samples we learned, however, that a careful pre-selection of cluster candidates can lead to high success rates during the spectroscopic follow-up observations and can thus save a large amount of telescope time (see also B\\\"ohringer et al. 1998, 2001, Guzzo et al. 1999).  In the following section we give a description of a general likelihood method for the detection of clusters of galaxies (Sect.\\,\\ref{DETECT}) which allows a simultaneous analysis of multiwavelength data. In Sect.\\,\\ref{RASS} we apply the method to the X-ray photon distributions as given by the 3rd version of the ROSAT All-Sky Survey (RASS-3, Voges et al. 1999). A new source detection in RASS is necessary because it will turn out that -- because we combine X-ray data with much deeper optical data -- the new method can go significantly deeper compared to the standard analyses of RASS. For the measurement of the source extent, hardness ratio and other important X-ray quantities the Growth Curve Analyses (GCA) method of B\\\"ohringer et al. (2000, 2001) is used. In Sect.\\,\\ref{SDSS} we apply the same likelihood method to the galaxy distribution as given by the Early Date Release of the Sloan Digital Sky Survey (SDSS, York et al. 2000) using the $r'$ band where completeness and sample sizes are highest down to the faintest magnitudes and which serves as a `reference' passband where most methods using SDSS data are first illustrated (future applications will, however, include all five SDSS bands).  The similarity of the source detection in X-rays and cluster search in the optical allows a straightforward merging of the two results (Sect.\\,\\ref{RS}).  In a test area of 140 square degrees, 75 clusters are detected down to the nominal flux-limit of $3-5\\times 10^{-13}\\,{\\rm erg}\\,{\\rm s}^{-1}\\,{\\rm cm}^{-2}$ in the ROSAT energy band $0.1-2.4$\\,keV, that is, 5--10 times deeper than REFLEX.  The measured parameters of the clusters are described in Sect.\\,\\ref{CAT}. Extrapolated to a final SDSS area of 7,000 square degrees, our method is thus expected to yield about 4,000 X-ray clusters. To test the reliability of our method, we compare in Sect.\\,\\ref{DISCUSS} our X-ray cluster sample with various SDSS and RASS-based catalogues. Further tests with realistic simulations are postponed to a second paper.  It will be seen that a detailed analysis of the RASS and SDSS data requires massive computing power. Moreover, our method allows the inclusion of further high-resolution maps from, e.g., temperature anisotropies of the cosmic microwave background radiation, centimeter or millimeter (Sunyaev-Zel'dovich) maps etc. The combined analysis of such large databases is a typical example where virtual observatories can help with a fast and easy data handling and a distribution of the time-consuming computations over many computer systems (grid computing). The present investigation was thus planned already from its beginning to be performed within the German Astrophysical Virtual Observatory (GAVO).  All computations assume a pressureless Friedmann-Le\\-ma\\^{\\i}\\-tre world model with the Hubble constant $H_0$ in units of $h=H_0/(100\\,{\\rm km}\\,{\\rm s}^{-1}\\,{\\rm Mpc}^{-1})$, the normalized cosmological constant $\\Omega_\\Lambda=0.73$, and the density parameter $\\Omega_m=0.27$, but the resulting cluster catalogues do not depend on the chosen values of the cosmic energy densities. ",
        "conclusions": "\\label{DISCUSS}  \\begin{figure*} \\vspace{0.0cm} \\centerline{\\hspace{-7.5cm} \\psfig{figure=lz1.ps,height=8.0cm,width=8.0cm}} \\vspace{-8.0cm} \\centerline{\\hspace{10.5cm} \\psfig{figure=lz2.ps,height=8.0cm,width=8.0cm}} \\vspace{0.0cm} \\caption{\\small X-ray luminosity vs. redshift for the total cluster sample (upper left), for clusters not in the Goto et al. (2002a) sample (lower left), for clusters not in the Bahcall et al. (2003) sample (upper right), and for the clusters neither in Goto et al. and Bahcall et al. (lower right). The redshifts used are estimated from the SDSS data with the matched filter described in Sect.\\,\\ref{SDSS}.} \\label{FIG_LZ} \\end{figure*}   The present paper describes a method to search for clusters of galaxies by combining data obtained in different energy ranges. The main goal is to reach lower flux limits and to improve the reliability of a first identification of the sources.  As an example, we combined X-ray data from RASS-3 and optical data from SDSS. The SDSS galaxy sample reaches much higher redshifts than RASS and can thus be used to guide a cluster detection in RASS down to lower X-ray flux limits. We could show that a detection of X-ray clusters in RASS-3 is possible down to a nominal X-ray flux limit of $3-5\\times 10^{-13}\\,{\\rm erg}\\,{\\rm s}^{-1}\\,{\\rm cm}^{-2}$ in the ROSAT energy band 0.1-2.4\\,keV. This limit is about 5--10 times deeper than REFLEX and gives a first qualitative impression how much deeper one can get with the new technique.  In the following we discuss the results which are already indicated by the analysis of the test area. Table\\,2 shows that only 40 out of our 75 X-ray clusters (53\\%) are listed as X-ray sources in the ROSAT Bright and Faint Source (RBF) catalogues. The overlap can be increased possibly to 73\\% when a preliminary source list obtained with RASS-3 is used. It is thus not only the X-ray flux limit which can be decreased but also the X-ray completeness which can be increased when RASS and SDSS data are combined.  At the nominal X-ray flux limit only 20\\% of the 75 clusters show a significant X-ray extent as measured with the GCA method (Fig.\\,\\ref{FIG_HREX}, Tab.\\,3). This small fraction is mainly caused by the comparatively small number of source counts (half of the sample has source counts between 5 and 10). Usually, more than 30 source counts are necessary for the GCA method for a significant detection of a source extent -- the precise numbers do, however, also depend on the intrinsic extent of the source.  One could worry about a significant contamination by AGN. However, only 12\\% of the clusters have a known nearby AGN (Tab.\\,3) and only 5\\% of the X-ray clusters have an X-ray hardness ratio measured with GCA which is consistent with an AGN (Fig.\\,\\ref{FIG_HREX}). The fraction of contaminated clusters should thus be not very large, possibly around 10\\%.  Before we discuss the direct comparisons with published SDSS-based cluster samples we want to note that it would not make much sense to test the number of SDSS clusters which are not present in our sample because SDSS-based cluster surveys are expected to reach much higher redshifts and less rich systems compared to RASS-based cluster surveys. We are thus mainly looking for X-clusters which are not found in SDSS samples.  The comparison of our X-ray cluster sample with public SDSS cluster samples yields the following results: 13 of our 75 X-ray clusters (17\\%) are not in the catalogue of Goto et al. (2002a) and 23 clusters (31\\%) are not in the catalogue of Bahcall et al. (2003). At the comparatively bright X-ray fluxes $S\\ge 5.0\\cdot 10^{-13}\\,{\\rm erg}\\,{\\rm s}^{-1}\\,{\\rm cm}^{-2}$ where spurious detections are unlikely, still 11 clusters (15\\% of our sample) are not in the Goto et al. sample and 19 clusters (25\\% of our sample) are not in the Bahcall et al. sample.  However, when we combine the catalogues of Goto et al. and Bahcall et al. then the incompleteness of the optical sample decreases to a level of 8\\%. The corresponding 6 X-ray clusters which are not in the two SDSS-based cluster catalogues are marked as BG in the last column of Table\\,2 and have the spectroscopic redshifts $z=0.0637$, 0.0650, 0.0663, 0.0425, 0.0206 and 0.1822. Most of these missing clusters (with one exception) are nearby and less rich systems. This is also seen in the X-ray luminosity/redshift plots shown in Fig.\\,\\ref{FIG_LZ}. In our X-ray cluster sample we have 12 clusters with redshifts $z<0.07$ and only 7 (58\\%) are in the combined optical catalogues of SDSS. One should note that the 6 X-ray clusters missing in the optical catalogues on our test area of 140 square degrees could lead to an extrapolated number of 300 clusters over a total SDSS survey area of 7000 square degrees. We are thus discussing here a significant fraction of nearby systems.  The difference between X-ray and optical cluster selection becomes apparent in Fig.\\,\\ref{FIG_NZ}, where the redshift distribution of the total sample (unfilled histogram) and the distribution of clusters which are members of both the Goto et al. and the Bahcall et al. samples are compared. Ignoring the redshifts $z>0.3$ where the sample of Bahcall et al. is incomplete (by construction), the redshift distributions are significantly different as seen by the strong deficiency of nearby clusters in the combined Goto plus Bahcall catalogue. Whereas the cluster counts of the combined optical catalogue increase with redshift -- which is typical for volume-limited samples -- the X-ray sample shows a steady decrease with redshift -- which is typical for flux-limited samples. Numerical simulations and comparisons with deep XMM and Chandra observations will give further information about the final incompleteness level which can be reached with our method.  We want to close the discussion of the sample with a few general remarks on the quite deep X-ray flux limit achieved from the combination of RASS and SDSS data. The resulting `superfaint' flux limit -- though optimal for the detection of high-redshift clusters and for compiling huge sample sizes -- yields X-ray cluster candidates which consists in many cases of only 5--10 RASS X-ray source counts. The corresponding formal flux errors are of the order of 30 and 40\\%. A careful evaluation of the resulting accuracies achievable in cosmological tests with such samples is thus necessary.  In addition, more work is needed to better calibrate and test our method: (1) The analysis of a larger test area (Data Release 2) is expected to give significantly smaller statistical errors of the cluster number counts which is necessary for a better calibration of the likelihood threshold of the optical data. (2) Tests are under way which quantify the effect an optical sample like SDSS has on the selection of X-ray clusters and vice versa. In Popesso et al. (2004) the correlations between basic optical and X-ray cluster properties are studied and more tests will follow. One of their conclusions is that the $z'$ SDSS photometric band and not the $r'$ band used here gives the strongest correlations between optical and X-ray data, and should thus be the preferred band when we want to work with the smallest bias caused by including optical data in the detection process. (3) The numerical simulations in Schuecker \\& B\\\"ohringer (1998) have to be performed under the specific conditions of the present survey in order to test selection effects caused especially by the small number of X-ray photons and surface brightness. (4) A direct comparison with ROSAT, XMM-Newton, and Chandra pointed observations will give model-independent estimates of the flux and completeness limits.   \\subsection{Role of Virtual Observatories}\\label{ROLE}  \\begin{figure} \\vspace{-2.2cm} \\centerline{\\hspace{1.0cm} \\psfig{figure=nz.ps,height=10.0cm,width=10.0cm}} \\vspace{-1.1cm} \\caption{\\small Redshift histogram of the clusters with spectroscopic redshifts. Filled steps show the redshift distribution of the clusters which are in common with Goto et al. (2002a) and Bahcall et al. (2003).} \\label{FIG_NZ} \\end{figure}   The present investigation gives a good illustration of the benefits achievable from the combination of high-quality optical and X-ray data. The combination of the two large databases allows us to fill the gap between nearby wide-angle and deep serendipitous X-ray cluster surveys with already existing data. The combination increases the completeness of the samples and reduces the time-consuming spectroscopic follow-up observations. This method is also expected to be helpful when samples of the order of several $10^4$ X-ray cluster candidates expected from future wide-angle X-ray surveys like DUO or ROSITA must be identified.  However, the quality and completeness of the catalogues depends also on the model used to detect and characterize the galaxy clusters. In the present investigation only one set of values of model parameters (cluster core radius, cutoff radius, Schechter characteristic magnitude and faint end slope) was used to develope and illustrate the method. More realistic applications should test different sets of model parameters and use a finer sampling grid for cluster detection. In addition, cluster radii are known to increase with mass so that some kind of iterative procedure is necessary to find the optimal model for each cluster.  Further applications should also include the detection of clusters in so-called Sunyaev-Zel'dovich maps obtained in the centimeter or millimeter range. Such kinds of maps will be provided by, e.g., the PLANCK mission, and will surely increase the reliability of the resulting catalogues.  We are thus left with the conclusion that the quality, the completeness limit, the detection limit, and the telescope time for source identification can be optimized when enough computer power to sample the high-dimensional parameter spaces is available. Moreover, a fast and easy handling of data is needed to work with very large databases probably located at different physical nodes. {\\it Virtual observatories} like the GAVO provide the infrastructure for such kind of investigation and are surely necessary to make full and efficient use of already existing and future observed data.  GAVO provides an all-sky X-ray event file including many characteristics of the invididual photons. The photons are retrieved from the archive with a general cone-search algorithm. We are currently in the process to implement the cluster search described here within the GAVO grid. This final step is crucial for the analyses of larger survey areas with more cluster models and a finer sampling grid of the detection cells. This will allow us to analyse in a much more efficient manner the SDSS Data Release 2 sample which will become available soon and RASS-3.  The method described in the present investigation needs about 1 hour computing time per square degree on a standard workstation with one processor. Running the method in a controlled manner over a grid of many computers is thus essential for a proper analysis within reasonable times. Our plans include the provision of the cluster (source) detection algorithm as a service of GAVO so that the procedures can be used in a user-friendly manner by a larger community of interested researchers."
    },
    "0403/astro-ph0403099_arXiv.txt": {
        "abstract": "We present the results of a high-resolution spectroscopic monitoring program of the eclipsing pre--main-sequence star KH~15D that reveal it to be a single-line spectroscopic binary. We find that the best-fit Keplerian model has a period $P = \\pn$ days, which is nearly identical to the photometric period. Thus, we find the best explanation for the periodic dimming of KH 15D is that the binary motion carries the currently visible star alternately above and below the edge of an obscuring cloud. The data are consistent with the models involving an inclined circumstellar disk, as recently proposed by Winn et al. (2004) and Chiang \\& Murray-Clay (2004). We show that the mass ratio expected from models of PMS evolution, together with the mass constraints for the visible star, restrict the orbital eccentricity to $0.68 \\le e \\le 0.80$ and the mass function to $0.125 \\le F_M/\\sin^3i \\le 0.5$~\\msun. \\\\ \\\\ \\noindent {\\em Subject Headings:} stars: pre--main-sequence --- stars: individual (KH~15D) --- binaries: spectroscopic --- techniques: radial velocities --- circumstellar matter ",
        "introduction": "\\kh is a K6-K7 pre--main-sequence star that exhibits dramatic photometric variability \\citep{kearns98}. Every 48 days, the star's brightness dims by 3.5 magnitudes and remains in this faint state for nearly half of the photometric period. These deep brightness minima are accompanied by a slight blueing of the star's color indices \\citep[][hereafter He02]{herbst02}, little or no change in spectral type \\citep{hamilton01} and an increase in linear polarization \\citep{agol04}. This implies that the star is completely eclipsed by an optically thick, extended collection of dust grains, possibly in the form of a circumstellar disk. If this is the case, the serendipitous alignment of the \\kh star/disk system may provide insights into the evolution of young stars and their interactions with their circumstellar environments.  While \\kh has periodic eclipses, it cannot be an ordinary eclipsing binary because of the long duration of minimum light. Recent theories postulate that there is, nonetheless, a currently unseen binary companion to the visible star. In a study of archival photographic plates, \\citet[][hereafter JW04]{johnson04} discovered that the historical light curve of \\kh is similar to the modern light curve but appears to be diluted by light from a second star. Motivated by these findings, \\citet[][hereafter W04]{winn04} constructed a model composed of a binary system with the orbital plane inclined with respect to the edge of an optically thick screen. As the two stars orbit one another, the reflex motion carries one star alternately above and below the edge of an opaque screen, causing the eclipses. The long-term evolution of the light curve is reproduced by allowing the screen to move slowly across the binary orbit---a feature of the model that led the authors to envision the screen as a precessing, circumbinary disk. A similar model is proposed by \\citet[][hereafter CM04]{chiang04} who also envision the opaque screen as an inclined, precessing circumbinary disk, or ``ring,'' with an inner edge truncated by tidal interactions with the binary and an outer edge possibly shepherded by a planetary companion. Both models provide explanations of the unique features of KH~15D's light curve. The W04 model makes quantitative predictions about the orbital parameters of the binary system, while the CM04 model provides a physical description of the circumbinary ring. \\begin{deluxetable*}{rlccccc} \\tablecaption{Spectroscopic Observations of \\kh \\label{observations}} \\tablewidth{0pt} \\tablehead{ \\colhead{UT Date} & \\colhead{Telescope/} & \\colhead{$\\Delta\\lambda$ [\\AA]} & \\colhead{R} & \\colhead{J.D.$- 2.4\\times 10^6$} & \\colhead{$v_r$ [km/s]} \\\\ \\colhead{} & \\colhead{Instrument} & \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{} & } \\startdata 25 Oct 2002 & Keck/HIRES & 3900-6200 & 70,000 & 52572.574 & $1.7 \\pm 0.2$\\tablenotemark{a}\\\\ 28 Oct 2002 & Keck/HIRES & 3900-6200 & 70,000 & 52575.511 & $3.0 \\pm 0.3$\\tablenotemark{a}\\\\ 16 Dec 2002 & Magellan/MIKE  & 4500-6800 & 45,000 & 52624.738 & $3.1 \\pm 0.2$\\tablenotemark{b}\\\\ 14 Jan 2003 & Keck/HIRES & 3900-6200 & 70,000 & 52653.468 & $3.3 \\pm 0.2$\\tablenotemark{a}\\\\ 8 Feb 2003  & Keck/HIRES & 4200-6600 & 70,000 & 52678.400 & $9.0 \\pm 0.5$\\tablenotemark{b}\\\\ 9 Feb 2003  & Keck/HIRES & 4200-6600 & 70,000 & 52679.410 & $11.5 \\pm 0.5$\\tablenotemark{b}\\\\ 9 Mar 2003  & Magellan/MIKE  & 4900-8700 & 22,000 & 52707.511 & $0.8 \\pm 0.2$\\tablenotemark{b}\\\\ 3 Nov 2003  & Keck/HIRES & 4200-6600 & 30,000 & 52946.584 & $1.2 \\pm 0.4$\\tablenotemark{b}\\\\ 4 Nov 2003  & Keck/HIRES & 4200-6600 & 30,000 & 52947.594 & $1.8 \\pm 0.3$\\tablenotemark{b}\\\\ 4 Jan 2004 & McDonald/CE & 5600-6900 & 30,000 & 53008.762 & $1.4 \\pm 0.6$\\tablenotemark{b}\\\\ 5 Jan 2004 & McDonald/CE & 5600-6900 & 30,000 & 53009.852 & $1.7 \\pm 0.4$\\tablenotemark{b}\\\\ 10 Jan 2004 & Keck/HIRES & 3900-6200 & 70,000 & 53014.320 & $5.4 \\pm 0.3$\\tablenotemark{a}\\\\ 10 Jan 2004 & McDonald/CE & 5600-6900 & 30,000 & 53014.760 & $5.8 \\pm 0.4$\\tablenotemark{b}\\\\ 11 Jan 2004 & Keck/HIRES & 3900-6200 & 70,000 & 53015.506 & $7.0 \\pm 0.2$\\tablenotemark{b}\\\\ 9 Feb 2004  & Keck/HIRES & 3900-6200 & 70,000 & 53044.834 & $1.8 \\pm 0.2$\\tablenotemark{a}\\\\ 10 Feb 2004 & Keck/HIRES & 3900-6200 & 70,000 & 53045.828 & $1.3 \\pm 0.2$\\tablenotemark{a} \\enddata \\tablenotetext{a}{Multiple reference stars were used to calculate the radial velocity. The reported velocity is the mean velocity measured from the various target-reference star pairs. The estimated uncertainty is the standard deviation of the mean velocity. See Table \\ref{multiobs} for a listing of individual reference stars and radial velocity measurements for each target-reference star pair.} \\tablenotetext{b}{Only one reference star was used in the radial velocity measurement. The reported velocity is the mean velocity from all echelle orders. The estimated uncertainty is the standard deviation of the mean velocity.} \\end{deluxetable*}  A fundamental question that has not yet been answered is whether \\kh~is a single or multiple stellar system. He02 \\citep[see also][hereafter Ha03]{hamilton03} first searched for evidence of orbital companions using high-resolution VLT spectra and reported a radial velocity change of $+3.3 \\pm 0.6$ km s$^{-1}$ over two widely spaced epochs. However, one of the measurements was made during egress when there was a strong possibility of contamination from scattered light or line profile distortions from a sharp-edged cloud only partially occulting the stellar disk.  Over the past two years we have conducted a high-resolution, multi-site spectroscopic monitoring campaign to determine whether or not \\kh exhibits orbital motion indicative of a multiple system. Here we present the results of our study, which show that \\kh undergoes significant radial velocity variations. The variations are consistent with a binary companion with an orbital period equal to the 48-day photometric period. In \\S \\ref{data} we summarize our observations and reduction procedures. The radial velocity measurements and best-fit Keplerian orbital parameters are presented in \\S \\ref{analysis}. We conclude in \\S \\ref{discussion} with a discussion of our findings and implications for existing models of the \\kh eclipse mechanism. ",
        "conclusions": "\\label{discussion}  The results of our spectroscopic monitoring campaign show that \\kh is a spectroscopic binary system. The observed radial velocities are consonant with a stellar companion with an orbital period equal to the photometric period. We now discuss the implications of the binarity of \\kh for current models of the photometric variability mechanism.  Existing models of the \\kh eclipse mechanism fall into two classes differentiated by whether it is the orbital motion of the star or a feature of a circumstellar disk that causes the photometric variability. The first class of models posits the existence of a single star surrounded by a circumstellar disk containing a nonaxisymmetric density enhancement or alternatively, a warp. As the disk feature orbits the star with a 48-day period, it periodically blocks the line-of-sight to the stellar surface. Based on their spectropolarimetric observations that showed an increase in polarization during minimum light, \\citet{agol04} developed a model of a warped disk with an extended atmosphere and obtained a reasonable fit to the 2001-2002 light curve. A similar analysis was performed by \\citet{barge03} using a model involving a large dusty vortex.  While these single--star models are able to match the light curve of \\kh at one point in time, they do not adequately explain the temporal changes in the observed light curve. Perhaps more importantly, it is not clear if a binary companion to \\kh is compatible with an eclipsing disk feature. If a density enhancement in a circumstellar disk is responsible for the 48--day photometric period, the feature must orbit the central star at a distance of $\\sim 0.22$~AU, assuming a Keplerian orbit. This geometry would restrict the stellar companion to orbit either at a distance less than 0.22 AU or beyond the outer extent of the circumstellar disk. However, since our best-fit orbit yields a period of $\\pn \\pm 0.01$ days and $a\\sin i = 0.21$ AU (assuming $e = 0.74$) for the binary companion, neither of these scenarios seems plausible. Thus, it does not seem possible for both an eclipsing disk feature and a stellar-mass companion to coexist with the same orbital period.  In the second class of models, proposed by W04 and CM04, there exists an unseen binary companion of comparable mass to the visible K-type star seen today. Surrounding the two stars is a circumbinary disk viewed nearly edge-on. The orbital plane of the two stars is tilted at a small angle with respect to the disk plane and the present day light curve is produced as the reflex motion of Star A carries it above and below the disk plane.  Independent evidence for the binary nature of \\kh has emerged from studies of archival photographic plates. In a study of photographic plates obtained from Asiago Observatory, JW04 show that the apparent magnitude of \\kh was variable from 1968 to 1983, but the light curve from this epoch was markedly different from the one observed today. The bright state was nearly a factor of 2 brighter in the past and the eclipse depth was a factor of 5 shallower. These findings can be explained by invoking the presence of a second star that was visible in the past but is unseen today (JW04).  Motivated by these findings, W04 constructed a model of an eccentric binary with a fraction of the orbital plane obscured by an opaque screen (or circumbinary disk) and found a quantitative orbital solution by fitting simultaneously to the 2001-2002 photometry of He02, the historical photometry of JW04 and the radial velocity measurements of Ha03. Based on a preliminary investigation of the radial velocities presented here, CM04 independently used physical arguments to propose a similar model of an eccentric binary surrounded by a nodally precessing circumbinary ring. As the ring precesses, the light curve gradually changes from the one recovered from the archival plates to the shape seen today. Perhaps the greatest advantage of the two-star models is their ability to explain not only the present-day light curve, but also its evolution over the past half century.  In addition to explaining the photometric phenomenology of KH~15D, the two-star models make predictions about the nature of the binary orbit. Both CM04 and W04 predict that (1) periastron passage occurs during minimum light, (2) the orbital companion has a mass comparable to the currently visible star, and (3) the binary orbit is highly eccentric. These predictions are precisely what we find from our orbit solution.  The W04 model makes additional, quantitative predictions about the orbital parameters of the binary. For an assumed fixed period of 48.35 days, the model produces a velocity semi-amplitude of 27.5~\\kms, eccentricity $e = 0.7$, a mass ratio $M_A/M_B = 1.6$, argument of pericenter $\\omega = -7^\\circ .2$ and a center of mass radial velocity $\\gamma = +15.5$~\\kms. As can be seen in Table \\ref{parameters}, the W04 Model 1 predictions agree well with our orbital solution with fixed eccentricity, $e = 0.74$ (see \\S \\ref{constrain}). We note here that our center of mass velocity for the binary system of $\\gamma = +14.7 \\pm 0.9$\\kms~clearly rules out W04 Model 2, which predicts $\\gamma = +5.7$~kms and negative velocity at periapse.  The W04 model also predicts the inclination of the binary orbit, a property of the system that our radial velocities do not. W04 Model 1 predicts $i = 84^\\circ.6$ or $\\sin i = 0.996$. Similarly, using the geometry of the circumbinary ring proposed by CM04, an upper limit on the inclination of $i < 80^\\circ$ can be assigned to the binary orbit based on the time lag between periastron passage and mid-eclipse (Chiang, private communication).  We note that the W04 model has the peculiar feature that the less massive star is the more luminous star. Using Equation \\ref{ma}, the values of $P$, $K$ and $e$ produced by the W04 model lead to a mass function $F_M = 0.038$~\\msun~and $M_A = 0.41$~\\msun. This mass is significantly less than the lower limit of 0.5~\\msun~measured by \\citet{park00} and the value of 0.6~\\msun~measured by \\citet{flaccomio99}. Figure \\ref{mvse} shows a plot of $M_A$ as a function of $e$ for the W04 mass ratio $M_B/M_A = 1.6$. For our best-fitting model parameters, only eccentricities between 0.58 and 0.65 yield a mass of A between 0.5 and 1.0 \\msun. Therefore, for a mass ratio of 1.6, the eccentricity reported by W04 ($e = 0.7$) is larger than the value allowed by our radial velocity measurements assuming $0.5 \\le M_A \\le 1.0$~\\msun. However, this discrepancy is not too surprising since the W04 model used only two radial velocity measurements. It is also important to note that the mass limits of star A are derived by placing \\kh on a theoretical HR diagram and are therefore subject to the accuracy of the PMS evolutionary model employed. In all other features of the \\kh binary system, there is a remarkable agreement between the W04 model and the orbital solution calculated from the radial velocities. Because of the strong evidence of a second star from our radial velocity measurements and the historical photometry, we find the two-star class of model to be the most compelling explanation of the \\kh photometric variability mechanism.  However, the case of the ``winking star'' is still far from closed. A key missing aspect of the two-star models is direct detection of a circumbinary disk around KH~15D. He02 report a lack of near-IR excess and a null detection at millimeter wavelengths. CM04 state that such findings are consistent with a circumbinary ring having an inner radius of $\\sim1$~AU that is tidally truncated by the central binary, and an outer radius of $\\sim5$~AU that is possibly shepherded by an as yet unseen planet. They predict mid-infrared fluxes that are observable with the Spitzer Space Telescope. Clearly such observations will be vital in further development of models of the KH~15D system."
    },
    "0403/hep-ph0403014_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "In this Letter we have studied a mechanism for CP violation (originally proposed by Dolgov \\cite{Dolgov}) which is effective in the early Universe but could shut off at late times. A key feature of our mechanism is that it does {\\it not} involve any new CP-violating phases in the Lagrangian. Instead, in a model containing several scalar fields, initial conditions will lead to an asymmetry in the phases of these fields. CP-violating processes then arise as these quantum fields relax dynamically towards their symmetric ground state values. The initial displacement of these fields is completely natural in the context of early Universe cosmology.  There are many possible applications for this mechanism. In this paper, we have concentrated on its contribution to leptogenesis and baryogenesis. Our calculation is applicable in the phase when the fields are rolling. This rolling phase will start when the Hubble constant drops to a value comparable to the mass of the scalar fields. It is at this time in the cosmological evolution that CP violation is most efficient. After the fields have relaxed to their vacuum values, our CP violation mechanism turns off. We plan to discuss more details, in particular applications to concrete baryogenesis models, in a future publication \\cite{Balaji}.  Note that string cosmology and brane world scenarios may provide natural settings for the origin of the scalar fields required for our mechanism (e.g.\\ see Ref.~\\cite{Gian} for a recent paper on how scalar fields from brane world scenarios can play a new role in spontaneous baryogenesis).  \\bigskip \\noindent {\\bf Acknowledgements}: We thank Guy Moore and Cliff Burgess for useful discussions. One of us (RB) is also grateful to Marco Peloso for interesting conversations. RB is supported in part by the US Department of Energy under Contract DE-FG02-91ER40688, TASK~A. He thanks the Fields Institute for Mathematical Sciences and the Perimeter Institute for hospitality and support during the final stages of this project. The work of KB, TB and DL is supported by NSERC (Canada) and by the Fonds de Recherche sur la Nature et les Technologies du Qu\\'ebec."
    },
    "0403/astro-ph0403521_arXiv.txt": {
        "abstract": "A detailed X-ray analysis of an \\XMM observation of the high-redshift (z=0.89) galaxy cluster ClJ1226.9$+$3332 is presented. After careful consideration of background subtraction issues, the X-ray temperature is found to be $11.5^{+1.1}_{-0.9}\\keV$, the highest X-ray temperature of any cluster at $z>0.6$. The temperature is consistent with the observed velocity dispersion. In contrast to MS1054-0321, the only other very hot cluster currently known at $z>0.8$, ClJ1226.9$+$3332 features a relaxed X-ray morphology, and its high overall gas temperature is not caused by one or several hot spots. The system thus constitutes a unique example of a high redshift (z$>$0.8), high temperature (T$>$10 keV), relaxed cluster, for which the usual hydrostatic equilibrium assumption, and the X-ray mass is most reliable.  A temperature profile is constructed (for the first time at this redshift) and is consistent with the cluster being isothermal out to $45\\%$ of the virial radius. Within the virial radius (corresponding to a measured overdensity of a factor of 200), a total mass of $1.4\\pm0.5\\times10^{15}M_\\odot$ is derived, with a gas mass fraction of $12\\pm5\\%$ (for a \\LCDM cosmology and H$_0$=70 km s$^{-1}$ Mpc$^{-1}$). This total mass is similar to that of the Coma cluster. The bolometric X-ray luminosity is $5.3^{+0.2}_{-0.2}\\times10^{45}\\ergps$. Analysis of a short {\\it Chandra} observation confirms the lack of significant point-source contamination, the temperature, and the luminosity, albeit with lower precision. The probabilities of finding a cluster of this mass within the volume of the discovery X-ray survey are $\\sim8\\times 10^{-5}$ for $\\Omega_M=1$ and $0.64$ for $\\Omega_M=0.3$, making $\\Omega_M=1$ highly unlikely.  The entropy profile suggests that entropy evolution is being observed. The metal abundance (of $Z=0.33^{+0.14}_{-0.10}\\Zsol$), gas mass fraction, and gas distribution are consistent with those of local clusters; thus the bulk of the metals were in place by z=0.89. ",
        "introduction": "\\label{sect:intro} Massive galaxy clusters form from the high-sigma tail of the initial cosmological density distribution. As a result they are rare, but also very powerful probes of cosmology. Given an assumed initial density distribution, the properties of the massive cluster population can be predicted under alternate cosmologies, and those predictions tested with observations. The predictions of different cosmologies diverge with redshift, making high-redshift, massive clusters the most useful objects to distinguish between them.  X-ray observations of galaxy clusters provide a useful means of measuring their properties. The intra-cluster gas is extremely luminous in X-rays, and measurements of the gas temperature and density distributions allow the total mass of the system (the property most directly related to cosmological predictions) to be inferred. This inference requires, however, that the intra-cluster medium (ICM) be in hydrostatic equilibrium. If, as believed, clusters form through a series of hierarchical mergers, then this will only be the case some time after the last merger. Thus the most useful objects to study for the purpose of constraining cosmological models in this way, are high-redshift, massive, relaxed clusters of galaxies. These are extremely rare.  However the question of when a cluster can be considered to be relaxed is something of a contentious issue. In a study of $368$ low$-z$ clusters observed by \\Einstein, \\citet{jon99} found $\\approx40\\%$ to have substructure in their X-ray images. This fraction is likely to be an underestimate, as \\Einstein was unable to resolve small scale substructure. More recently, the high resolving power of \\Chandra has revealed substructure in clusters that were previously considered to be relaxed, such as A1795 \\citep{fab01,mar01}, and MS 1455.0+2232 \\citep{maz02a}. On the other hand, X-ray derived masses of clusters that appear relaxed in \\Chandra observations have been found to agree well with independent weak lensing mass measurements, at least in the inner regions \\citep[][and references therein]{all01b}.  While \\Chandra can accurately probe the gas properties in the central regions of clusters, the strength of \\XMM lies in its large collecting area, which allows it to trace the gas density and temperature structure out into the low surface-brightness emission at large radii, even at high redshifts. This minimises the uncertainties involved in extrapolating these properties out to the virial radius when deriving the total mass of the system. The mass composition of massive galaxy clusters (\\eg the baryonic to total mass fraction) is believed to be representative of the universe as a whole, due to their large size \\citep[\\egc][]{all02a}. Thus by directly observing the ICM out to large radii, one obtains a more representative measurement of these properties.  In June 2001, \\XMM made a $30\\ks$ observation of galaxy cluster ClJ1226.9$+$3332, one of the most distant, luminous clusters found in the WARPS X-ray selected survey \\citep{sch97,jon98a,ebe00,per02}. The cluster was positioned $\\approx4\\arcm$ off-axis in order to investigate other candidate clusters in the field, to be described in a future paper. The discovery \\ROSAT data indicated a high X-ray luminosity, but were insufficient to accurately determine morphology or temperature. Optical follow up found the cluster's redshift to be $0.89$ \\citep{ebe01}, corresponding to a look back time of over half of the age of the universe, and Sunyaev-Zel'dovich effect imaging confirmed that the cluster is both hot and massive \\citep{joy01}. We present here the results of a detailed analysis of the \\XMM data. There is also a fairly short ($10\\ks$) archived \\Chandra observation of ClJ1226.9$+$3332, which has been examined by \\citet{cag01}. We have analysed these data in a way consistent with our \\XMM analysis in order to check consistency, and we draw comparisons at several relevant points.  Throughout this paper, a cosmology of $H_0=70\\kmpspMpc$, and $\\Omega_M=0.3$ ($\\Omega_\\Lambda=0.7$) is adopted, unless stated otherwise, and all errors are quoted at the $68\\%$ level. At the cluster's redshift, $1\\arcs$ corresponds to $7.8\\kpc$ in this cosmology. The virial radius ($\\r200$) is defined as the radius within which the mean density is $200$ times the critical density at the redshift of observation. ",
        "conclusions": "ClJ1226.9$+$3332 is a remarkable and unique cluster. We have performed a detailed analysis of an \\XMM observation, and after careful comparison of background subtraction methods, we have confirmed its high temperature, and produced a temperature profile for the first time at this high redshift ($z=0.89$). The total mass is found to be extremely high ($1.4\\pm0.5\\times10^{15}M_\\odot$) and similar to that of the Coma cluster. The probability of such a cluster being found in the discovery survey is $0.64$ (assuming a \\LCDM cosmology).  The relaxed, and generally isothermal, X-ray appearance, together with the gas mass fraction, metal abundance, and gas density profile slope ($\\beta$) all being consistent with those of local clusters, suggests that this cluster was assembled significantly earlier than z=0.9.  The high luminosity and relaxed nature make it an extremely useful subject for further studies of the gas, dark matter and galaxy properties out to large radii at high redshift. Deeper \\Chandra and \\XMM observations are planned, in part to test the assumptions of isothermality and hydrostatic equilibrium which underpin the derivations of many of the cluster properties."
    },
    "0403/astro-ph0403717_arXiv.txt": {
        "abstract": "Circumbinary (CB) disks have been proposed as a mechanism to extract orbital angular momentum from cataclysmic variables (CVs) during their evolution.  As proposed by Taam \\& Spruit, these disks extend outwards to several a.u. and should be detected observationally via their infrared flux or by absorption lines in the ultraviolet spectra of the CV.  We have made use of archival {\\it HST}/STIS spectra as well as our own near-IR imaging to search for observational evidence of such CB disks in seven CVs.  Based on the null result, we place an upper limit on the column density of the disk of $\\nh\\sim10^{17}\\cm^{-2}$. ",
        "introduction": "Cataclysmic variables (CVs) are mass transferring binary star systems consisting of a white dwarf primary star and red dwarf secondary star. The evolutionary path from detached binary star system to close, interacting binary star system is greatly affected by orbital angular momentum loss (OAML), of which there must be significant amounts in order to transform the originally long period progenitor into an interacting binary \\citep[see][]{how01}.  A large amount of the OAML occurs during the common envelope (CE) phase of the CV evolution.  Although the exact mechanism for OAML is still a matter of debate, several theories have been advanced to explain the observed orbital period-mass transfer distribution of CVs.  Generally, magnetic braking is invoked for long period systems ($P_{\\rm orb}>3$ hr), and although these models are marginally successful in explaining the observed distributions, they predict a strong correlation between mass transfer rate ($\\dot{M}$) and orbital period, along with an excessive number of ultra-short period, very low luminosity systems; neither the stringent $\\dot{M}-P_{\\rm orb}$ correlation, nor the preponderance of low luminosity systems are observed.  The discord between theory and observation could be due to a number of different factors, the most intriguing of which is that magnetic braking may not be solely responsible for OAML, but rather angular momentum may be extracted from the binary system via a circumbinary (CB) disk \\citep{spr01,taa01,dub02,taa03}.  The gravitational torque exerted by the CB disk on the binary will cause significant OAML leading to an accelerated period evolution and allowing the secondary star to completely dissolve within a finite time (less than a Hubble time). These two consequences are the attraction of the CB disk theory, as they would cause a broad range of $\\dot{M}$ for any given period and remove the need for a large population of faint CVs.  If CB disks exist, they likely begin formation during the CE phase and acquire material via mass outflow from the white dwarf (novae) or accretion disk (winds and outbursts) \\citep{spr01}. For the CB disk to become a significant source of OAML, the mass input rate into the CB disk must be $10^{-5}$ to 0.01 of the material transferred from the secondary to the white dwarf throughout the CB disk lifetime.  CB disks are assumed to be geometrically thin with a pressure scale height to radius ratio of $H/R\\sim0.03$, have a viscosity $\\alpha\\la0.01$, and of Keplerian structure.  \\citet{dub02} have proposed that the inner edge is truncated due to tidal forces from the binary, has a radius near $1\\times10^{11}\\cm$, effective temperature $T\\sim3000$ K, and is optically thick, while the majority of the outer disk is cool, $T<1000$ K, optically thin, and can extend out to several a.u.  Continuum contribution comes from dust emission, with opacities due to dust grains, and includes contributions from iron-poor opacities at $T\\la1500$ K where dust particles exist.  Based on models of protoplanetary disks \\citep[e.g.,][]{bel97}, \\citet{dub02} find that their CB disks reach a maximum vertical height above the plane of the disk at a radius less than the outer edge of the disk.  However, the exact physical height of the outer edge of the CB disk is not known (R. Taam, private communication). Figure~\\ref{cb1} shows a schematic model of a CB disk based on specifications given in \\citet{dub02}.  Two properties of the proposed CB disks will manifest observationally; the large extent of the CB disk and its surface mass density, $\\Sigma \\sim10^2-10^4\\, {\\rm g}\\cm^{-2}$, should make the disk a source of ''ISM-like'' absorption lines appearing at the systemic velocity of the system, and, if the CB disk emission (from the warm inner edge) can be considered to be a sum of blackbodies, the flux from the CB disk will peak near $3\\,\\micron$ \\citep{dub02}.  These properties lend themselves to detection in two different wavelength regions.  Flux measurements via near-IR ($1-10\\,\\micron$) imaging of face-on systems (offering a full view of the CB disk) may detect the IR flux of the CB disk.  A second detection technique is to search for the ISM-like absorption lines produced by the CB disk in UV spectra of the binary systems.  {\\it Hubble Space Telescope}/ Space Telescope Imaging Spectrograph ({\\it HST}/STIS) data are ideal for this as the high spectral resolution will enable detection and measurement of ISM-like lines whose radial velocities may then be compared to the space motion of the binary.  This technique is best done for CVs with inclinations that match the opening angle of the CB disk, as the binary is then observed through the plane of the CB disk.  The theoretical work \\citep[e.g.,][]{taa01} predicts that the higher mass transfer rate CVs are the best candidates for a CB disk, as material may continuously be entering the CB disk, and may still retain some vertical extent above the disk (i.e., the CB disk has not had time to ``flatten out''). Thus, we have begun an observational search for CB disks with long period nova-likes and old novae. ",
        "conclusions": "\\label{disc} For each spectrum, we identified absorption lines of ionic species commonly associated with the ISM and CV systems, such as C, N, Mg, Si, and S.  We considered absorption lines that are above the $2\\sigma$ error level in flux and that have full width at half maximum (FWHM) values $<0.2$ \\AA, as absorption lines with larger FWHM values would not be associated with the slowly rotating CB disk or the ISM. Absorption lines were fit with a single Gaussian function to determine line centers and equivalent widths (EWs). Some absorption lines are composed of well-separated multiple absorption components (for example, the absorption lines of QU~Car); for these we measured the line centers of the components individually.  At the resolution of the E140M grating, our spectra are not sufficient to resolve closely spaced ISM features; for absorption lines containing unresolvable components, we measured the lines with a single Gaussian.  Radial velocities for all measured absorption components were then calculated and are given in Table~\\ref{all_tab}, along with the EWs of the absorption lines.  The measurement errors on the radial velocities for narrow lines that are unblended and unsaturated are within the velocity resolution of the spectrum, which is typically $\\sim3\\kms$. Lines with EW $\\,>100$ m\\AA\\ are likely saturated or close blends of two strong components.  In such circumstances, it is difficult to assign a meaningful centroid velocity to these features; we therefore assign uncertainties of $\\pm7\\kms$ to the radial velocities calculated for these lines.  Absorption lines of excited state ions are noted in the table by an asterisk.  \\subsection{ISM Lines of Sight} Many of the identified absorption lines will be due to absorption by the ISM, in particular the local interstellar cloud (LIC); the warm, partially ionized cloud in which the Sun resides.  The heliocentric velocity of the LIC is $25.7\\kms$ towards the Galactic coordinates $l=186\\fdg1$, $b=-16\\fdg4$ \\citep[e.g.,][]{woo02}.  Evidence of the LIC is frequently observed over a large fraction of the sky as a distinct range of velocities associated with ISM absorption components present in high dispersion UV spectra of nearby stars \\citep[e.g.,][]{lal92,hol98,woo02,kim03}.  In addition to the LIC, many other velocity components are also observed in certain directions.  The LIC and these other components are commonly discussed in terms of an assemblage of warm interstellar clouds within 30 pc of the sun.  A good up-to-date and lucid description of these ISM velocity components is provided in \\citet[hereafter FGW]{fri02}.  In Table~\\ref{objects}, we give the velocity of the LIC in the direction of each of our candidate objects.  The LIC velocity is easily recognized in the histograms we have created for the ISM-like absorption lines found in the spectra of the six CV systems we studied.  These histograms are shown in Figures~\\ref{dilac_hist} - \\ref{v3885_hist}.  For four of the six objects, DI~Lac, IX~Vel, TZ~Per, and QU~Car, the LIC velocity is at or near the peak velocity in the histograms.  EX~Hya and V3885~Sgr do not follow suit; the peak in the EX~Hya histogram is $\\sim7\\kms$ from the LIC velocity, and neither of the peaks in the V3885~Sgr histogram are aligned with the LIC velocity.  The line of sight to DI Lac lies in a direction expected to show LIC components.  However, the projected LIC velocity very nearly coincides with the systemic velocity of the system, making it difficult to observationally distinguish either (see Figure~\\ref{dilac_hist}). Complicating matters is the presence of lines due to interstellar \\ion{C}{1}, an ion not present in the LIC because it is easily destroyed by the ambient stellar UV radiation field.  The presence of \\ion{C}{1} lines indicates that the line of sight to DI Lac passes through a diffuse cloud with sufficient density to preserve the \\ion{C}{1} ion.  The presence of an unusually dense component of the ISM is also indicated by the very strong and saturated lines of \\ion{S}{2}.  The mean velocity of the \\ion{C}{1} lines and the \\ion{S}{2} lines ($-1.78\\kms$) contributes significantly to the peak in the velocity distribution adjacent to the predicted LIC velocity. It would be tempting to associate these lines with a CB disk, however, the velocities are not consistent and the physical conditions in the CB disk are not appropriate to the existence of the \\ion{C}{1} ion. Rather, it appears that the line of sight to DI Lac passes through a relatively dense interstellar cloud, unrelated to the CV system. There is no evidence in the spectrum of DI Lac of lines arising from a CB disk.  The line of sight to EX Hya lies in a direction that shows no LIC components.  The peak of the radial velocity histogram (Figure \\ref{exhya_hist}) does, however, coincide with the expected velocity of the Filament structure described in FGW, consistent with several stars that lie in the same general direction as EX Hya.  There is no evidence for a CB disk and the line of sight to this star is consistent with the ISM only.  For IX Vel and QU Car, the line of sight to each object lies in a direction expected to show LIC components and the peaks of the histograms coincide with the expected LIC velocity (see Figures \\ref{ixvel_hist} and \\ref{qucar_hist}).  The line of sight to IX Vel is consistent with the ISM only.  The spectrum of QU Car shows several relatively weak \\ion{C}{1} lines, indicating the line of sight passes through a relatively dense interstellar cloud.  There is no indication of CB disk related absorption for QU Car.  The line of sight to TZ Per lies in a direction expected to show LIC components, however, the histogram of velocities (see Figure \\ref{tzper_hist}) shows a broad distribution of velocities.  Like the spectrum of DI Lac, the spectrum of TZ Per contains \\ion{C}{1} lines as well as relatively strong \\ion{S}{2} lines, consequently a similar interpretation - a line of sight through a dense interstellar cloud - holds for TZ Per.  Many of the features in the TZ Per spectrum are either saturated or quite strong, indicating that several unresolved velocity components are present and contributing to the broad distribution of velocities.  There is one \\ion{C}{1} line appearing near the systemic velocity, however, this is not sufficient evidence to support absorption from a CB disk and as discussed above, the existence of this ion in a CB disk is not probable.   The line of sight to V3885 Sgr contains two distinct, well separated velocity components (see Figure \\ref{v3885_hist}).  The component at $+2.5\\kms$ can be confidently identified with the South Polar Cloud (see FGW).  The second peak near $20\\kms$ corresponds to no previously described interstellar cloud in the local ISM.  There is no clear contribution from the LIC or a CB disk.  The line of sight to this star is consistent with the ISM only.  \\subsection{Absorption Line Velocities} Figures~\\ref{di_vel} - \\ref{v3_vel} display line velocity profiles for select absorption lines from the spectra of each of the six CVs.  On each plot, the vertical dashed line represents the LIC velocity towards the object, and the vertical dotted line represents the systemic velocity of the system.  For systems that have systemic velocities with large errors, the dotted line represents the measured value.  We selected a sampling of absorption lines to illustrate the line structure and velocity behavior for each system.  Shown are the common absorption lines of the \\ion{N}{1} $\\lambda\\lambda 1199.55,1200.22,1200.71$ triplet, lines from several \\ion{C}{1} multiplets, the \\ion{S}{2} $\\lambda\\lambda 1250.58,1253.81,1259.52$ triplet, lines from several \\ion{Si}{2} multiplets, along with the \\ion{Fe}{2} $\\lambda1608.46$ and \\ion{Al}{2} $\\lambda 1670.79$ absorption lines (excited transitions are noted with an asterisk).  Figure~\\ref{di_vel} displays absorption lines found in the spectrum of DI~Lac.  Each line exhibits a single velocity structure, however, because the LIC velocity and systemic velocity are separated by less than the velocity resolution of the spectrum, it is not possible to identify the absorption lines with either velocity.  As discussed in the previous section and based on the results from the remaining objects of our study (see below) it is likely that the absorption lines in the DI~Lac spectrum are ISM in nature.  The 1910 nova eruption of DI~Lac implies the presence of nova shells, which may be observed via absorption lines, however, these lines should be blue-shifted with respect to the systemic velocity.  The absorption lines found in the EX~Hya spectrum (shown in Figure~\\ref{ex_vel}) have a single velocity structure.  They are shifted $\\sim7-10\\kms$ with respect to the LIC velocity of $-14.8\\kms$, but are aligned with the Filament structure discussed in FGW.  The line velocities of IX~Vel (Figure~\\ref{ix_vel}) and QU~Car (Figure~\\ref{qu_vel}) are aligned with the LIC velocity.  In each case, the systemic velocity of the system is clearly not associated with the absorption line velocities.  TZ~Per (Figure~\\ref{tz_vel}) is another system for which the LIC and systemic velocities fall within the velocity width of the absorption line profiles.  The two velocities are separated by more than the velocity resolution of the spectrum, therefore the absorption line velocities may be identified with the LIC velocity, as it is more closely aligned with the absorption line centers (the LIC velocity alignment is also seen in Figure~\\ref{tzper_hist}).  The absorption lines in the V3885~Sgr spectrum (Figure~\\ref{v3_vel}) have two velocity components; one that is aligned with the South Polar Cloud (FGW) and one that does not match any known ISM velocity system. Sagittarius is in the direction of the galactic center, so it is likely that different clouds along this line of sight are causing absorption along the line of sight to V3885~Sgr.   \\subsection{Infrared Observations} Although the outer regions of the proposed CB disks are cool and optically thin and not a source of IR blackbody flux, the warm, optically thick inner regions of the disk will make a contribution to the IR flux of a CV system.  Given that the inner region of the CB disk is predicted to have a temperature of up to $\\sim3000$ K and is optically thick, we approximated a mean temperature of the CB disk that agrees well with Figure 2 of \\citet{dub02} to create a blackbody model representative of the CB disk.  Determining this temperature is important, as the peak of the blackbody curve of the CB disk should cause an observable 'bump' in the spectral energy distribution (SED) of the CV.  V592~Cas has an inclination of $i=27\\degr$, and is therefore ideal for detecting the full CB disk via IR observations.  In Figure~\\ref{fred}, we plot $JHKL'$ points for V592~Cas, along with blackbody curves representing the secondary star for the system and a steady state accretion disk.  Also plotted is the sum of the IR flux from all components of the CV system, including the flux contribution expected from a CB disk; it is a sum of blackbodies and peaks near $3\\,\\micron$.  This model summed spectrum overestimates the IR flux observed for V592~Cas; the data points do not fall along the model curve that includes the flux contribution of the CB disk, but rather follow the blackbody curve of the steady state accretion disk.  It is apparent that flux from a CB disk is not evident in the IR data of V592~Cas.   \\subsection{Observations of the ISM} Most of the spectral features observed in the six CV systems studied can be reasonably assigned to the LIC or other recognized velocity structures in the local ISM near the Sun.  However, in several CVs (DI Lac, QU Car, and TZ Per), we do observe unusually strong lines due to \\ion{S}{2} and features due to \\ion{C}{1}, indicating that the line of sight to these stars passes through a dense, warm ISM component where \\ion{C}{1} is shielded from ionizing radiation.  It is regarded as unlikely that these regions are associated with either CB disks or other gas related to the CV, as the observed velocities of these components are distinctly different from the CV system velocity.     \\subsection{Evidence for CB Disks?} In the spectroscopic and photometric data presented here, we have found no convincing evidence to support the theory of CB disks around CVs.  The ISM-like lines appearing in the spectra of the six CVs have velocities consistent with the ISM in the direction of the object. The two cases in which the systemic velocity fell within the velocity width of the measured absorption lines (DI~Lac and TZ~Per) were complicated by the fact that the systemic velocity was close to the LIC velocity and that the identified lines were unlikely to occur in a CB disk.  Based on the results of the other four objects, the match of the systemic velocity is due purely to its proximity to the LIC velocity.  Evidence for a CB disk was not found in the $JHKL'$ photometric data of V592~Cas.  There is no excess of flux around the expected peak wavelength of $3~\\micron$, and models of a CB disk overestimate the flux measured in V592~Cas.  In the absence of any absorption features clearly associated with CB disks, the upper limits for the disk material may be derived based on 3$\\sigma$ upper limits for the EWs of \\ion{N}{1}, \\ion{Si}{2}, \\ion{O}{1}, and \\ion{S}{2}, evaluated at the systemic velocity in the spectrum of each CV (except for DI Lac, which exhibits some lines near the systemic velocity).  The 3$\\sigma$ EW values for all lines measured are very low - they range between 5 m\\AA\\ - 25 m\\AA\\ - so we can use the linear portion of the curve of growth to determine the species column densities: $N_i = 1.129\\times10^{20}(\\frac{W_i}{\\lambda_i}) (\\frac{1}{f_i \\lambda_i})$. Wavelengths and $f$ values are taken from \\citet{mor91}.  To calculate equivalent H column densities, we use typical abundances of each species observed for the warm phase of the ISM, $A_i$, such that $\\nh = \\frac{N_i}{A_i}$.  The best upper limits we calculate are from the \\ion{N}{1} and \\ion{O}{1} lines.  For EX Hya we determine a value of $\\nh \\sim 3\\times10^{17}\\cm^{-2}$, for IX Vel $N_{\\rm H} \\sim 1\\times10^{17}\\cm^{-2}$, for QU Car $\\nh \\sim 2\\times10^{17}\\cm^{-2}$, for TZ Per $\\nh \\sim 2\\times10^{17}\\cm^{-2}$, and for V3885 Sgr $\\nh \\sim 2\\times10^{17}\\cm^{-2}$.  We have indicated the observational upper limit of $\\nh \\sim 10^{17}\\cm^{-2}$ as the dashed horizontal line in Figure \\ref{diskmodel} for comparison of the observational numerical values with the theoretical disk models.  The vertical dotted line at $z/H\\sim7$ represents the value of $z/H$ for EX Hya, while the vertical dotted line at $z/H\\sim17$ represents the value of $z/H$ for IX~Vel, QU~Car, and V3885~Sgr (the values for DI~Lac and V592~Cas are past the limits of the plot).  If CB disks exist, we can eliminate Models C and D as the proposed CB disk structure for the four CVs, as the model column densities are significantly above the detection limit and absorption by such disk structures would be detectable at the given inclinations.  A CB disk with the structure of Model A would be detectable for the inclination of EX Hya and can therefore be ruled out, as no absorption from a CB disk is seen in EX Hya.  Model B would not be observed in any of our CV systems, but we note that this model is the most conservative disk assumption of the four and we do not expect it to be physically realistic (i.e., we would expect the CB disk to have some form of an atmosphere).  As this investigation has only considered a small sample of CV systems that could contain a CB disk, we plan to continue our photometric and spectroscopic investigation for CB disks.  In particular, we plan to analyze additional archival STIS data and to continue to look for their signatures in the IR SEDs of candidate CVs."
    },
    "0403/astro-ph0403467_arXiv.txt": {
        "abstract": "As part of a detailed study of the gas content in a sample of Early-Type galaxies, we present 21\\,cm \\ion{H}{1} line maps of the S0 galaxy NGC\\,404, at a spatial resolution of $15.2''\\times14.4'' \\, (\\alpha\\times \\delta)$ and a velocity resolution of $\\rm 2.6\\, km\\,s^{-1}$. The \\ion{H}{1} has been traced out to a radius R $\\sim 8R_{25}$ or 48 disc scale-lengths, making it one of the largest \\ion{H}{1} extents reported ($800''$ or 12.6\\,kpc at the assumed distance of 3.3\\,Mpc). Approximately 75\\% of the \\ion{H}{1} resides in a doughnut which is seen close to face--on with inner and outer radii of $\\sim R_{25}$ and $\\sim 4R_{25}$. The optical galaxy fits snugly within the hole of the doughnut. The remaining 25\\% of the neutral gas is found in an annulus concentric with the doughnut and with a somewhat larger ellipticity, extending from a radius of $\\sim 4R_{25}$ to $\\sim 8R_{25}$.  A total \\ion{H}{1} mass of $1.52\\pm 0.04\\times10^8\\, M_\\odot$ is found, which implies an $\\mhi/L_{B} = 0.22$ in solar units. We argue that most if not all of this gas is of external origin, most likely due to the merger of a dwarf irregular galaxy with $M_B \\sim -15.5$\\,mag.  The velocity field shows a steeply declining observed rotation curve, compatible with Keplerian decline. However, because the galaxy is close to face--on there is a degeneracy in the determination of the intrinsic rotation curve and inclination. We therefore analysed two extreme cases, producing tilted--ring model fits forcing either a Keplerian or a Flat rotation curve through the observations; both approaches result in equally plausible fits. In both model fits, the position angle of the kinematical major axis of the annulus is distinct from that of the doughnut and ranges from $160^\\circ$ to $120^\\circ$ (for the doughnut these values are $100^\\circ$ to $60^\\circ$).  Assuming a distance of 3.3\\,Mpc, a total mass of $3\\times10^{10}\\, M_\\odot$ is found on the basis of the Keplerian rotation curve. This implies a global $M_{\\rm T}/L_{B}$ ratio of $\\sim 44$ in solar units which is high and likely a reflection of the low blue luminosity of the galaxy ($\\sim 15$ times lower than the average S0 luminosity). Values for a Flat rotation curve are a factor of 4 higher. ",
        "introduction": "Traditionally, S0 galaxies were thought to represent a transition between elliptical (E) and spiral (S) galaxies. They share with the latter a central bulge and an exponential disc, and with the former the lack of spiral structure or arms. They are considered inert objects with very little gas, composed almost exclusively of Population II type stars, and virtually a total absence of on--going star formation. This picture has changed considerably over the past 20 years thanks mainly to tremendous progress in observational techniques which are now routinely delivering high sensitivity data over a wide range of wavelengths.  The interstellar medium of early--type galaxies, and especially that of S0s, has attracted a great deal of interest ever since \\citet{fg76} showed that a considerable amount of mass ejected by evolved stars could pile up over the lifetime of the galaxy. Surprisingly, though, high resolution, sensitive \\ion{H}{1} observations demonstrated that in general hardly any gas is detected and that most of it seems to have been accreted (\\citealt{vgea86,wh90}). The scarcity of material detected in \\ion{H}{1}, H$_\\alpha$, CO, and at FIR wavelengths and the origin of that material form two puzzles of contemporary astronomy.  \\citet{fg76} discuss several ways in which gas could have been expelled or eliminated from early--type systems. They reach the conclusion that gas is particularly hard to remove from these systems via star formation. Although star formation has virtually ceased, there are indications that some star formation is taking place in S0s, as based on IRAS observations (\\citealt{j86,kgkj89}) and CO observations (\\citealt{wh89,s90}) which show that some S0s rival spirals in their CO content, but that the molecular gas is concentrated within 1--2 kpc from the centre.  We have started a project, mapping a few early--type galaxies in detail, in order to shed light on the origin of the ISM in these objects. In this paper we present the results obtained for one of the galaxies under study, NGC\\,404.  NGC\\,404 is a typical example of an early--type galaxy which shows more activity than expected from such systems. It contains an appreciable amount of atomic \\citep{bw76} and molecular gas \\citep{wh90}, a prominent, semicircular dust lane encircling the nucleus \\citep*{bcr82} and various bright UV point sources \\citep{mea98}. It was classified a LINER by \\citet{sea90}. \\citet{m58} classified this galaxy as E$_{\\rm pec}$, because of the dust lane.  \\citet*{hms56} and \\citet{s61} classified it as an S0$_3$. Based on photographic surface photometry, \\citet{bcr82} show that the object consists of bulge, lens, and exponential disc components, characteristic of a lenticular galaxy. This easily rules out a classification as (dwarf) elliptical, as several authors have proposed.  For many years the distance to NGC~404 was basically an unknown quantity. Estimates ranged from 1.5\\,Mpc \\citep{bw76} to 10\\,Mpc \\citep{wh90}. It is only recently that a firm distance to this object has been established. \\citet{tdb} used the surface brightness fluctuation method to come up with a distance of 3.3 Mpc. This was confirmed independently by \\citet{kara02} who used the tip of the red giant branch and determined a distance of $3.15 \\pm 0.32$\\,Mpc.  Although NGC~404 is listed as a member of the loose group LGG~011 \\citep{g93}, it is unlikely to be physically associated with the rest of the group as at 3.3~Mpc \\citep{tdb} the nearest member, NGC~598 (M\\,33), is located at 2.6 Mpc. Rather, NGC~404 is likely the nearest isolated S0 galaxy, especially as there are no other known galaxies within 1.1~Mpc of it \\citep{kara02}. It is interesting to note that the field is only two degrees away from the centre of the Piscis cluster of galaxies ($\\rm \\alpha = 01^h\\,10^m \\;\\delta = 33.4^\\circ$). A summary of optical data on NGC~404 is given in Table~\\ref{optical-datos}. ",
        "conclusions": "NGC~404 is an exceptionally gas rich early--type galaxy.  The \\ion{H}{1} is distributed in two separate systems, a doughnut which extends from $R_{25}$ or $100''$ to $4R_{25}$ (1.6 to 6.4 kpc for a distance of 3.3\\,Mpc) and an annulus of much lower surface brightness which extends from $4R_{25}$ to roughly $8R_{25}$.  We derive a total \\ion{H}{1} mass of $1.52\\pm 0.04\\times10^8\\,M_\\odot$ of which 75\\% resides in the doughnut and the rest in the outer annulus. This translates to a distance independent $\\mhi/L_{B}$ ratio of 0.22 in solar units which is high for an S0 galaxy knowing that the median for S0's galaxies is $0.04 \\,M_\\odot/L_\\odot$ whereas for Sb galaxies this is $\\sim 0.22\\,M_\\odot/L_\\odot$ \\citep{rh94}.  The velocity field shows a steeply declining rotation curve. Unfortunately, the galaxy is oriented almost face--on and, assuming circular rotation of the gas in annuli, good fits can be obtained when forcing both a flat rotation curve {\\em and} a rotation curve that falls off like a Keplerian. In both cases the position angle (PA) of the kinematical major axis varies with radius by $\\sim 40^\\circ$.  In the Keplerian decline model, the gas in the outer annulus is warped, the annulus moving to a less face--on orientation with radius.  A total mass of $3\\times10^{10}\\, M_\\odot$ is found when using the Keplerian rotation curve. This implies a global $M_{\\rm T}/L_{B}$ ratio of $\\sim 44$ in solar units, which is exceptionally high. This is partly due to the fact that the blue luminosity of the galaxy is $\\sim 15$ times lower than average for an S0. For a flat rotation curve, a lower mass limit of $12\\times10^{10}\\, M_\\odot$ or four times higher is derived, leading to a consequently four times higher total mass--luminosity ratio of $\\sim 175\\, M_\\odot/L_\\odot$.  The nuclear source is detected at 20\\,cm. If nonthermal, its luminosity of $\\rm \\log P_{21} (W\\,Hz^{-1})=17.9$ is a factor of ten lower than Cas~A and more in line with that of the Crab nebula. If dominated by star formation, the SFR is a low SFR($ M\\ge 5\\,M_\\odot)\\,\\approx 2 \\times 10^{-4} M_\\odot\\,\\rm yr^{-1}$, 2--6 O stars being enough to explain the observed luminosity.  Given the mass and luminosity of NGC~404, we are dealing with a dwarf lenticular galaxy. Assuming Keplerian decline, the DM halo is compact and ends near $2R_{25}$. Our preferred explanation for the origin of the ISM in this system is that it has been accreted through a catastrophic merger with a gas--rich, dwarf irregular system, less than a tenth of the mass of NGC~404 which has taken place anywhere between $0.5-1.0\\times10^9$ yr ago (no other objects having been encountered within a volume of 1.1\\,Mpc radius). Gas at small radius has already settled within the plane whereas gas in the outer annulus is still gravitating towards the midplane of the disc defined by the S0, resulting in a warp of the gas distribution.  Although none of our conclusions depends strongly on the rotation curve being flat or declining, it is of course extremely frustrating that it hasn't proven possible to decide between these two extreme models, especially since the observed radial velocities decline so strongly. This leaves NGC~404 as {\\em potentially} the first galaxy with a rotation curve that shows such a pronounced decline over such a large extent which, if confirmed, would be of extreme importance."
    },
    "0403/astro-ph0403698_arXiv.txt": {
        "abstract": "On large scales galaxies and their halos are usually assumed to trace the dark matter with a constant bias and dark matter is assumed to trace the linear density field. We test these assumption using several large N-body simulations with $384^3-1024^3$ particles and box sizes of $96-1152h^{-1}$Mpc, which can both resolve the small galactic size halos and sample the large scale fluctuations. We explore the average halo bias relation as a function of halo mass and show that existing fitting formulae overestimate the halo bias by up to 20\\% in the regime just below the nonlinear mass. We propose a new expression that fits our simulations well. We find that the halo bias is nearly constant, $b \\sim 0.65-0.7$, for masses below one tenth of the nonlinear mass. We explore next the relation between the initial and final dark matter in individual Fourier modes and show that there are significant fluctuations in their ratio, ranging from 10\\% rms at $k\\sim 0.03$h/Mpc to $50\\%$ rms at $k\\sim 0.1$h/Mpc. We argue that these large fluctuations are caused by perturbative effects beyond the linear theory, which are dominated by long wavelength modes with large random fluctuations. Similar or larger fluctuations exist between halos and dark matter and between halos of different mass. These fluctuations must be included in attempts to determine the relative bias of two populations from their maps, which would otherwise be immune to sampling variance. ",
        "introduction": "Determination of the power spectrum of mass fluctuations and its redshift evolution is one of the main goals of modern observational cosmology. Its accurate measurement would allow us to test some of the most fundamental questions in cosmology today, such as the shape of primordial power spectrum and its relation to fundamental theories of structure formation, the mass of neutrino and the nature of dark energy.  In general there are two approaches to the measurement of the matter power spectrum. One is to measure galaxies, either in redshift space or in angular position (perhaps supplemented by photometric redshift information) and to assume they trace the dark matter. This assumption is believed to be valid on large scales, where the so called linear bias model assumes that the galaxy density field is proportional to the matter density field times a free parameter called bias. While power spectrum measurements of galaxies with modern surveys such as SDSS \\citep{2003astro.ph.10725T} or 2dF\\citep{2002MNRAS.337.1068P} have enormous statistical power, they can only determine the shape of the matter power spectrum and not its amplitude because of the bias uncertainties. This limits their use in the study of the growth factor evolution, important for investigations of dark energy models. In addition, on small scales information from galaxy clustering is limited by the uncertainties in the relation between the galaxies and the dark matter, which make the bias scale dependent. For this reason the small scale information is usually discarded.  The situation with galaxies would not be as dimm if we could determine the bias. Here we explore a method to determine galaxy bias based on determination of clustering amplitude of faint galaxies. These are likely to occupy low mass halos which, as we show in this paper, have a well determined large scale bias that is nearly independent of halo mass. While some fraction of these galaxies are satellites in larger halos, this can be quantified and corrected for (see \\cite{2004astro.ph..6594S} for a first application of this method to the real data). In addition, there exist populations, such as IRAS galaxies, for which this fraction may be small. In these cases measuring the large scale power spectrum amplitude for these galaxies determines the matter power spectrum amplitude as well.  The main problem with using faint galaxies as tracers of large scale structure is that in a typical flux limited survey faint galaxies occupy a small nearby volume, so the sampling variance errors for power spectrum on large scales are large. However, we can still determine their bias relative to a population of brighter galaxies. If there is no stochasticity between the two populations then a direct comparison of the maps gives an accurate determination of the relative bias with no sampling variance. We can then use the power spectrum determination of the brighter population, with smaller sampling variance errors because of larger volume covered, to determine the power spectrum of the fainter population and of the matter itself. The limiting source of noise is the stochasticity between these fields, which we explore in this paper.  Another approach to determine matter fluctuations is to use weak lensing induced correlations between background galaxy ellipticities. These are sensitive to the dark matter fluctuations directly and as such this approach holds the promise to improve upon the limitations of the galaxy clustering methods. Its main limitation is that it traces the dark matter in angular projection and has large sampling variance errors on large scales. This limits the statistical power of the weak lensing surveys. On small scales the nonlinear corrections, noise, intrinsic correlations and other systematic contaminations become significant, all of which may complicate the modelling.  A possible approach to achieve the best of both worlds is to combine the weak lensing and galaxy clustering surveys: one can use the weak lensing to determine the galaxy bias and then use the 3-d galaxy clustering information to improve on the statistical errors. One way to do this is to use galaxy-dark matter cross-correlation analysis from weak lensing and combine it with the galaxy auto-correlation analysis. If galaxies are tracing perfectly the dark matter then it suffices to have a few well measured modes in both fields to determine the galaxy bias. This has been proposed as a way to get around the sampling variance in weak lensing surveys \\citep{2004astro.ph..2008P}. In the absence of stochasticity it gives the bias (and so the dark matter power spectrum) without the usual sampling variance errors, assuming the analysis is done on the same patch of the sky and with the correct radial weighting of the galaxies to match that of the dark matter.  In both of these cases the underlying assumption is that there is no stochasticity between these fields on large scales. While there have been analytic attempts to address this assumption \\citep{1999ApJ...525..543M}, it has not been tested well with simulations in the past due to the lack of sufficient dynamic range (but see \\cite{2002MNRAS.333..730C} for a related study). One must resolve the halos small enough to be suitable as galaxy hosts (with typical masses at or below $10^{12}h^{-1}M_{\\sun}$). At the same time, the simulations must be large enough so that many long wavelength modes are sampled to determine the statistics of interest. We achieve this by using a set of new simulations with a larger dynamical range. The number of particles in these simulations, $10^8-10^{9}$, and their box size, 100-1000$h^{-1}$Mpc, allow a much better exploration of the halo bias and stochasticity on scales larger than available before.  In addition to exploring the relation between halos and matter we can also investigate the relation between the initial and final matter distribution. Weak lensing measures the nonlinear matter field, while for the study of linear growth factor one would like to know the relation between galaxies and linear matter field instead. We explore the relation between the final and initial dark matter density field on large scales, where this relation is believed to be perfect. This case is amenable to perturbation theory analysis and as such allows one to interpret and verify the numerical simulation results.  Finally, we revisit the question of halo bias as a function of halo mass with the new simulations. This has been addressed by previous generation of simulations using $256^3$ particles and box sizes of order (100-140)$h^{-1}$Mpc \\citep{1998ApJ...503L...9J,1999MNRAS.308..119S,2001MNRAS.323....1S}. This, as noted by the authors themselves, is barely adequate for this purpose because of large shot noise at high halo masses and insufficient number of large scale modes where linear evolution is valid. The goal of these papers was to provide expressions which fit over a range of power law simulations and were not specifically optimized for realistic $\\Lambda $CDM models. They provide expressions that fit the simulations available at the time to a reasonable accuracy, but which can be systematically wrong by as much as 10-20\\%. To put things into a current context, the statistical error on the amplitude of galaxy clustering from SDSS using $k<0.2h$/Mpc modes is 1\\% \\citep{2003astro.ph.10725T}, so a perfect bias determination, for example using the faint galaxies as described above, would allow us to reach this accuracy on the matter power spectrum. In this era 20\\% accuracy no longer suffices and the goal of the present paper is to provide more accurate expressions for halo bias as a function of mass and cosmological model. ",
        "conclusions": "In this paper we have addressed the relations between the matter density field, halos and initial density field, focusing on large scales where these are often assumed to be proportional to each other. We focus on two issues. First, what is the scatter between these fields around the average relation? This is expressed here in terms of relative scatter between the mode amplitudes, which is related to the stochasticity parameter $r$, defined as the cross-correlation coefficient between the two fields. While the two are related we emphasize that even small deviations of $r$ from unity may lead to large relative fluctuations between the two fields. These are of interest whenever one is trying to relate the fields to one another to determine their relative amplitudes. One example is the bias determination using the cross-correlation between the weak lensing signal (tracing the matter density) and the galaxies. Another example is the relative bias determination between two different galaxy populations, which we propose here as an alternative method to determine the galaxy bias, because galaxies in low mass halos have a bias of $b\\sim 0.7$ independent of their mass. In all cases we find the scatter between the fields in individual modes is significant and one cannot assume the fields are simply proportional one to another. This scatter, coupled with a small number of modes on large scales, makes it difficult to accurately determine the bias (or relative bias) and needs to be included in the predictions of how accurately can one determine the matter power spectrum with these methods.  The second goal of this paper was to revisit the halo bias as a function of halo mass. This relation is a fundamental ingredient of any halo model \\citep[see][ for a recent review]{2002PhR...372....1C} and plays an important role if one is trying to model galaxy clustering by connecting it to the underlying halos. The previous generation of simulations \\citep{1998ApJ...503L...9J,1999MNRAS.308..119S} had a limited dynamical range and the predictions were not tuned specifically for $\\Lambda$CDM models. As a result the existing expressions overestimate the bias by as much as 20\\% in the range below the nonlinear mass, which is likely to be the mass range for halos that host most of the galaxies. We propose a new expression that fits the simulations better. We argue that this expression should be fairly accurate for other cosmological models of interest as well, as long as the mass is expressed in units of nonlinear mass. We give corrections for small deviations from this model. The overall accuracy on bias-halo mass relation is at the level of 0.03 or better (for $b<1$), which should help with the bias determination from the current generation of observations.  We thank  P. McDonald, N. Padmanabhan and L. Teodoro for help and useful comments, S. Habib for encouraging this collaboration and Zheng Zheng and David Weinberg for comments on the manuscript. US is supported by Packard Foundation, Sloan Foundation, NASA NAG5-1993 and NSF CAREER-0132953. Simulations by MSW were performed on the Space Simulator Beowulf Cluster and ASCI Q at Los Alamos.  Work by MSW was performed under the auspices of the U.S. Dept. of Energy, and supported by its contract \\#W-7405-ENG-36 to Los Alamos National Laboratory."
    },
    "0403/astro-ph0403651_arXiv.txt": {
        "abstract": "s{ We  present  preliminary results   of a   systematic analysis  of  the XMM-Newton spectra of nearby optically  bright QSOs. The objects  have been selected  from the Bright Quasar Survey  sample. The goal of this project is to characterise the X-ray  spectral properties of optically selected QSOs in  the  0.3--12 keV  energy  band.  In most cases,  two component continua are    sufficient  to model  reasonably  well   the observed spectra.  All  but  one  detected  Fe K$\\alpha$  line have narrow profiles. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/hep-ph0403291_arXiv.txt": {
        "abstract": "We show that spontaneous breaking of global $B-L$ symmetry responsible for small neutrino masses via the seesaw mechanism provides a unified picture of hybrid inflation and  dark matter if the scale of $B-L$ breaking is close to the GUT scale. The majoron which acquires a small mass in the milli-eV range due to Planck scale breaking of $B-L$ is the dark matter candidate. The coupling of the majoron field to the neutrino induces a small violation of CPT and Lorentz invariance at the present epoch. We discuss some of the phenomenological and cosmological implications of the model. ",
        "introduction": "While a non-zero neutrino mass is at the moment the only experimental evidence for physics beyond the standard model, as far as particle physics is concerned, in the domain of cosmology, several other oquestions remain unanswered within the standard model framework and also cry out for new physics. They are: ($i$)  particle physics candidates for dark matter, ($ii$) a deeper understanding of the mechanism of inflation \\cite{Lyth:1998xn}, ($iii$) the nature of matter anti-matter asymmetry responsible for the observable universe\\cite{Riotto:1999yt} as well as ($iv$) an explanation of dark energy. There are many interesting and compelling models of new physics that can explain the different cosmological phenomena listed above individually. However, it is always much more desirable that a single model explain more than one phenomenon (ideally, of course, all of them). In this paper we discuss a model\\cite{zhang,cmp} originally designed to explain neutrino masses via the seesaw mechanism that seems to provide, after supersymmetrization, an explanation of both inflation and dark matter in a rather novel manner. The model has other desirable features such as gauge coupling unification, stable dark matter, stable proton due to R-parity conservation, etc.  We use  the supersymmetric singlet majoron \\cite{cmp} model described in Ref.\\cite{zhang} where the seesaw mechanism\\cite{seesaw} for small neutrino masses is implemented by breaking a global $B-L$ symmetry. The idea is to extend the standard model by the addition of one right handed neutrino per family and three standard model singlet superfields $S, \\Delta$ and $\\overline{\\Delta}$. Of the new Higgs fields, $S$ has $B-L=0$, $\\Delta$ has $B-L =-2$ and $\\overline{\\Delta}$ has $B-L=+2$. The theoretical rationale for $B-L$ symmetry is rooted in the present neutrino oscillation results which require that the seesaw scale be much lower than the Planck scale. It is therefore natural to think that it is protected by some symmetry. The simplest symmetry that does this is the $B-L$ symmetry. $B-L$ symmetry could be a global or a local symmetry. If we take it to be a global symmetry, its spontaneous breaking  leads to a Nambu-Goldstone boson, the majoron. Since it is natural to expect that all global symmetries are broken by nonperturbative gravitational (or stringy) effects, we will parameterize these symmetry breaking effects by nonrenormalizable terms in the effective Lagrangian which are suppressed by the Planck scale $M_{P\\ell}$. These Planck scale $B-L$ breaking effects then give a tiny mass to the majoron\\cite{akhm}.  In this paper we show that this model leads to: ($i$) a picture of hybrid inflation of the universe that links the neutrino mass (more precisly the seesaw or $B-L$ scale of order $10^{15}$ GeV) to inflation with all the desired features and ($ii$) Planck scale $B-L$ breaking effects that lead to a mass for the majoron in the milli-electron volt range and make it a candidate for dark matter of the Universe in the same way as the familiar axion, even though the parameters of the model are very different.  It is, of course, interesting that this model ties the neutrino mass and the scale of inflation to the scale of gauge coupling unification.  An interesting implication of the model is that it leads to a cosmologically induced Lorentz violation for neutrinos. We also find that the superpartners of the majoron must be in the MeV range in order to be compatible with the successes of big bang nucleosynthesis (BBN). ",
        "conclusions": ""
    },
    "0403/astro-ph0403137_arXiv.txt": {
        "abstract": "We report on the design, first observing season, and analysis of data from a new prototype millimeter-wave interferometer, MINT\\@. MINT consists of four 145\\,GHz SIS mixers operating in double-sideband mode in a compact heterogeneous configuration.  The signal band is subdivided by a monolithic channelizer, after which the correlations between antennas are performed digitally. The typical receiver sensitivity in a 2\\,GHz band is 1.4\\,mK$\\sqrt{\\mbox{s}}$. The primary beams are 0.45\\degr\\ and 0.30\\degr\\ FWHM, with fringe spacing as small as 0.1\\degr. MINT observed the cosmic microwave background (CMB) from Cerro Toco, in the Chilean Altiplano.  The site quality at 145\\,GHz is good, with median nighttime atmospheric temperature of 9\\,K at zenith (exclusive of the CMB).  Repeated observations of Mars, Jupiter, and a telescope-mounted calibration source establish the phase and magnitude stability of the system. MINT is the first interferometer dedicated to CMB studies to operate above 50\\,GHz.  The same type of system can be used to probe the Sunyaev-Zel'dovich effect in galaxy clusters near the SZ null at 217\\,GHz.  We give the essential features of MINT and present an analysis of sideband-separated, digitally sampled data recorded by the array. Based on 215 hours of data taken in late 2001, we set an upper limit on the CMB anisotropy in a band of width $\\Delta \\ell=700$ around $\\ell=1540$ of $\\delta T < 105\\ \\mu$K (95\\% conf.). Increased sensitivity can be achieved with more integration time, greater bandwidth, and more elements. ",
        "introduction": "Measurements of the CMB at small angular scales ($\\ell>1000$) complement those from {\\sl WMAP} \\citep{bennett03} and previous experiments including TOCO  \\citep{miller99}, BOOMERanG \\citep{mauskopf/etal:2000,ruhl/etal:2003}, MAXIMA \\citep{lee/etal:2001}, DASI \\citep{halverson/etal:2002}, VSA \\citep{grainge/etal:2003}, and ARCHEOPS \\citep{benoit/etal:2003}, allowing yet more precise constraints on the standard cosmological model \\citep{spergel03}.  Additionally, such measurements will enable us to understand the transition from the linear to the non-linear growth regimes for the formation of cosmic structure. One promising measurement technique is interferometry because of the ability to control systematic errors. Existing CMB interferometers, CAT~\\citep{robson93}, CBI~\\citep{padin01}, DASI~\\citep{leitch02}, IAC~\\citep{dicker00}, VSA~\\citep{watson03} are based on HEMT amplifiers~\\citep{posp92,posp97,posp95} and operate at 30\\,GHz where the contamination from point sources can be large; MMIC technology at 90\\,GHz will enable a straightforward switch to higher frequencies.  A compact interferometer using SIS (superconductor-insulator-superconductor) mixers at $f>100$\\,GHz has the potential to measure the CMB anisotropy and SZ effects where the point source contamination is at a minimum~\\citep{toffolatti} and where the SZ~effect \\citep{sunyaev70,sunyaev72} has a characteristic spectral null. However, the technical complications are considerable, as might be surmised by the existence of only a few such interferometers (e.g., SMA, OVRO, BIMA, IRAM, NRO/NMA).  We have built and observed with a four element, heterogeneous, compact, fully-digital, SIS-based D-band ($\\sim 145$\\,GHz) interferometer designed as a prototype for a larger, more ambitious instrument. Novel elements of the interferometer include the channelizer, a custom microwave integrated circuit; digital correlators for CMB measurements; and mixed antenna sizes. We describe key features of the interferometer, the observations in Chile, and the data analysis. ",
        "conclusions": "The Millimeter Interferometer has operated successfully as a prototype for a high-frequency CMB instrument.  MINT combines heterogeneous antennas, double-sideband SIS mixers, monolithic channelizers, and digital correlators.  All of these elements may be features in future CMB interferometers.  To make a competitive detection of the CMB anisotropy, such an instrument would require wider bandwidth, longer observing time, and more receivers (having more distinct baselines would also improve the sampling in $\\ell$-space). Usable bandwidth in the D-band is ultimately limited by atmospheric emission lines \\citep{matsushita}, but sideband-separating mixers---under development at NRAO \\citep{sb_sep_sis}---offer an instant factor of two improvement, given equal system noise temperatures. Consider an improved MINT, with 8\\,GHz bandwidth (4\\,GHz upper and 4\\,GHz lower sideband), four times the actual 2001 observing time, and fully working polarizers.  Assuming 50 separate fields of view and the power spectrum of the 2002 WMAP best fit $\\Lambda$CDM model \\citep{spergel03}, the MINT window functions produce CMB bandpowers of 1450 and 590 $\\mu\\mathrm{K}^2$ at effective $\\ell=900$ and $\\ell=1390$. In this case, the improved MINT would have the raw sensitivity to see anisotropy at the level of $7\\sigma$ and $6\\sigma$ in the short and medium baselines.  We conclude that the techniques described here constitute a promising, if challenging, method for observing CMB anisotropies at or above 100\\,GHz."
    },
    "0403/astro-ph0403301_arXiv.txt": {
        "abstract": "SN~1998bw, especially after the discovery of GRB~030329/SN~2003dh, seems to be the equivalent of the Rosetta stone for the SN/GRB connection. In this paper I review optical and near IR observations that have been carried out for this uncanny object, which has probably confirmed suspicions and ideas originally formulated in the early seventies of last century.   \\vspace{5mm} \\begin{quote} {\\it \\footnotesize Thus, the observation of gamma-ray line emission from a young supernova seems very promising in the near future. The observation, or even a null observation at a low threshold, will give significance in the fields of nuclear astrophysics and supernova theory. The scientific importance of a positive measurement would be analogous with and comparable to the importance of successful detection of neutrinos from the Sun.}  \\vspace{5mm} \\hspace{60mm}Clayton, Colgate \\& Fishman \\cite{clayton}. \\end{quote}    \\vspace{10mm}  This story probably begins in 1969, with what I like to call a prophecy, and it is right with it that I wish to start this review on the optical and near-IR observations of SN~1998bw\\footnote{This talk was given in Valencia on April 25, 2003. For some cabalistic reason, this coincided exactly with the fifth anniversary of GRB~980425/SN1998bw.}.  As J. Sollerman said in one of his papers on this striking object, SN~1998bw was born famous. And it was doomed to become even more famous as time went by, so famous that it was sometimes named {\\it the supernova of the century}. And this is indeed interesting, since it was born in the same century of SN~1987A, one of the most studied and referenced objects in the sky.  Just from the optical and near-IR observations point of view, this is witnessed by the large number of papers which have been published in the first four years. Starting with the {\\it Nature} papers by Galama et al. \\cite{galama} and Iwamoto et al. \\cite{iwamoto}, a number of authors have presented the results of their observational campaigns: McKenzie \\& Schaefer \\cite{mckenzie}, Galama et al. \\cite{galama99}, Stathakis et al. \\cite{stathakis}, Fynbo et al. \\cite{fynbo}, Sollerman et al. \\cite{sollerman00}, Patat et al. \\cite{patat01} and Sollerman et al. \\cite{sollerman02}.  The reader is referred to these papers for a detailed account on the observations, while here I will try to give only a general view of the SN~1998bw phenomenon.  SN~1998bw was discovered by Galama et al. \\cite{galama98a} in the BeppoSAX Wide Field Camera error box of GRB~980425 (Soffita et al. \\cite{soffita}, Pian et al. \\cite{pian}) close to a spiral arm of the barred galaxy ESO~184--G82 (see Fig.~\\ref{fig:contour}), by comparing two frames taken at the ESO New Technology Telescope on Apr 28.4 and May 1.3 UT.  Spectroscopic and photometric observations, both in the optical and in the near IR, started at ESO--La Silla immediately after the discovery, and showed that this object was profoundly different from all then known SNe (Lidman et al. \\cite{lidman}).  \\begin{figure} \\centering \\includegraphics[height=10cm]{fpatat_fig1.ps} \\caption{Contour plot of ESO-184-G82. The original image stack was obtained at VLT+FORS1 at about 900 days after the explosion. Spatial scale was computed for a distance of 40 Mpc. The SN position is marked.} \\label{fig:contour} \\end{figure}   Its peculiar spectroscopic appearance, its unusually high radio luminosity at early phases (Kulkarni et al. \\cite{kulkarni}), its optical luminosity ($M_V\\sim -$19.2 + 5 log $h_{65}$) and, in particular, the probable association with GRB~980425 through positional and temporal coincidence (Galama et al. \\cite{galama98a}, Pian et al. \\cite{pian}) placed SN~1998bw at the center of discussion concerning the nature of Gamma Ray Bursts. The object was tentatively classified as a peculiar Ic (Patat \\& Piemonte \\cite{patat98a}, Filippenko \\cite{filippenko98}), I would say {\\it by definition} more than anything else, due to the complete absence of H lines, the weakness of the Si~II 6355~\\AA\\/ line and no clear He~I detection in the optical spectra (see Fig.~\\ref{fig:firstspec}).  Main spectral features were identified as O~I, Ca~II, Si~II and Fe~II (Iwamoto et al. \\cite{iwamoto}). The estimated expansion velocities were exceptionally high ($\\sim$30,000 km s$^{-1}$) and this caused a severe line blending. The evolution during the first months was unusually slow compared to known Ic, with the nebular spectra still retaining many of the features present during the photospheric phase (Stathakis et al. \\cite{stathakis}, Patat et al. \\cite{patat01}).   \\begin{figure} \\centering \\includegraphics[height=10cm]{fpatat_fig2.ps} \\caption{Spectrum of SN~1998bw taken at ESO-La Silla on May 5, 1998.} \\label{fig:firstspec} \\end{figure}    The late onset of the fully nebular phase has been interpreted as an indication for a large ejected mass (Stathakis et al. \\cite{stathakis}) as it was predicted by the early light curve models. During the intermediate phase, the emission lines were definitely broader than in known Type Ib/c SNe and the simultaneous presence of iron--peak and $\\alpha$-elements indicated unusual relative abundances or physical conditions in the SN ejecta (Patat et al. \\cite{patat01}). The late spectroscopy presented by Sollerman et al. \\cite{sollerman00} showed that the tentative morphological classification of SN~1998bw as a Type Ic event was indeed appropriate. The main features have been identified as [O~I], Ca~II, Mg~I and Na~I~D, the latter possibly contaminated by He~I~5876~\\AA.  As far as the Gamma-ray burst is concerned, GRB~980425 was pretty weak, since the implied energy for a 40 Mpc distance was 8.1$\\pm$1.0 $\\times$10$^{47}$ erg (Pian et al. \\cite{pian}), which is definitely smaller than the usual 10$^{53}$ erg value typical for the so-called cosmological GRBs. This has led the community to believe that GRB~980425 is a member of an unusual class of GRBs/SNe (see for example Fynbo et al. \\cite{fynbo}).  The galaxy which hosted GRB~980425/SN~1998bw, ESO~184--G82, is an SBc galaxy with a recession velocity $v_r$=2532 km s$^{-1}$ (Patat et al. \\cite{patat01}). Its luminosity is $L\\sim$0.5$-$1.2 $L_{LMC}$, it is currently undergoing strong star formation, it shows a bar and rather clear indications of morphological disturbances (Fynbo et al.\\cite{fynbo}). This is clearly visible in the late VLT images, which show a possibly double nucleus, isophotal twisting and asymmetry (see Fig.~\\ref{fig:contour}). All this suggests that the observed star formation is related to galaxy interaction/merging.  The HST images have shown that SN~1998bw exploded in a star-forming region (Fynbo et al.\\cite{fynbo}), containing several bright and young stars within a projected distance of 100 pc. This is consistent with the progenitor of SN~1998bw being a young and massive star  Due to its relatively high apparent brightness, Kay et al. \\cite{kay} and Patat et al. \\cite{patat01} have attempted to perform some polarimetric measurements at different epochs. After correcting for the interstellar polarization in the direction of the host galaxy, the estimated optical linear polarization was 0.6\\% (day $-$7), 0.4\\% (day +10) and 0.5\\% (day +42). The fluctuations in these values seem to suggest that the observed polarization is intrinsic to the SN, although a dusty medium in the parent galaxy cannot be ruled out.  \\begin{figure} \\centering \\includegraphics[height=12cm,angle=-90]{fpatat_fig3.ps} \\caption{Comparison between SN~1998bw and other SNe at maximum light.} \\label{fig:compmax} \\end{figure}   The small degree of polarization at optical wavelengths can be explained in terms of a moderate departure from sphericity (axial ratio less than 2:1; H\\\"oflich et al. \\cite{hoflich}), either in the photosphere or in the outer scattering envelope when the line of sight is not coincident with an axis of symmetry.   Interestingly, no polarization, either circular or linear, was detected in the radio, and this has been interpreted as the signature of a spherically symmetric blast wave (Kulkarni et al. \\cite{kulkarni}). This apparent discrepancy might suggest that the radio and the optical radiation were generated in regions of different geometry. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/gr-qc0403036_arXiv.txt": {
        "abstract": "We study  gravitational collapse of rapidly rotating relativistic polytropes of the adiabatic index $\\Gamma = 1.5$ and $2$, in which the spin parameter $q \\equiv J/M^{2} > 1$ where $J$ and $M$ are total angular momentum and gravitational mass, in full general relativity. First, analyzing initial distributions of the mass and the spin parameter inside stars,  we predict the final outcome after the collapse. Then, we perform fully general relativistic simulations on assumption of axial and equatorial symmetries and confirm our predictions. As a result of simulations, we find that in contrast with the previous belief, even for stars with $q > 1$, the collapse proceeds to form a seed black hole at central region, and the seed black hole subsequently grows as the ambient fluids accrete onto it. We also find that growth of angular momentum and mass of the seed black hole can be approximately determined from the initial profiles of the density and the specific angular momentum.  We define an effective spin parameter at the central region of the stars, $q_{c}$, and propose a new criterion for black hole formation as $q_{c} \\alt 1$. Plausible reasons for the discrepancy between our and previous results are clarified. ",
        "introduction": "One of the fundamental issues in numerical general relativity is to explore the final fate after gravitational collapse of rotating stellar cores. If their mass is not as large as the maximum mass of a neutron star $\\sim 2 M_{\\odot}$, a rotating neutron star will be the outcome of the collapse. On the other hand, for stars of sufficiently large mass, the collapse will proceed completely and a spacetime singularity will be formed according to the singularity theorems of Hawking and Penrose\\cite{Singu0,Singu1}. If the cosmic censorship conjecture suggested by Penrose\\cite{Pen} is correct, any singularity should be surrounded by an event horizon. Then the black hole uniqueness theorems of Israel\\cite{Isr}, Carter\\cite{Car}, and Robinson\\cite{Rob} tell that a collapsed star will consequently settle down to a Kerr black hole. On the other hand, if this conjecture is not correct, a rotating stellar collapse might form a state with naked singularities.  It is well known that in a Kerr spacetime, the singularity is covered by an event horizon only if a non-dimensional spin parameter defined as $q \\equiv J/M^{2}$, where $J$ and $M$ are the angular momentum and the gravitational mass of the system, does not exceed unity\\cite{Singu0,Singu1}. Otherwise the singularity is naked. This implies that the value of $q$ for any black hole cannot be larger than unity. For the realistic progenitor of black holes, however, the value of $q$ may be larger than unity (see Sec. \\ref{Candidates}). Thus, it is interesting to explore the final fate after gravitational collapse of rotating stars with $q>1$. Numerical relativity is the unique approach to resolve this problem.  There have been several studies in full general relativity with regard to the above subject on the assumption of axial symmetry\\cite{Nak,NOK,SP,ACST,Shiba2000,S03,SS}. A series of simulations of rotating stellar collapse in full general relativity was first performed by Nakamura\\cite{Nak} and his collaborators\\cite{NOK} using the (2+1)+1 formalism developed by Maeda {\\it et al}.\\cite{MSNM}. They adopted differentially rotating massive stars that are to collapse (i.e., whose masses are much larger than the maximum allowed mass for a neutron star formation). An interesting finding in their simulations is that $q$ is an important parameter for determining the prompt black hole formation. Their results suggest that for $q>1$, no black hole is formed and the stars bounce back due to centrifugal force, indicating that the cosmic censorship conjecture holds. Stark and Piran \\cite{SP} performed simulations for the collapse of polytropes with the adiabatic index $\\Gamma=2$ and with an artificially given rigid rotation, using the Bardeen-Piran formalism \\cite{BP}. They reconfirmed the result found by Nakamura; $q \\sim 1$ is the demarcation between black hole formation and bounce. Abrahams {\\it et al}.\\cite{ACST} studied the collapse of axisymmetric tori consisting of collisionless matter, and also found that black holes form only from initial configurations with $q \\alt 1$. Shibata \\cite{Shiba2000} investigated the effects of rotation and shock heating on the criteria for prompt formation of black holes, for rotating equilibrium polytropes with $\\Gamma = 2$ and $q < 1$, and found that effects of rotation and shock heating raise the critical mass for black hole formation. Shibata \\cite{S03} also studied formation of black holes from marginally stable supramassive rotating neutron stars, modeled by $\\Gamma = 2.5$, $2.25$, $5/3$, and $1.5$ polytropes rotating at the mass shedding limit for which $q < 1$, and found that the final state of such collapse is a Kerr black hole with no appreciable disk. Shibata and Shapiro \\cite{SS} studied the collapse of a rigidly rotating polytrope at the mass-shedding limit with $\\Gamma=4/3$. The value of $q$ for such a configuration is close to unity as $\\approx 0.96$ \\cite{BS}. They indicate that the final state is a black hole surrounded by an appreciable disk ($M_{\\rm disk}/M \\approx 0.1$). All these results seem to suggest that $q \\approx 1$ is the maximum value for the black hole formation, and the final state after the gravitational collapse of rotating stars with $q < 1$ is a rotating Kerr black hole with a small disk mass.  However, the rotating stellar collapse in general relativity has been not sufficiently studied for {\\it realistic} stellar cores; e.g., iron stellar cores and pair-unstable oxygen cores for which $\\Gamma \\alt 4/3$ at the onset of the collapse. Recently, Shibata \\cite{ShibaSA} determined marginally stable and rigidly rotating stars with soft equations of state ($\\Gamma \\alt 4/3$) against gravitational collapse, which are plausible initial conditions for rotating stellar core collapse. He indicates that (i) even for a rigidly rotating star with $q > 1$, a central region, in which an approximate {\\it local} value of $q$ ($q_c$) is smaller than unity, will first collapse to form a black hole, unless $q$ is not too large ($q < 2.5$), and that (ii) as a result of the collapse for $q_c < 1$ and $q \\gg 1$, a massive disk ($M_{\\rm disk}/M=O(0.1)$) will be formed around the final black hole.  In this paper, we examine the validity of the conjectures indicated in \\cite{ShibaSA} by axisymmetric simulations in full general relativity. The essence of the points in \\cite{ShibaSA} is that the density profile for the equilibrium stars with $\\Gamma \\alt 4/3$ is significantly different from that for stiff equations of state, only for which previous works have paid attention \\cite{NOK,SP,Shiba2000,S03}. This fact could modify the criterion of black hole formation from previous studies. As a first step toward a more realistic simulation, in this paper, we focus on extracting a physical essence for the criterion of black hole formation using simple toy models. We perform simulations for rotating stellar collapse with a {\\it moderately} soft equations of state of $\\Gamma=1.5$ and with a stiff equation of state of $\\Gamma=2$ and compare the results. From the results of these simulations, we illustrate that the global value of $q$ is an adequate parameter for predicting the black hole formation only for the stiff equations of state but inadequate for the soft equations of state.  The paper is organized as follows. In Sec. II, we briefly describe our formulation, gauge conditions, and boundary conditions. In Sec. III, we first review plausible initial conditions for black hole formation in which the value of $q$ can be larger than unity. Then, we describe initial conditions adopted in this paper and predict the final outcomes of the stellar collapse for our initial models following \\cite{ShibaSA}. Sec. IV presents numerical results, emphasizing that the predictions made in Sec. III are correct. Sec. V is devoted to a summary. Throughout this paper, we adopt the geometrical units $G = c = 1$, where $G$ and $c$ denote the gravitational constant and the speed of light, respectively. We use the Cartesian coordinates $x^{k} = (\\, x, y, z\\,)$  as the spatial coordinates, with $r = \\sqrt{x^{2} + y^{2} + z^{2}}$. ",
        "conclusions": "We perform fully general relativistic simulations for rotating stellar collapse in axial symmetry using the same formulation as in~\\cite{S2003}, to which the reader may refer for details and basic equations. The fundamental variables for the hydrodynamics are: \\beqn \\rho &&:{\\rm rest~mass~ density},\\nonumber \\\\ \\varep &&: {\\rm specific~ internal~ energy}, \\nonumber \\\\ P &&:{\\rm pressure}, \\nonumber \\\\ u^{\\mu} &&: {\\rm four~ velocity}, \\nonumber \\\\ v^i &&={dx^i \\over dt}={u^i \\over u^t}, \\eeqn where subscripts $i, j, k, \\cdots$ denote $x, y$, and $z$, and $\\mu$ the spacetime components. As the variables to be evolved in the numerical simulations, we define a weighted density $\\rho_* = \\rho \\alpha u^t e^{6\\phi}$ and a weighted four-velocity $\\hat u_i = h u_i= (1+\\varepsilon+P/\\rho)u_i$ where $h$ denotes a specific enthalpy. {}From these variables, the total baryon rest mass and angular momentum of system, which are conserved quantities in axisymmetric spacetime, can be defined as \\beqn M_*&=&\\int d^3 x \\rho_*, \\\\ J  &=&\\int d^3 x \\rho_*\\hat u_{\\varphi} . \\eeqn General relativistic hydrodynamic equations are solved using a so-called high-resolution shock-capturing scheme \\cite{Font,S2003} on the $y=0$ plane with the cylindrical coordinates $(x, z)$ (in the Cartesian coordinates with $y=0$).  To model initial conditions we adopt the polytropic equations of state \\beq \\label{poly-EOS} P =  K \\rho^{1+\\frac{1}{n}}, \\eeq where $n$ is the polytropic index and $K$ polytropic constant. Then physical units enter the problem only through the polytropic constant $K$, which can be chosen arbitrarily or else completely scaled out of the problem. Thus, in the following we display only the dimensionless quantities which are defined as \\beq \\label{rel-nondim} \\bar{M}_{\\ast} = M_{\\ast}K^{-n/2}, \\ \\ \\bar{M} = M K^{-n/2}, \\ \\ \\bar{R} = R K^{-n/2}, \\ \\ \\bar{J} = J K^{-n}, \\ \\ \\bar{\\rho} = \\rho K^{n}, \\ \\ \\bar{\\Omega} = \\Omega K^{n}, \\eeq where $M$, $R$, and $\\Omega$ denote an ADM mass, a radius, and an angular velocity. Hereafter, we adopt the units of $K = 1$ so that we omit the bar.  On the other hand, during the time evolution, we use the so-called $\\Gamma$-law equations of state of the form \\beq P = (\\Gamma -1)\\rho \\epsilon, \\eeq where the adiabatic index $\\Gamma$ is set as $1+ 1/n$. In the absence of shocks, no heat is generated and the collapse proceeds in an adiabatic manner, preserving the polytropic form of the equations of state.  We neglect effects of viscosity and magnetic fields. The timescale of dissipation and angular momentum transport due to these effects are much longer than the dynamical timescale, unless the magnitude of viscosity or magnetic fields is extremely large\\cite{BSS}. Thus neglecting them is an appropriate assumption.  The fundamental variables for geometry are: \\beqn \\alpha &&: {\\rm lapse~function}, \\nonumber \\\\ \\beta^k &&: {\\rm shift~vector}, \\nonumber \\\\ \\gamma_{ij} &&:{\\rm metric~in~3D~spatial~hypersurface},\\nonumber \\\\ \\gamma &&=e^{12\\phi}={\\rm det}(\\gamma_{ij}), \\nonumber \\\\ \\tilde \\gamma_{ij}&&=e^{-4\\phi}\\gamma_{ij}, \\nonumber \\\\ K_{ij} &&:{\\rm extrinsic~curvature}. \\eeqn As in the series of our papers, we evolve $\\tilde \\gamma_{ij}$, $\\phi$, $\\tilde A_{ij} \\equiv e^{-4\\phi}(K_{ij}-\\gamma_{ij} K_k^{~k})$, and trace of the extrinsic curvature $K_k^{~k}$ together with three auxiliary functions $F_i\\equiv \\delta^{jk}\\pa_{j} \\tilde \\gamma_{ik}$ with an unconstrained free evolution code as in \\cite{AMS,SBS,bina1,bina2,SN,Shiba2000,S2003}.  The Einstein equations are solved in the Cartesian coordinates. To impose axisymmetric boundary conditions, the Cartoon method is used \\cite{alcu}: Assuming the reflection symmetry with respect to the equatorial plane, simulations are performed using a fixed uniform grid with the grid size $N \\times 3 \\times N$ in $(x, y, z)$ which covers a computational domain as $0 \\leq x \\leq L$, $0 \\leq z \\leq L$, and $-\\Delta \\leq y \\leq \\Delta$. Here, $N$ and $L$ are constants and $\\Delta = L/N$. In the Cartoon method, the axisymmetric boundary conditions are imposed at $y= \\pm \\Delta$. Details are described in \\cite{S2003} to which the readers may refer.  As the slicing condition we impose an ``approximate'' maximal slicing condition in which $K_k^{~k} \\approx 0$ is required\\cite{AMS}. As the spatial gauge, we adopt a dynamical gauge condition \\cite{DynGauge}. In the present work, the equation for the shift vector is written as \\beq \\pa_t \\beta^k = \\tilde \\gamma^{kl} (F_l +\\Delta t \\pa_t F_l), \\label{dyn} \\eeq where $\\Delta t$ denotes a time step in numerical computation\\cite{S03}. Note that in this gauge condition, $\\beta^{i}$ obeys a hyperbolic-type equation for sufficiently small value of $\\Delta t$ because the right-hand side of the evolution equation for $F_i$ contains a vector Laplacian term \\cite{SN}. The outstanding merit of this gauge condition is that we can save computational time significantly. It has already been found that stable simulations for rotating stellar collapse and merger of binary neutron stars are feasible in this gauge \\cite{bina2,S03}.  An outgoing wave boundary condition is imposed for $h_{ij}(\\equiv \\tilde \\gamma_{ij}-\\delta_{ij})$, $\\tilde A_{ij}$, and $F_{i}$ at the outer boundaries of the computational domain. The condition adopted is the same as that described in \\cite{SN}. On the other hand, for $\\phi$ and $K_k^{~k}$, outer boundary conditions are imposed as $r\\, \\phi = {\\rm const}$ and $K_k^{~k}=0$, respectively.   A black hole may be formed as a result of collapse. We determine the location of it using an apparent horizon finder developed in \\cite{ShibaAH}. As the system approaches a stationary state, the apparent horizon will approach the event horizon. In a dynamical spacetime we compute the apparent horizon mass $M_{\\rm AH}$ which is defined as\\cite{YP} \\beq M_{\\rm AH} = \\sqrt{\\frac{A}{16\\pi}} , \\eeq where $A$ denotes area of an apparent horizon.  During the numerical simulations,  conservation of mass and angular momentum and a violation of the Hamiltonian constraint, $H_{\\rm error}$, are monitored as code checks. Here $H_{\\rm error}$ is evaluated with a weighted average by $\\rho_*$ as \\begin{equation} H_{\\rm error} \\equiv \\frac{\\displaystyle \\int \\rho_{\\ast} \\left| V \\right| d^{3}x} {\\displaystyle \\int \\rho_{\\ast} d^{3}x} = \\frac{1}{M_{\\ast}} \\int \\rho_{\\ast} \\left| V \\right| d^{3}x, \\end{equation} where $V$ is defined as \\begin{equation} V \\equiv \\frac{\\displaystyle \\; \\tilde \\Delta \\psi - \\frac{1}{8}\\psi \\tilde{R} + 2\\pi \\psi^{5} \\left(\\rho h w^{2} - P \\right) + \\frac{\\psi^{5}}{8}\\tilde{A}_{ij}\\tilde{A}^{ij} - \\frac{\\psi^{5}}{12}(K_k^{~k})^{2}\\; } {\\displaystyle \\; \\left|\\tilde \\Delta \\psi\\right| + \\frac{1}{8}\\left|\\psi\\tilde{R}\\right| + 2\\pi \\psi^{5} \\left(\\rho h w^{2} - P \\right) + \\frac{\\psi^{5}}{8}\\tilde{A}_{ij}\\tilde{A}^{ij} + \\frac{\\psi^{5}}{12}(K_k^{~k})^{2} \\; }. \\label{eq11} \\end{equation} In Eq. (\\ref{eq11}), $\\psi \\equiv e^{4\\phi}$, $w \\equiv \\alpha u^t$, and $\\tilde \\Delta$ and $\\tilde{R}$ are the Laplacian and the Ricci scalar with respect to the conformal metric $\\tilde{\\gamma}_{ij}$. Numerical results for several test calculations, including stability and collapse of spherical and rotating neutron stars, have been described in~\\cite{S2003}.    We have reported new results about black hole formation in collapse of rapidly rotating stars with $q > 1$ and with moderately soft ($\\Gamma = 1.5$) and stiff ($\\Gamma = 2.0$) equations of state, analyzing initial conditions and performing axisymmetric simulations in full general relativity.  The initial conditions are given, following Stark and Piran \\cite{SP}; we first give marginally stable spherical polytropes with $\\Gamma = 1.5$ or $\\Gamma = 2.0$ and then artificially add an angular momentum. Adopting the same method described in \\cite{ShibaSA}, we have predicted that (I) inner region in which $q_{c} \\alt 1$ will collapse first to form a seed black hole even if the global value of $q$ is much larger than unity; (II) the formed seed black hole will subsequently grow as the ambient fluids accrete onto it; (III) evolution of the relation between the rest mass $m_{\\ast}(j)$ and angular momentum $J(j)$ enclosed inside a growing black hole will agree approximately with the initial relation between $m_{\\ast}(j)$ and $J(j)$; (IV) whether black hole is formed or not can be determined by $q_{c}$ irrespective of $q$; (V) the final outcome of dynamical collapse will be a black hole surrounded by an appreciable disk.  To confirm these conjectures, we performed fully general relativistic hydrodynamic simulations using a high-resolution shock capturing scheme with the $\\Gamma$-law equations of state. As a result of numerical simulations, we confirmed the conjectures (I), (II), (III), and (IV). {}From these results we conclude that the previous criterion for black hole formation that no black hole is formed for $q \\agt 1$ is no longer valid in particular for soft equations of state. As a new criterion for no black hole formation, we propose the condition $q_{c} \\agt 1$.  Associated with conjecture (II), we have found that the process of black hole formation is divided into two phases: the first phase is that a seed black hole is formed and the other is the accretion phase that a part of the ambient fluids are swallowed into the formed seed black hole. At the time of the accretion phase, the system is composed of a black hole and a surrounding massive disk accreting onto the central black hole. In the collapse of massive iron cores, such accretion phase may be associated with the gamma ray bursts \\cite{GRB}.  In this paper, we have focused only on the criterion of black hole formation analyzing simple toy models. To obtain a scientific result that can be compared with observational data such as gravitational waves, however, a simulation with a realistic initial condition and a realistic equation of state is necessary. As we indicate in this paper, a rapidly rotating black hole will be formed after rapidly rotating massive stellar core collapse and pair-unstable collapse, if the progenitors are not extremely rapidly rotating at the onset of collapse as $q_c < 1$. Formation of a rapidly rotating black hole with $q \\sim 1$ and accretion of a large mass onto such black hole of a large value of $q$ are likely to be strong burst sources of gravitational waves for laser interferometric detectors. Thus, performing realistic numerical simulations of black hole formation is an important subject for predicting the gravitational waveforms. We plan to attack these computations in a fully general relativistic manner extending previous works \\cite{S-Sekig1,S-Sekig2}.  For $q_c > 1$, a black hole will not be formed promptly. In such case, the collapse leads to formation of a self-gravitating disk or torus. They will be subsequently unstable against nonaxisymmetric deformation. After the nonaxisymmetric instabilities turn on, angular momentum will be transported from the inner region to the outer region, decreasing the value of $q(j)$ around the inner region below unity. As a result, a seed black hole may be formed. To follow these processes, it is necessary to perform a numerical simulation without assuming the axial symmetry. During the collapse, the typical length scale may change by a factor of $10^4$ from $R / M \\sim 10^4$ to 1. To follow the collapse by numerical simulation, very large computational resources will be necessary and, thus, the simulation for this phenomenon will be one of the computational challenges in the field of numerical relativity.   \\vspace{4mm}  \\begin{center} {\\bf Acknowledgments} \\end{center}  Numerical computations were performed on the FACOM VPP5000 machine in the data processing center of National Astronomical Observatory of Japan. This work is in part supported by Japanese Monbu-Kagakusho Grants (Nos. 14047207, 15037204, and 15740142)."
    },
    "0403/astro-ph0403071_arXiv.txt": {
        "abstract": "{The transition between the Main Sequence and the Red Giant Branch in low mass stars is powered by the onset of the CNO burning, whose bottleneck is the $^{14}$N(p,$\\gamma)^{15}$O. The LUNA collaboration has recently improved the low energy measurements of the cross section of this key reaction. We analyse the impact of the revised reaction rate on the estimate of the Globular Clusters ages, as derived from the turnoff luminosity. We found that the age of the oldest Globulars should be increased by about 0.7-1 Gyr with respect to the current estimates. ",
        "introduction": "Globular Clusters (GCs) represent the oldest resolved stellar populations. Their age practically coincides with the time elapsed since the epoch of the formation of the first stars in the Universe and provides an independent check of the reliability of standard (and non standard) cosmological models. Moreover, the age spread in the GC system is a primary indicator of the time scale of the halo formation. Among the various methods to date stellar Clusters, the most reliable and widely adopted is that based on the measurement of the luminosity of the turnoff (i.e. the bluest point on the main sequence). This dating technique requires the knowledge of the Cluster distance, the light extinction along the line of sight and the chemical composition (see \\cite{Gra.2} for an exhaustive analysis of the present state of the art). In addition, a reliable theoretical calibration of the turnoff luminosity-age relation (TOL-A) is needed. This relies on our knowledge of the physical processes of energy generation (e.g. nuclear reactions) and transport (e.g. opacity) taking place in H burning low mass stars. An adequate description of the thermodynamics of stellar matter is also required. Finally we have to consider any mechanism capable to modify the internal chemical stratification (once again nuclear reactions, convective mixing, rotational induced mixing, microscopic diffusion or levitation induced by radiation pressure). \\cite{chab} discuss the influence of various theoretical uncertainties on the calibration of the turnoff luminosity-age relation and conclude that the total uncertainty due to the theory may be confined within 0.5 Gyr.  This paper is devoted to the evaluation of the impact on the theoretical calibration of the Globular Clusters ages of the improved determination of the rate of the key reaction $^{14}$N(p,$\\gamma)^{15}$O, as obtained by the LUNA collaboration (\\cite{stra.20}). Since thermonuclear reactions are responsible for chemical modifications occurring in stellar interiors and supply most of the energy irradiated from the stellar surface, the estimated stellar lifetime depends on accurate measurements of their rates. In the last few years, many efforts have been spent in improving these measurements at energy as close as possible to the Gamow peak, namely the relevant energies at which nuclear reactions take place in stars. This is a mandatory requirement for the calibration of stellar ages.  The main sequence stars presently observed in Globular Clusters have masses smaller than that of the Sun. As in the Sun, these low mass stars burn H in the center, mainly through the pp chain. However, towards the end of their life, when the central hydrogen mass fraction becomes smaller than about 0.1, the nuclear energy released by the H burning becomes insufficient and the stellar core must contract to extract some energy from its own gravitational field. Then, the central temperature (and the density) increases and the H burning switches from the pp chain to the more efficient CNO cycle. Thus, the departure from the Main Sequence is powered by the CNO cycle, whose bottleneck is the $^{14}$N(p,$\\gamma$)$^{15}$O reaction. The luminosity of the turnoff depends on the rate of this key reaction: the larger the rate, the fainter the turnoff. On the contrary, the total lifetime is only marginally affected by a change in the CNO, because it is mainly determined by the rate of the $^1$H(p,$e^+,\\nu_e)^2$H. As a consequence, an increase of the CNO rate would imply fainter turnoff points, for a given age, or younger ages, for a given turnoff luminosity (see also \\cite{chab98}). Note that an equivalent effect may be caused by the enhancement of the CNO abundances (\\cite{rood}, \\cite{stra.6}).  In the next section we remind the new measurements of the stellar cross section of the $^{14}$N(p,$\\gamma$)$^{15}$O reaction. Then, in section 3 we present the revised turnoff luminosity-age relation (TOL-A). We show that this revision leads to systematically larger estimates of the age of the Globular Clusters. Implications for cosmology are briefly discussed in the conclusive section. ",
        "conclusions": ""
    },
    "0403/astro-ph0403592_arXiv.txt": {
        "abstract": "We have detected sub-TeV gamma-ray emission from the direction of the Galactic Center (GC) using the CANGAROO-II Imaging Atmospheric Cherenkov Telescope (IACT). We detected a statistically significant excess at energies greater than 250\\,GeV. The flux was one order of magnitude lower than that of Crab at 1~TeV with a soft spectrum $\\propto E^{-4.6\\pm0.5}$. The signal centroid is consistent with the GC direction and the observed profile is consistent with a point-like source. Our data suggests that the GeV source 3EG~J1746$-$2851 is identical with this TeV source and we study the combined spectra to determine the possible origin of the gamma-ray emission. We also obtain an upper limit on the cold dark-matter density in the Galactic halo. ",
        "introduction": "Many observations have been made of the Galactic Center (GC) region over a wide range of wavelengths \\citep[e.g.,][]{pauls76,sofue86,pedlar89,koyama96,buckley97, purcell97,mh98,maeda02}. At the dynamical center of the Galaxy is the source Sgr~A$^*$, containing a 4$\\times 10^6$\\,M$_\\sun$ black hole, which is a strong, variable, radio source \\citep[e.g.,][]{zhao03} and a weak, but occasionally flaring, infra-red \\citep{genzel} and X-ray \\citep{baganoff} source. Surrounding Sgr~A$^*$ are the (presumed) supernova remnant Sgr~A~East, and a number of filamentary radio structures, including one known as the Arc, extending over tens of parsecs. One filamentary structure was recently found to be a non-thermal X-ray source \\citep{sakano}. The radio emission is interpreted as synchrotron radiation from high-energy electrons in the ambient magnetic fields. The observed magnetic fields are as high as several mG, although whether they are distributed over the whole GC region or are localized is not known \\citep{morris}.  The bright EGRET source 3EG~J1746$-$2851 is listed as unidentified in the third EGRET catalog \\citep{hartman99}, but as the position, given in 3EG as ($l=0.11, b=-0.04$) with an 95\\% confidence contour radius of 0.13 degrees, is consistent with the GC, it has been considered as the gamma-ray counterpart to GC region \\citep{mh98}. Three different source scenarios were suggested by \\cite{mh98}: emission from one or more pulsars, inverse Compton radiation from relativistic electrons in the Arc, or Cold Dark Matter (CDM) annihilation. N-body simulations suggest a significant enhancement of the CDM density around galaxy cores \\citep{navarro96}. Better constraints on CDM can be obtained by high-energy gamma-ray observations of the GC region. More recently, 3EG~J1746$-$2851 has been associated with Sgr~A~East, with the gamma-ray emission arising from the decay of neutral pions produced by high-energy protons, accelerated in the remnant, interacting with the ambient matter \\citep{fm03}. Time variability was reported, although its amplitude is not large \\citep{no03}.  Spectra derived from the EGRET data \\citep[][shown in Fig. 3]{mh98,hartman99} are not well fit by a single power-law: the spectra are quite hard at low energies but flatten, or possibly turn over, at high energies. Simple extrapolations of the high energy spectra, however, suggest a potentially detectable TeV flux.  We have observed the GC region with the CANGAROO-II Imaging Atmospheric Cherenkov Telescope (IACT). The 10\\,m diameter CANGAROO-II telescope \\citep{kawachi} detects gamma-rays above several hundred GeV by detecting optical Cherenkov radiation generated by relativistic secondary particles in the cascades produced when high-energy gamma-rays (and background cosmic rays) interact with the Earth's upper atmosphere. The optical images provide information on the direction and energy of gamma-ray events. The telescope is located near Woomera, South Australia (136$^\\circ$47$'$\\,E, 31$^\\circ$06$'$\\,S) and thus the GC culminates within $\\sim$2$^\\circ$ of the zenith, enabling observations with a lower energy threshold. Observations with the Whipple IACT from Mt Hopkins (31$^\\circ$58$'$\\,N) yielded an 2\\,$\\sigma$ upper limit of $0.45\\times 10^{-11}$\\,cm$^{-2}$\\,s$^{-1}$ at 2\\,TeV \\citep{buckley97}.  The angular resolution of CANGAROO-II is 0.24$^\\circ$ (36\\,pc at a distance of 8.5\\,kpc) with an energy threshold of 400\\,GeV in the case of a Crab-like energy spectrum ($\\propto E^{-2.5}$). For the softer spectrum, $\\propto E^{-4.6}$, we find (see \\S\\,4), the angular resolution becomes 0.32$^\\circ$ (47\\,pc at 8.5\\,kpc) and the energy threshold is lowered to 250\\,GeV. With this angular resolution we cannot resolve Sgr~A$^*$ ($\\sim$1\\,pc in size) and the elliptical Sgr~A~East (major axis of 10.5\\,pc), which surrounds Sgr A$^*$ \\citep{yusef,maeda02}. The location of the EGRET source 3EG~J1746$-$2851 is also too close to the GC for us to resolve it. ",
        "conclusions": "Due to the complex structure of the GC region different scenarios for the origin of the EGRET gamma-ray flux have been considered \\citep[see, e.g.,][and references therein]{mh98}. As mentioned in \\S1, \\cite{fm03} have recently attributed the GeV emission to $\\pi^0$ decay resulting from high-energy protons interacting with the ambient matter in Sgr~A East. They argued that due to synchrotron cooling in the high average magnetic field, primary (accelerated) and secondary leptons would have a steep spectrum, and their contribution to the high energy gamma-ray flux via inverse Compton emission and bremsstrahlung would be small. The leptonic contribution to the TeV gamma-ray flux would be even smaller. Indeed, the synchrotron cooling time can be estimated as $\\tau_s=16\\times B_{mG}^{-2} E_{e,TeV}^{-1}$~yr. Thus for leptons of a few TeV, which might contribute to the sub-TeV gamma-ray flux, the cooling time in a mG magnetic field is even less than those considered by \\cite{fm03}.  Another factor which will reduce the amount of secondary TeV leptons, the key products from {\\it pp} interactions,  is the presence of a cut-off in the accelerated proton spectrum at a few TeV, suggested by this observation (see below). In the case of Sgr~A$^*$, an even higher magnetic field is expected, which would lead to the same conclusion.  In Fig.~\\ref{fig3} we reproduce the $\\pi^0$ decay emission to fit both the EGRET data and ours. The input proton spectrum was assumed to be proportional to $E^{-\\gamma}e^{-E/E_{max}}$. The number density of the ambient gas ($n$) was taken to be $\\sim 10^3$ cm$^{-3}$ \\citep{maeda02}, however this affects only the estimation of the total proton cosmic-ray energy, as $E_c\\propto n^{-1}$. The various lines shown in Fig.~\\ref{fig3} are for differing values of $\\gamma$ and $E_{max}$. The GeV spectrum given in the 3EG catalog \\citep{hartman99} differs from that of the dedicated analysis of \\cite{mh98}, however, in both cases the EGRET and CANGAROO-II data can be relatively smoothly connected, with a cutoff energy of 1--3\\,TeV. The spectra for $\\gamma=2.4$ yield slightly worse fits to the data at the lowest energies. The inferred total cosmic-ray energy greater than 1\\,GeV (see Fig.~\\ref{fig3}) corresponds to $\\leq 10$\\% of a typical supernova energy, a quite plausible conversion efficiency. The details, however, are dependent on the EGRET flux. In the case of a lower gas density ($n$), the cosmic ray energy exceeds that of a single supernova, and we should consider the possibility of other sources in the GC region, such as Sgr A$^*$ itself, contributing to the gamma-ray flux. Our data determine the maximum energy to which cosmic rays are accelerated in this region, which does not depend upon other physical parameters. For example, the range of values is consistent with recent theoretical predictions for shock acceleration in supernova remnants \\citep{pz03}.  Thus, identifying the TeV source with 3EG\\,J1746$-$2851, and associating these in turn with Sgr~A~East and/or Sgr A$^*$, we find that the gamma-ray fluxes and spectra can be naturally explained as arising due to $\\pi^0$ decay.  We can, in addition, use our observations to derive upper limits to the CDM abundance in the GC region, assuming the GeV and TeV emission is centered on Sgr~A$^*$ and following the method of \\cite{enomoto253}. The emission region is assumed to be a sphere with a radius of 47\\,pc. The annihilation mode of $\\chi \\chi \\rightarrow q\\bar{q}$ was first considered. The cross section ($\\sigma$) multiplied by the relative velocity ($v$) and the branching fraction ($B$) is a free parameter that was normalized to $10^{-26} cm^3s^{-1}$. The experimental data of $e^+e^-\\rightarrow q\\bar{q}\\rightarrow \\gamma X$ was used for the gamma-ray multiplicity \\citep{lep}. The derived 2\\,$\\sigma$ upper limits for the CDM densities ($\\rho_{CDM}^{47pc}\\sqrt{\\sigma vB/10^{-26} cm^3s^{-1}}$) are 9300, 7300, 5800, 5300, and 5800\\,GeV/cm$^3$ for an assumed WIMP mass of 0.7, 1, 2, 4, and 6\\,TeV, respectively. Here, although we took the EGRET flux as an upper limit, above 0.6\\,TeV it is the CANGAROO-II data that dominates these results. \\cite{navarro96} and subsequent analyses \\citep{fukushige,moore,ghigna,jing,power} adopted a cusp structure, i.e., the density profile proportional to $r^{-\\beta}(1+r/r_s)^{\\beta-3}$, where $r_s$ is a free parameter and $\\alpha$ ranges from 1 to 1.5. In the case of $\\beta$=1.3, the upper limits of local densities ($\\rho_{CDM}^{\\odot}\\sqrt{\\sigma vB/10^{-26} cm^3s^{-1}}$) are of the order of several GeV/cm$^3$. Higher upper limits are obtained when the decay mode of $\\chi \\chi \\rightarrow \\gamma \\gamma$ was assumed with $\\sigma vB_{\\gamma \\gamma}=10^{-29} cm^3s^{-1}$."
    },
    "0403/astro-ph0403253_arXiv.txt": {
        "abstract": "We present the age and mass distribution of star clusters in M51. The structural parameters are found by fitting cluster evolution models to the spectral energy distribution consisting of 8 {\\em HST-WFPC2} pass bands. There is evidence for a burst of cluster formation at the moment of the second encounter with the companion NGC5195 (50-100 Myr ago) and a hint for an earlier burst (400-500 Myr ago). The cluster IMF has a power law slope of -2.1. The disruption time of clusters is extremely short ($< 100$ Myr for a 10$^4 M_{\\odot}$ cluster). ",
        "introduction": "$N$-body models (Salo \\& Laurikainen 2000) suggest 2 interactions between M51 and the companion NGC5195. The age distribution shows clear evidence for a burst $\\sim$50-100 Myr ago and there is a hint for another burst at the epoch of the early encounter (see Fig.~1). The age distribution is complete for the last Gyr only for clusters with $M_{\\rm cl} > 10^{4.7} M_{\\odot}$ for which two bursts are clear. This makes it difficult to confirm the burst for all cluster masses.  \\begin{figure}[!h] \\plotone{gieles1_fig1.ps} \\caption{Age histogram in absolute numbers (left) and in number of clusters per Myr (right). The shaded regions are the predicted moments of encounter with NGC5195.} \\end{figure} ",
        "conclusions": ""
    },
    "0403/astro-ph0403065_arXiv.txt": {
        "abstract": "The completed 2dF QSO Redshift (2QZ) Survey has been used to search for extreme large-scale cosmological structure ($\\sim$200\\,h$^{-1}$\\,Mpc) over the redshift range $0 < z < 2.5$.  We demonstrate that statistically significant overdensities and underdensities do exist and hence represent the detection of cosmological fluctuations on comoving scales that correspond to those presently detected in the cosmic microwave background.  However, the fractional overdensities on scales $>100$\\,h$^{-1}$\\,Mpc are in the linear or only weakly non-linear regime and do not represent collapsed non-linear structures.  We compare the measurements with the expectation of the $\\Lambda$CDM model by measuring the variance of counts in cells and find that, provided the distribution of QSOs on large scales exhibits a mild bias with respect to the distribution of dark matter, the observed fluctuations are found to be in good agreement with the model.  There is no evidence on such scales for any extreme structures that might require, for example, departures from the assumption of Gaussian initial perturbations.  Thus the power-spectrum derived from the 2QZ Survey appears to provide a complete description of the distribution of QSOs. The amount of bias and its redshift dependence that is required is consistent with that found from studying the clustering of 2QZ QSOs on $\\sim$10\\,h$^{-1}$\\,Mpc scales, and may be adequately described by an approximately redshift-invariant power spectrum with normalisation $\\sigma_8 \\simeq 1.0$ corresponding to a bias at $z=0$ of $b \\simeq 1.1$ rising to $b \\simeq 2$ at the survey's mean redshift $z \\simeq 1.5$. ",
        "introduction": "It has long been recognised that active galaxies and QSOs may used to probe structure in the universe on the largest observable scales, and statistical analysis of the clustering of samples of such cosmological objects was first developed thirty years ago (e.g. \\citealt{webster76,seldner78}).  Even in those early studies it became apparent that, although the universe appeared generally homogeneous on the largest scales, individual clusters of active galaxies could be identified \\citep{webster82}.  In modern ideas of active galaxies and of their link to the process of galaxy formation, it is now thought that luminous QSOs lie at the heart of the most massive elliptical galaxies (\\citealt{kukula01,dunlop03,floyd04}) and that, provided the clustering bias expected for such a population is accounted for \\citep{sheth99}, studying the clustering of active galaxies is a means of testing the predictions of hierarchical galaxy formation theory both on the largest cosmological scales and as a function of cosmic epoch.  Whether non-uniform distributions of QSOs on large scales could in fact be detected with a high degree of statistical confidence was an open question for many years, however, largely because of the limited sample sizes and systematic uncertainties and varying selection biases present in samples of QSOs.  A statistically significant detection of QSO clustering on $\\sim$10\\,Mpc scales was achieved primarily from the AAT QSO redshift survey \\citep{shanks94}.  With the advent of the much larger 2dF QSO Redshift Survey (2QZ: \\citealt{croom01a,smith04,croom04}) not only has an extremely secure detection of QSO clustering been achieved, but also the shape of the correlation function, its bias with respect to the underlying dark matter distribution, the cosmic evolution in that bias, and the luminosity dependence of QSO clustering have now all been measured on scales out to about 80\\,Mpc (\\citealt{croom01b,croom02,croom03,croom04b}).  On larger scales, the shape of the power spectrum of QSO clustering has been measured and shown to be consistent with the expectations for a low density CDM universe \\citep{outram03}. The aim of this paper is to look at the distribution of QSOs in redshift space, rather than Fourier; to investigate on what scales departures from randomness may be detected; and to test whether those departures are consistent with the expectation of the widely-assumed cosmological model in which structures have grown by gravitational collapse of initially Gaussian fluctuations in a universe dominated by a combination of cold dark matter and dark energy.  Over the years there have been a number of claims of the detection of extremely large (50--200\\,Mpc) groupings of AGN and QSOs \\citep{crampton87,clowes91,williger02,brand03}: at the time of these claims there was no consensus on the cosmological model or on the relationship of the distribution of QSOs to the distribution of matter, and it was hoped that the measurement of such large groupings might constrain models of one or both of these.  Now, however, there have emerged in the literature preferences for one particular cosmological model \\citep{spergel03} and for the bias of QSOs measured by the 2dF QSO Redshift Survey \\citep{croom01b,croom02,croom03,croom04b}, so it is timely to ask whether the very large scale distribution of QSOs is consistent with these recent developments. If such large-scale clusters were shown to be collapsed structures with high values of overdensity this would be likely to be in significant disagreement with the cosmological model and measurements of QSO bias. In fact, we shall see in this paper that indeed very large groupings of QSOs may be identified, possibly on scales as large as 300\\,Mpc, but that the groupings are still in the linear regime of gravitational collapse and are in fact entirely in accord with current expectations. ",
        "conclusions": "Fluctuations in QSO space density on scales $\\sim 200$\\,h$^{-1}$\\,Mpc may be seen directly in the QSO distribution (Figs\\,\\ref{fig:SN50}-\\ref{fig:Bayesianmap}). It is likely that these fluctuations may be traced out to larger scales, and inspection of the maps and analysis of the variance of the QSO density field indicates the detection of fluctuations on scales possibly as large as 300\\,h$^{-1}$\\,Mpc.  The fluctuations are in good agreement with those expected in a $\\Lambda$CDM cosmology with WMAP parameters if we assume that QSO clustering is a biased tracer of dark matter fluctuations, with a bias value that is approximately redshift independent and which is close to the value inferred from QSO clustering on 10\\,Mpc scales.  Since it appears that $b > 1$, the fluctuations in mass associated with the observed QSO fluctuations are certainly in the linear or weakly non-linear regime of gravitational collapse, and hence they do not represent collapsed non-linear structures, but are merely a reflection of the large-scale fluctuations expected given the $\\Lambda$CDM power spectrum.  Figs\\,\\ref{fig:SN50}-\\ref{fig:Bayesianmap} represent direct detection of structure in the distribution of discrete objects on comoving scales that correspond to the scales measured in the cosmic microwave background by WMAP \\citep{spergel03}.  Consideration of two methods of correcting the measurements for observational selection shows that on the largest scales there may remain systematic variations in survey uniformity at levels that are comparable to the measured fluctuations.  For this reason we believe that the 2QZ survey should not be used to attempt to infer more precise values for cosmological parameters such as $\\Omega_M$ until those systematic non-uniformities can be better corrected for.  One of the most significant features of the results presented here is that there is no evidence for any non-Gaussian initial conditions, which might have been implied had any of the $> 50$\\,Mpc structures previously claimed for QSO distributions turned out to represent collapsed structures. As the counts in cells are completely in accord with the simulations generated from a $\\Lambda$CDM power spectrum, we conclude that, as far as this test is concerned, the power spectrum remains a complete description of the large-scale distribution of QSOs.  We have not in this paper attempted to measure the topology of the observed fluctuations. This would be an interesting exercise in order to test whether the observed structures are consistent in every way with the expectations of linearly collapsing structures forming from Gaussian initial conditions.  \\vspace*{5mm}  \\noindent {\\large \\bf ACKNOWLEDGEMENTS}\\\\  The 2dF QSO Redshift Survey was based on observations made with the U.K.~Schmidt and Anglo-Australian Telescopes.  We thank all the present and former staff of the Anglo-Australian Observatory for their work in building and operating the 2dF facility."
    },
    "0403/astro-ph0403315_arXiv.txt": {
        "abstract": "Inflation universally produces classical almost scale free Gaussian inhomogeneities of any light scalars. Assuming the coupling constants at the time of inflation depend on some light moduli fields, we encounter the generation of modulated cosmological fluctuations from (p)reheating. This is an alternative mechanism to generate observable (almost) scale free adiabatic metric perturbations. We extend this idea to the class of hybrid inflation, where the bifurcation value of the inflaton is modulated by the spatial inhomogeneities of the couplings. As a result, the symmetry breaking after inflation occurs not simultaneously in space but with the time laps in different Hubble patches inherited from the long-wavelength moduli inhomogeneities. To calculate modulated fluctuations we introduce techniques of general relativistic matching conditions for metric perturbations at the time hypersurface where the equation of state after inflation undergoes a jump, without evoking the detailed microscopic physics, as far as it justifies the jump. We apply this theory to the modulated fluctuations from the hybrid and chaotic inflations. We discuss what distinguish the modulated from the inflation-driven fluctuations, in particular, their spectral index, modification of the consistency relation and the issue of weak non-Gaussianity. ",
        "introduction": "One of the  generic predictions of inflation lies in the fact that vacuum fluctuations of all light scalar fields, $\\chi_a$, minimally coupled to gravity and with a mass smaller than the Hubble parameter ($m \\ll H$), are universally unstable and appear after inflation as classical random a priori Gaussian inhomogeneities with (almost) scale free spectrum ${H}/{\\sqrt{2k^3}}$.  The wavelengths of the fluctuations, $\\delta \\chi_a(t, \\vec x)$, of such a light scalar field are stretched by inflation and exceed the Hubble patch after inflation.  One can relate $\\delta \\chi_a(t, \\vec x)$ to the cosmological scalar metric perturbations in different ways, depending on the composition of the underlying theory. Indeed, the simplest and most studied possibility is to assume that there is a single light scalar field that is the inflaton itself, $\\varphi$. The inflaton fluctuations $\\delta \\varphi$ are transferred to the scalar metric perturbations through gravitational interaction~\\cite{scalar}.  On general grounds, we may however expect many scalar fields playing roles during inflation.  There is a broad range of multiple fields inflationary models with different motivations behind them, like e.g. double inflation~\\cite{KLS85} or hybrid inflation~\\cite{hybr}, where different fields dominate at different stages of the cosmological evolution.  Multiple fields were also evoked to design departures from the standard inflationary predictions: existence of isocurvature modes~\\cite{hybr,iso}, non-scale free spectrum of primordial fluctuations~\\cite{non-flat} or deviation from Gaussianity~\\cite{ng,bu}. In these cases some fields non-necessarily give dominant contribution to the background evolution, but during some time give dominant contribution to the perturbations.  Another corner of the multiple field parameter space is related to the curvaton scenario~\\cite{curvaton}. There, a newly introduced scalar field, the curvaton, that is indeed not the inflaton, should be light during inflation, plus dominate after inflaton decay, plus decay after inflaton but prior BBN, plus gives a dominant contribution to the metric fluctuations.  A new and more economic idea to generate cosmological perturbations from modulated fluctuations of couplings was proposed recently~\\cite{modul0,modul1,modul2}.  In the context of multiple scalar fields theories, it is assumed in these models that some of the light fields never give dominant contribution neither into background nor in perturbations during inflation, but contribute to the coupling constants \\begin{equation}\\label{coupl} \\alpha=\\alpha(\\chi_a) \\ . \\end{equation} Indeed, in string theory~\\cite{strings} couplings are in fact the VEV of moduli fields and, in SUSY theories couplings can also depend on scalars. As a result, fluctuations of the moduli $\\delta \\chi_a$ generated during inflation, will manifest themselves in spatial inhomogeneities of the couplings \\begin{equation}\\label{coupl1} \\delta \\alpha= \\frac{\\partial \\alpha}{\\partial \\chi_a} \\, \\delta \\chi_a \\ . \\end{equation} The interaction is not as important during inflation as it is after inflation. Consider for example simple chaotic inflation. The background scalar field rolls towards the minimum of the potential, where it begins to oscillate. Due to the coupling to the other fields the inflaton decays into radiation in the process of (p)reheating~\\cite{KLS94}. Because of the large scale coupling inhomogeneities, Eq.~(\\ref{coupl1}), at scales much larger than the causal horizon after inflation, transition from the matter dominated regime of inflaton oscillations to radiation occurs in different causal patches not simultaneously, which leads to small adiabatic metric perturbations after (p)reheating.  After the moduli $\\chi_a$ get pinned down to their minima, the spatial variations of coupling constants in the late time universe will be erased.  However, the large scale metric fluctuations that are produced due to interactions survive as a memory of the primordial moduli inhomogeneities.  It is actually not mandatory for scalar fields to be light during inflation. In fact, in the context of $N=1$ supergravity with the minimal K{{{\\\"a}}}hler potential during inflation the scalars typically acquire the mass $m_a \\sim H$. Cosmological fluctuations neither in inflaton sector nor in moduli sector are produced unless a special care is taken to make at least some of them light.  This, together with the options to built up inflationary models, motivate the study of different mechanisms of generation of cosmological perturbations.  In this article, we investigate this mechanism of {\\it modulated cosmological fluctuations}.  Original papers on the modulated fluctuations \\cite{modul1,modul2} discussed metric perturbations generated from the decay of inflaton oscillations.  However, there is another important class of hybrid inflationary models, which typically emerges in supergravity and the string theory/brane cosmology.  The two field $(\\varphi, \\sigma)$ hybrid inflation has the effective potential \\begin{equation}\\label{veff} V_{\\rm eff} = \\frac{\\lambda}{4}\\left(\\sigma^2 - v^2\\right)^2 + \\frac{1}{2}g^2\\varphi^2\\sigma^2 + V(\\varphi) \\end{equation} where $V(\\varphi)$ is the inflaton potential.  This potential contains two couplings $\\lambda$ and $g^2$, which define the end-point of inflation and the field dynamics.  Assuming these couplings are moduli dependent, we encounter modulated fluctuations in the hybrid inflation.  This is the main novel idea of the paper which we describe and develop in Section~\\ref{sec_hybrid}.  One may think of different aspects of the theory of modulated fluctuations, related to the nature of the moduli and dependencies, Eq.~(\\ref{coupl}), details of preheating, thermalization, moduli evolution etc.  Here we concentrate on the cosmological General Relativistic part of the story, namely, how to derive cosmological metric perturbations from the coupling perturbations, Eq.~(\\ref{coupl1}). In the chaotic inflation typically inflaton oscillations decay through the non-perturbative effect of parametric resonance of particle creation~\\cite{KLS94}. This leads to copious production of particles in out-of-equilibrium state. Further interactions between particles relax them towards thermal equilibrium. Perturbative regime of particle interactions takes place at the latest stages of transition from inflation towards hot radiation~\\cite{Felder:2000hr}. In hybrid inflation, preheating after inflation has the character of the tachyonic preheating~\\cite{tachprech}, accompanying the symmetry breaking after inflation~\\cite{tachprech2}. Tachyonic preheating leads to generation of particles out-of-equilibrium with subsequent thermalization of them. Coupling constants of interaction appears at different stages of preheating and thermalization. However, for the production of modulated cosmological perturbations it is essential that couplings are responsible for the change of the equation of state.  If the equation of state is not changing, modulated perturbations are not generated.  In this respect we shall concentrate at the first instance where equation of state after inflation is changing due to the couplings. In the case of chaotic inflation it happens when an effective matter equation of state of the coherent inflaton oscillations is replaced by the radiation in the very fast process of preheating.  In the case of hybrid inflation, the vacuum-like (inflation) equation of state is replaced by radiation in the very quick process of tachyonic preheating.  So far the formalism of the modulated cosmological fluctuations was considered for the toy model of slow perturbative reheating \\cite{modul1,modul2,mata,vernizzi}, often with the Yukawa type interaction $\\phi \\bar \\psi \\psi$ between inflaton and fermions. This is fair enough to see in principle that metric fluctuations are generated from the moduli inhomogeneities. However, the modern theory of the transition between inflation and radiation is described by the theory of preheating. Even for case of  fermions their production is non-perturbative and significantly modified by preheating \\cite{fermion}.  In this paper, on the technical side, we suggest the method to treat generation of modulated cosmological fluctuations in the inflationary models with preheating. We will use the fact that in all cases preheating is very short, and from the general relativistic point of view can be considered as the instant jump of the equation of state. It is then convenient to use the formalism of matching conditions of the geometrical quantities at the time of the transition~\\cite{Deruelle:1995kd}. For this, we do not need to know the microscopic details of preheating and thermalization.  In Section~\\ref{sec_hybrid}, we first discuss the basics of the modulation of the couplings and some model building aspects that have to be fulfilled for such a mechanism to be efficient. After recalling the basics of the matching conditions in Section~\\ref{Sec-match}, we then apply them to the standard single field inflationary case in Section~\\ref{sec_infl}.  We then turn to the case of modulated fluctuations in Section~\\ref{sec_modfluc}. We will discuss the main features of the mechanism in Section~\\ref{sec_concl} and in particular emphasize that it allows to extend the standard consistency relation of inflation. ",
        "conclusions": "Modulated curvature fluctuations in hybrid models can be put on firm grounds whether it is from a model building perspective on for the computation of its phenomenological consequences. That such models are generic might reveal interesting at different levels. In such classes of models indeed the metric fluctuations are not necessarily associated to field fluctuations in the slow roll direction. It allows to break the relations between the slow roll parameters and the shape of the power spectrum. Analogous conclusions were reached in Ref.~\\cite{picon} where a model that decouples the spectral index from the inflationary stage is presented. Moreover some of the difficulties encountered in the usual hybrid models might be circumvent. Indeed the production of topological defects at an energy scale comparable to that of the adiabatic fluctuations, no contribution of which have been detected, is natural in such models forcing a tuning of the parameters subsequently causing the suppression of ``good\" scalar fluctuations (see Ref.~\\cite{Pterm}). Here we avoid this problem. Actually what we have obtained here is a kind of decoupling of the background evolution sector form the generation of fluctuations.  Are thus models arbitrary constructions or could they be physically motivated? Let us recall that light fields, such as moduli, are generic in e.g. string theory. They may have different effects related to the theory of structure formation and of interest for cosmology. They may induce a modulation of the coupling constants that can lead (i) to a spacetime modulation of the constants of nature, as e.g. the fine structure constant (see Ref.~\\cite{const2} for a the effect of a fluctuating light dilaton, Ref.~\\cite{const3} for an example of signature on the cosmic microwave background and Ref.~\\cite{const1} for a review), (ii) to the generation of non-Gaussianity and (iii) to the mechanism of modulated fluctuations described in this article.  The mechanism presented in the article shows that modulated inflationary models correspond to a large class of models that naturally emerges from supergravity where the couplings depending on light moduli fields can modulate the transition to the radiation era. In such models we have more specifically shown that gravitational waves were not affected by the modulation while scalar modes receive an extra-contribution. The primordial power spectra however follow a modified consistency relation, Eq.~(\\ref{lala}), which depends on a new parameter, $\\gamma$, that characterizes the dependence of the bifurcation point on the light fields. Moreover the fact that the contribution of the gravitational waves is always smaller in this context than in the standard inflationary case has two consequences: (i) it allows to have a higher energy scale for inflation and still have undetectable tensor modes but (ii) makes the possibility of their detection via $B$-polarization measurements more difficult. An interesting situation arises when $\\gamma\\sim 1$. In that case, the tensor modes still have a relative amplitude comparable to the standard case but (i) we have a deviation, Eq.~(\\ref{lala}), from the standard, Eq.~(\\ref{consist}), that can be hoped to be measured and (ii) the scalar power spectrum is the sum of two power laws, opening the door for a possible break in the range of wavelengths of interest for cosmology.  These results should be put in parallel to other extensions of standard inflation where the consistency relation can be modified. Generically in multifield inflation~\\cite{2field}, Eq.~(\\ref{consist}) becomes $T/S=\\varepsilon\\sin^2\\Theta$ where $\\Theta$ is the isocurvature-adiabatic correlation angle. As in our case, $T/S<\\varepsilon$. On the other hand in the curvaton scenario~\\cite{curvaton}, the density perturbations arise from the fluctuation of a light field and primordial perturbations are entirely from an isocurvature mode and the consistency relation becomes $T/S\\ll\\varepsilon$ (see e.g. Ref.~\\cite{vl}). This is analogous to a limit in which $\\gamma\\gg1$.  This class of models also opens a new phenomenological avenue. Indeed nothing prevent the fields $\\chi_a$ to develop non-Gaussianity prior to the transition as in the mechanism of Ref.~\\cite{bu}. In modulated inflation the transfer of modes is not not due to a bent of the trajectory but to the modulation. The outcome of such a mechanism could be the superposition of Gaussian perturbations and a non-Gaussian ones of the same variance. Let us stress that no non-Gaussianity is generated by the transition itself so that if fluctuations of the light field are initially Gaussian so will be the modulated fluctuations.  The last remarks are about the nature of modulated scalar fluctuations. First, note that this mechanism offers an example of situation in which the adiabatic mode is not conserved, an issue raised in Ref.~\\cite{adiab}. Second, before the end of inflation, the moduli field fluctuations are subdominant gravitationally and have a character of isocurvature fluctuations.  When the equation of state is changed the moduli field remains to be gravitationally subdominant (contrary to the curvaton scenario where dominance of the curvaton is additionally assumed). However, the moduli field fluctuations vary the time hypersurface where equation of state is altered.  At this moment their isocurvature fluctuations are transferred into adiabatic, scalar metric fluctuations (vis the coupling which control the change of the equation of state).  We have cosmological examples where initially isocurvature inhomogeneities are transferred into adiabatic fluctuations -- examples include isocurvature CDM scenario or curvaton scenario -- in both cases the carrier of isocurvature fluctuations become gravitationally dominant.  What is qualitatively new in the modulated fluctuations is the fact that the carrier of fluctuations remains always gravitationally subdominant, however, its isocurvature fluctuations are transformed into the adiabatic one after it modulates the jump in the equation of state."
    },
    "0403/astro-ph0403123_arXiv.txt": {
        "abstract": "In this model a collimated ultra-relativistic ejecta  collides with an amorphous dense cloud surrounding the central engine, producing gamma-rays via synchrotron process. The ejecta is taken as a standard candle, while assuming a gaussian distribution in thickness and density of the surrounding cloud. Due to the cloud high density, the synchrotron emission would be an instantaneous phenomenon (fast cooling synchrotron radiation), so a GRB duration corresponds to the time that the ejecta takes to pass through the cloud. Fitting the model with the observed bimodal distribution of GRBs' durations, the ejecta's initial Lorentz factor, and its initial opening angle are obtained as $\\Gamma_{0}\\lesssim 10^{3}$, and $\\zeta_{0} \\approx 10^{-2}$, and the mean density and mean thickness of the surrounding cloud as $\\overline{n} \\sim 3 \\times 10^{17} cm^{-3}$ and $\\overline{L} \\sim 2 \\times 10^{13} cm$. The clouds maybe interpreted as the extremely amorphous envelops of Thorne-Zytkow objects. In this model the two classes of long and short duration GRBs are explained in a unique frame. ",
        "introduction": "Undoubtedly, gamma-ray bursts (GRBs) have remained to be one of the most exiting, intriguing, and enigmatic astrophysical phenomena since their mysterious discovery in the past several decades (for a recent expository review of GRBs the interested reader is referred to the excellent work by J. I. Katz \\cite{katz}). Although no two GRBs resemble each other and each one has its own peculiarities which makes the problem of modeling GRBs very difficult, the whole of GRBs reveals several interesting features. Since the publication of the BATSE data \\citep{1st} which included the observation of over $200$ GRBs and revealed an almost uniform distribution of the location of GRBs in the sky, combined with the deficiency of faint GRBs, the association of GRBs with the galactic plane has been ruled out. The successive publications confirmed the figure more and more \\citep{3th,4th}. However, since the observation of afterglows in X-ray \\citep{Costa}, optical \\citep{van Paradijs} and radio spectrum \\citep{Frail} and the advert of Robotic Optical Transient Search Experiment (ROTSIE) telescope \\citep[e.g.]{Akerlof,Gisler} which has revealed the red shifts for several GRBs, their cosmological origin is widely accepted.\\\\ The BATSE data \\citep{1st,3th,4th} has also revealed another equally important overall feature of the GRBs. The distribution of time duration in observed GRBs shows a double heap distribution, which the smaller one peaks around 0.2 sec and the larger one peaks around 20 sec (Fig.1). This two peaked distribution which apparently separated GRBs into the so called short duration and long duration ones was referred to as $\\textit{bimodality}$ \\citep{kouveliotou,Norris} and led some investigators to believe that there are two distinct populations of GRBs.\\\\ It has been widely believed that whatever the central engine is, the radiation reaching us originates from the space surrounding the central engine. It is also believed that during the collapse the energy release streams out in relativistic confined ejectas (not isotropically) and thus the total energy release during each event is far less than the unbelievable amount that one might obtain by assuming isotropic radiation \\citep{kulkarni}.\\\\ The aim of this paper is to present a rather simple model, based on these general ideas, and show that there is no need for assuming two distinct populations of GRBs, and to show that once the geometrical considerations and cosmological effects are fully accounted for a genetic standard candle ejecta, crossing an amorphous dense cloud, the so called $\\textit{bimodality}$ can be deduced.\\\\ In \\S~2 the model and its general formulation is introduced. In \\S~3 the computational results and the fit of model parameters with BATSE data will be presented as well. \\S~4 is devoted to a discussion on the results. Some needed calculations and discussions are presented in appendices. ",
        "conclusions": "\\subsection{The GRB Source} Though the original shapes of equations (13) and (14) are naturally normalized, the resulting final equations (19) and (23) are not, because:\\\\ 1)the bursts that their symmetry axes make an angle $\\theta>\\zeta_{rad}(\\eta_{_{L}})$ can not be detected (see appendix C),\\\\ 2)the bursts for which $\\eta_{m}<\\eta_{_{L}}$ were not considered in the numerical computations (because they do not manage to cross out the cloud).\\\\ The total probability of observing the occurred bursts, $\\int_{-3}^{3} (dp / d\\log T_{\\oplus}) ~ d\\log T_{\\oplus}$, is obtained to be $1.47 \\times 10^{-6}$, using our best fit parameters(Eqn.[26]). This small probability must be interpreted to be due to the above two reasons. The first reason describes the suppression of observed GRBs by the term $ (\\Omega / 4\\pi)\\sim \\zeta_{0}^{2} / 2 \\sim 10^{-4}$, while the second is responsible for the remaining factor of $10^{-2}$. So, only about one percent of the bursts manage to produce a real GRB, and only about $10^{-4}$ of these GRBs occur in our line of sight. Of course, the lateral spreading of the ejecta+swept mass may modify these two factors, increasing the first and decreasing the second.\\\\ With the values for the fitted parameters in equation(25) or in equation(26), the mean mass of the clouds $\\overline{M}\\equiv\\frac{4}{3}\\pi \\overline{n} m_{p}\\overline{L}^{3}$ is about $1.2 \\times10^{34}\\:gr \\approx 6 M_{\\odot}$, which is of the order of the envelop mass in massive stars. Such amorphous clouds seem strange in stars, but in close neutron star-supergiant binaries where the neutron star orbits around the core and accrete the envelope, the spherical symmetry is likely removed, as pointed out by Podsiadlowski et al. \\cite{podsiadlowski}. Terman, Taam, \\& Hernquist \\cite{terman} show that the system would emerge to form a red supergiant with a massive Thorne-Zytkow Object (TZO) \\citep{thorne}. Podsiadlowski et al. \\cite{podsiadlowski} also estimated a TZO birth-rate of $\\geq 10^{-4} yr^{-1}$ in the galaxy. Considering equation(D4) and the obtained total probability of observing a \"real\" GRB, which is $1.47\\times 10^{-6}$ in our model (see above), it is seen that the GRB observation rate ($2$ events per day) implies a total (real+failed) GRB rate of the order of $\\sim 10^{-3}Mpc^{-3} yr^{-1}$. This is not too far from what Podsiadlowski et al. \\cite{podsiadlowski} theoretically estimated for TZO birth-rate ($\\geq 10^{-4} galaxy^{-1} yr^{-1}$). Qin et al. \\cite{qin} introduced AICNS (Accretion-Induced Collapse of Neutron Stars) scenario as GRB engines. Katz \\cite{katz3} introduces a dense cloud model to explain the observed $Gev$ gamma-rays in a number of GRBs \\citep[e.g.]{dingus,jones}, and suggests the amorphous envelops of TZOs for these clouds.\\\\ The initial opening angle of the ejecta in our model ($\\zeta_{0}\\simeq 10^{-2}rad\\approx 0.6 ^{\\circ}$) is obtained during the fitting process (\\S~3). Such a small opening angle for an ejecta or a jet might be explained by attributing it to the collimating process of an ultra-relativistic ejecta with $\\Gamma_{0}>1/\\zeta_{0}$ in a sufficiently high magnetic field \\citep{begelman}. The initial Lorentz factor of the ejecta, as obtained in our model ($\\Gamma_{0}\\simeq 10^{3}$, see Eqn.[25]) provides the necessary condition of $\\Gamma_{0}>1/\\zeta_{0}$. Aside from these theoretical justifications for the idea of a highly collimated ejecta at the source, there are some pieces of evidences supporting this idea \\citep{lamb,waxman,granot2}. \\subsection{The Bimodality} At the mean time the observed bimodality must be interpreted as a result of the second reason expressed in \\S~4.1. Though in our model one may expect only one heap in the $log T_{\\oplus}$ distribution associated with the directions in the clouds having both the most probable $L$ and $n$ (which are equal or near to $\\overline{L}$ and $\\overline{n}$), but as the numerical computation shows, such directions in the clouds are too thick and too dense to be crossed out by the ejecta, and therefore a real GRB would not be produced. So, we are left with four other regions with high probabilities:\\\\ 1)  $n\\sim\\overline{n}$ and $L<\\overline{L}$ (short GRBs)\\\\ 2)  $n\\sim\\overline{n}$ and $L>\\overline{L}$ (no GRBs)\\\\ 3)  $n>\\overline{n}$ and $L\\sim\\overline{L}$ (no GRBs)\\\\ 4)  $n<\\overline{n}$ and $L\\sim\\overline{L}$ (long GRBs)\\\\ Now, we show that the directions trough the clouds associated with region(1) produce the short duration GRBs, and ones associated to regions (2) and (3) produce no GRB, while the others in the forth region produce the long duration ones.\\\\ As to the equation(9) the quantity $l(n/E)$ (which is of the order of the sedov length) is $\\sim 7\\times 10^{11}cm \\ll \\overline{L}$ for $n\\sim \\overline{n}$ (see Eqn.[25]). As seen in Fig.7, the time duration of GRBs associated to such directions is of the order of the time duration of short GRBs (Case (1) above). Fig.8 shows that in our model the calculated time duration of GRBs for $n\\sim 10^{-6}\\overline{n}$ is of the order of the duration of long GRBs. Using equation(9), we see that in this case ( $n\\sim 10^{-6}\\overline{n}$ ), $l(n/E)$ is about $5\\times 10^{13}cm \\sim \\overline{L}$. So in our model the long duration GRBs are due to the passing of ejecta through the directions where $n \\ll \\overline{n}$ and $L \\sim\\overline{L}$ (case (4) above). Furthermore, Since $l(n/E) \\ll \\overline{L}$ when $n\\sim\\overline{n}$, we see that in cases (2) and (3) above, the ejecta would stop in the dense cloud and the produced photons can not scape from the optically thick cloud. In Fig.9, $T_{rec}(L;n,\\theta)$ is plotted for a number of densities, when $\\eta_{_{L}}=L / l(n/E)$ has the highest permitted value $\\eta_{m}$.\\\\ These general features of our calculations result from the general features of our model, and therefore we speculate that any distributions for the clouds thickness and density which are picked around a mean value could explain the general features of the duration distribution. Our chose of gaussian distributions for thickness and density was only for the few parameters needed to describe them.\\\\ \\subsection{The Opacity} In \"The Dense Cloud Model\" at this stage we have simply omitted the bursts which ejectas can not go out of the cloud and stop in it (because the emitted photons would be scattered by the dense cloud), and we have claimed that the produced photons by all bursts that succeed to cross out the dense cloud can finally enter the free space.\\\\ At the first glance the model might appear to have a serious problem in opacity, namely, in this model the best fit mean density and mean thickness of the clouds are found to be of the orders of $n\\sim 10^{17}cm^{-3}$ and $L\\sim 10^{13}cm$. So, as to the relation $\\tau_{op}=\\sigma_{_{T}}\\:n\\:L$, ($\\sigma_{_{T}}=6.65\\times10^{-25}cm^{2}$) one would expect an optical depth of the order of $10^{6}$. But there are two factors that remedy the situation:\\\\ (i) As the numerical calculation shows (see \\S~4.2), the most probable directions characterized by $n\\sim\\overline{n}$ and $L\\sim\\overline{L}$ are dynamically too thick to be crossed out by the ejecta and the radiation produced in this case would be completely scattered by the dense cloud. On the other hand, as discussed in \\S~4.2, the long duration GRBs are due to the crossing of ejecta through directions where $n\\sim 10 ^{-6}\\overline{n}$ and $L\\sim\\overline{L}$. since the density is reduced by the factor $10^{-6}$, the optical depth drops to the order of $1$. As explained in \\S~4.2, the short duration GRBs are due to the directions in the cloud where $n\\sim\\overline{n}$ and $L\\sim 10^{-2}\\overline{L}$. In this case the optical depth of the cloud reduces to $\\sim10^{4}$, which is yet too high. But,\\\\ (ii) most of photons emitted off the shock front have the chance to be overtaken by the moving shock (appendix B). Moreover, it has been shown \\citep{Momeni} that in high temperatures $kT\\sim10^{6}\\:m_{e}c^{2}$ the cross section of Compton scattering for instance for $Mev$ photons effectively drops to $\\sim 10^{-6}\\:\\sigma_{_{T}}$, so such a high temperature plasma is much more transparent than what may seem at first. The same is true for the case of a power law distribution of electrons (see Eqn.[A4]). Most of the electrons in such a distribution have energies of the order of $\\gamma_{e,min}m_{e}c^{2}\\sim 2 \\times 10^{5}m_{e}c^{2}$, in which we have used equation(A5), taking $p=2.5$, $\\xi_{e}=1/3$ and $\\Gamma=10^{3}$. For this case the Compton cross section suppresses by the factor of $\\sim10^{-3}-10^{-4}$ for $100ev-1Kev$ photons (the energy of the emitted photons in the shocked medium comoving frame is less than their observed energy by the factor of $\\Gamma$), and therefore, the optical depth drops to the order of $1$. So the photons overtaken by the shocked medium may remain in it without being scattered until the shocked medium crosses up the dense cloud. Briefly, due to this second factor, the optical depth in short duration case diminishes to $\\tau_{op}\\sim1$, and the optical depth in long duration case diminishes to $\\tau_{op}\\ll1$.\\\\ The variability seen in the light curves of long duration GRBs maybe attributed to the heterogeny of cloud's density in the ejecta's trajectory. This maybe the case in long duration GRBs for which $\\tau_{op}\\ll1$, but in short duration GRBs for which $\\tau_{op}\\sim1$, we expect that such a heterogeny would not be appeared.\\\\ The lateral expansion of the shocked medium can reduce its density and therefore its optical depth. This effect of course may cool the shocked medium and cause an increase in effective Compton cross section in it. The consideration of dynamical feedback makes our modeling much more complicated and it has been neglected here. In \"The Dense Cloud Model\" at this stage we have simply omitted the bursts which ejectas can not go out of the cloud and stop in it, but have claimed that the produced photons by all bursts that succeed to cross out the dense cloud can finally enter free space.\\\\  A complete study needs to include flux computations. An exact comparison with observed time duration data would be possible only when the theoretical time duration appearing in this work were exactly evaluated in the same manner as the observable duration $T_{90}$ is defined (as the time interval over which $5$ percent to $95$ percent of the burst counts accumulate). Moreover, the BATSE's triggering mechanism made it less sensitive to short GRBs than to long ones and therefore short GRBs were detected to smaller distances \\citep{mao,cohen,katz2}, so a smaller number of them have been observed. Lee \\& Petrosian \\cite{petrosian} studied this effect and corrected the number of short GRBs. In a more exact work this correction must be considered too.  We acknowledge the anonymous referee for valuable comments. F. S. acknowledges Mr. Mehdi Haghighi for his guides on computational methods. This research has been partly supported by Grant No. NRCI 1853 of National Research Council of Islamic Republic of Iran. \\newpage \\appendix"
    },
    "0403/astro-ph0403409_arXiv.txt": {
        "abstract": "I present an observational review of the five accretion-driven millisecond X-ray pulsars currently known, focusing on the results obtained with the {\\it Rossi X-ray Timing Explorer} ({\\it RXTE}) satellite.  A prominent place in this review is given to the first such system discovered, SAX J1808.4--3658. Currently four outbursts have been detected from this source, three of which have been studied using {\\it RXTE}. This makes this source the best studied example of all accretion-driven millisecond pulsars. Its October 2002 outburst is of particular interest because of the discovery of kilohertz quasi-periodic oscillations and burst oscillations during its thermonuclear X-ray bursts.  The other four accreting millisecond pulsars were discovered within the last two years and only limited results have been published so far. A more extended review can be found at http://zon.wins.uva.nl/$\\sim$rudy/admxp/ ",
        "introduction": "Ordinary pulsars are highly-magnetized ($B\\sim 10^{12}$ G), rapidly rotating ($P\\sim 10$ ms) neutron stars which spin down on timescales of 10 to 100 million years due to magnetic dipole radiation. By contrast, millisecond ($P<10$ ms) radio pulsars have ages of billions of years and weak ($B\\sim 10^{8-9}$ G) surface magnetic fields.  Since many of these millisecond radio pulsars are in binary systems, it has long been suspected (see, e.g., \\cite{bvdh1991} for an extended review) that the neutron stars were spun up by mass transfer from a stellar companion in a low-mass X-ray binary (LMXB). If this scenario is correct, LMXBs should harbor a fast spinning neutron star which might be visible as an accretion-driven millisecond X-ray pulsar. However, before the launch of the {\\itshape Rossi X-ray Timing Explorer} ({\\itshape RXTE}), all searches for coherent millisecond X-ray pulsations in LMXBs failed to yield a detection \\cite[and references therein]{v1994}. Only after {\\it RXTE} was launched did we obtain conclusive evidence that at least some LMXBs indeed harbor weakly magnetic neutron stars with millisecond spin periods.  In April 1998 the first accreting millisecond X-ray pulsar was discovered \\cite{wvdk1998} followed by the discovery of four additional systems during the last two years \\cite{markwardt2002,galloway2002,markwardt2003_1807,markwardt2003_1814}. ",
        "conclusions": "From the current review it is clear that {\\it RXTE} is vital to the discovery and study of accretion-driven millisecond X-ray pulsars. Thanks to {\\it RXTE} we know about the existence of five such systems. The detailed studies performed with {\\it RXTE} for those systems have yielded break-throughs in our understanding of kHz QPOs and burst oscillations. Furthermore, three of these accreting pulsars are in ultrashort binaries which will constrain evolutionary paths for this type of system (e.g., see \\cite{np2003}). However, it is also clear that the five systems do not form a homogeneous group; their pulsation frequencies span the range between 185 Hz and 435 Hz (Fig.~\\ref{fig:pulsations}), their orbital periods between 40 minutes and 4.3 hrs, and their X-ray light curves are very different (Fig.~\\ref{fig:lightcurves}). More well studied outbursts of the currently known systems are needed as well as discoveries of additional systems. At the moment, only {\\it RXTE} is capable of performing the necessary timing observations. After {\\it RXTE} the need for an instrument with at least similar or better capabilities is highly desirable for our understanding of accretion-driven millisecond pulsars and their connection with the non-pulsating neutron-star LMXBs.  For lack of space this review has focused on observational findings on accretion-driven millisecond pulsars during outbursts. For information on their quiescent states and on theoretical progress, I refer to \\cite{wijnands2003_review}.     \\begin{theacknowledgments} I thank Craig Markwardt for kindly providing the {\\it RXTE}/PCA data on XTE J1807--294 which I used to make Figures~\\ref{fig:pulsations} and \\ref{fig:broad-band}. I also thank Steve van Straaten for help in making Figures~\\ref{fig:1808_lc_2002}, \\ref{fig:1808_2002_1Hzflares}, and \\ref{fig:lightcurves}. \\end{theacknowledgments}"
    },
    "0403/hep-th0403208_arXiv.txt": {
        "abstract": " ",
        "introduction": "Theories with large extra dimensions provide a new framework for addressing the gauge hierarchy problem \\cite{led}. There are examples how such frameworks may arise from string theory compactifications \\cite{string,hw}. Much interest has been devoted to the models with an exponentially warped extra dimension \\cite{rs}, where the hierarchy arises from the gravitational redshift due to the curvature in the bulk rather than from the sheer size of the extra dimensions. These models can be linked with AdS/CFT correspondence in string theory \\cite{cft}, especially in the case when the extra dimension is of infinite proper size, but ends on a brane in the UV \\cite{rs2}. In this case, one encounters a new mechanism for generating 4D gravity out of infinite, noncompact extra dimensions. By respecting boundary conditions on the UV brane consistent with 4D general covariance, one finds quite generally that there is a normalizable gravitational mode localized on the UV brane, whose exchange generates 4D gravitational force \\cite{rs2,lykran}. The localized graviton mode persists for a large class of intrinsic geometries on the UV brane, most notably for de Sitter brane \\cite{bent}, lending to the construction of interesting cosmologies. This phenomenon of gravity localization does not depend on the codimension of the UV brane. Four-dimensional gravity can also be localized on intersections of codimension-one branes on a higher-codimension brane in an $AdS_{4+n}, n> 1$ environment \\cite{addk}. Some simple extensions of the example \\cite{addk} were considered in \\cite{csn}. However, one needs to consider the general case where all the branes have nonzero tension in order to address issues of possible vacua, stability, cosmological evolution and general multi-brane setups \\cite{bent}, \\cite{nkal}-\\cite{crystals}.  In this paper we explicitly derive a general class of solutions describing 4-branes intersecting on a tensionful 3-brane in a locally $AdS_6$ environment, such that the intrinsic 3-brane geometry is maximally symmetric, being Minkowski, de Sitter or Anti-de Sitter. These exhaust the vacua on the 3-brane. In the case of Minkowski and de Sitter they localize 4D gravity to the intersection. The solutions resemble an infinitely tall $4$-sided pyramid. We find them by cutting and pasting sections of $AdS_6$ bulk, such that the 4-branes reside at the seams and the 3-brane at the tip. The bulk cosmological constant is really a constant, rather than a step potential, since the $AdS$ patches between the branes are locally identical, and hence we can orbifold the configuration by its discrete symmetries. Our solutions can be straightforwardly generalized to the case of $AdS_{4+n}$, $n>2$, where $n$ $n+2$-branes intersect on a 3-brane, with a more general relation between brane tensions and angles between them. For the solutions with Minkowski metric along the intersection of 4-branes with identical tensions, we find the 4D graviton zero mode localized to the intersection, explicitly solving for its wavefunction, and we compute the couplings of the states in the Kaluza-Klein continuum to the matter stress-energy on the intersection. Just like in the Randall-Sundrum case \\cite{rs2,tunneling}, the couplings of the continuum modes are suppressed due to the warping of the bulk, except in this case as $\\sim m^2 L^2$, yielding them negligible at long range. Indeed, at large distances, $r \\gg L$ the leading order correction to the Newton's law is softened by additional powers of $L/r$ because of the tunnelling suppression, $\\delta V \\sim - G_N \\frac{m_1 m_2 L^5}{r^6}$. This shows explicitly that at large distances the objects localized to the 3-brane interact with the usual 4D Newtonian gravitational force. Similar situation persists for de Sitter origami, because the bulk volume is finite. We comment on the cosmic history of these models, indicating how the usual 4D FRW universes could be recovered. ",
        "conclusions": "In this paper we have derived exact solutions describing the intersection of two 4-branes in $AdS_6$ on a 3-brane of arbitrary non-negative tension. These solutions correspond to the vacua of the theory on the intersection, and generalize the analysis for the RS2 case of \\cite{bent}. In the case of the Minkowski and de Sitter origami, they localize 4D graviton on the intersection, yielding a fully consistent 4D effective theory in the low energy limit. We have explicitly computed the zero-mode wavefunction, and couplings of KK continuum, and evaluated the gravitational potential between two masses localized on the intersection. In the case of the $AdS$ origami there is no 4D localized graviton, but they may be an interesting arena to study the phenomenon of quasilocalization of \\cite{aklr}.  The Minkowski and de Sitter origami which localize 4D gravity to the intersection may be a natural background to formulate the low energy cosmology on codimension-2 braneworlds. We have elucidated the right limit where the homogeneous cosmological perturbations on the intersection gravitate like normal matter in 4D universe governed by the usual 4D Einstein gravity, as long as the order of scales $\\rho \\ll \\lambda \\ll \\sigma_k$ is maintained. We stress that the current precision of cosmological data allows for the inequalities to be satisfied by roughly two orders of magnitude, without spoiling the existing bounds on the validity of Einstein gravity. This makes the inequalities rather easy to satisfy, and opens up a possibility of developing models where such limits could be natural future attractors of cosmological evolution. In this context, one should also ask questions about the stabilization of the shape moduli, which are the angles between the 4-branes, which are fixed once the tensions are assigned. In the early universe, when the origami is perturbed, the shape moduli could roll, and it would be of interest to find out the details of their dynamics. We hope to return to these questions elsewhere.    \\vspace{.5cm} {\\bf Acknowledgements}  We thank N. Arkani-Hamed, S. Dimopoulos, G. Dvali, G. Gabadadze, M. Kaplinghat, M. Kleban, L. Randall and L. Sorbo for interesting discussions. This work was supported in part by the DOE Grant DE-FG03-91ER40674, in part by the NSF Grant PHY-0332258 and in part by a Research Innovation Award from the Research Corporation."
    },
    "0403/astro-ph0403645_arXiv.txt": {
        "abstract": "We present the results from a study of the host galaxies of 15 optically selected AGNs with $0.5<z<1.1$ from \\gems. \\gems\\ is a Hubble Space Telescope imaging survey of a $\\sim 28\\arcmin\\times 28\\arcmin$ contiguous field centered on the Chandra Deep Field South in the F606W and F850LP filter bands.  It incorporates the SEDs and redshifts of $\\sim$10~000 objects, obtained by the \\combo~ project.  We have detected the host galaxies of all 15 AGNs in the F850LP-band (and 13/15 in the F606W-band), recovering their fluxes, morphologies and structural parameters. We find that 80~\\% of the host galaxies have early-type (bulge-dominated) morphologies, while the rest have structures characteristic of late-type (disk-dominated) galaxies.  We find that 25~\\% of the early types, and 30~\\% of the late types, exhibit disturbances consistent with galaxy interactions.  The hosts show a wide range of colors, from those of red-sequence galaxies to blue colors consistent with ongoing star formation.  Roughly 70~\\% of the morphologically early-type hosts have rest-frame blue colors, a much larger fraction than those typical of non-active morphologically early-type galaxies in this redshift and luminosity range.  Yet, we find that the early-type hosts are structurally similar to red-sequence ellipticals, inasmuch as they follow an absolute magnitude versus half-light size correlation that are consistent with the mean relation for early-type galaxies at similar redshifts. ",
        "introduction": "\\label{sec:intro}   There is now a wide agreement that most if not all galaxies with bulges harbor massive black holes in their centers \\citep[e.g.,][]{mago98,korm01}. However, only a fraction of the galaxies harbor a quasar-like active nucleus (AGN). AGNs are often seen as compact sources in the nucleus of galaxies, with strong highly-ionized broad and narrow emission lines. They are thought to be powered by a central supermassive black hole, fed by the accretion of gas from the inner $\\sim$100 pc of the host galaxy \\citep{anto93,urry95}. While both a massive black hole and a reservoir of gas are believed to be important for triggering high level AGN activity, it is as yet unclear which conditions are needed to trigger such activity. Fundamental questions remain regarding the relationship between the AGN, its host galaxy, and star formation, the gas inflow rates needed to fuel the nuclear source, and the relevant fueling mechanisms on different physical scales.  The morphology and spectral properties of AGN host galaxies over different redshift regimes can bring important insights to these issues.   According to the currently favoured view, galaxy-galaxy interactions or mergers are important mechanisms for (re-)igniting AGN \\citep[e.g.,][]{cana01}, except perhaps for very low-level nuclear activity \\cite[e.g.,][]{coll03,ho03}. Mergers and interactions can induce large-scale radial motion of the interstellar gas within the AGN host, and trigger its infall into the inner regions to be accreted at later times by the AGN. \\cite{sand88} first proposed an evolutionary scenario in which galaxy encounters could produce ultraluminous infrared galaxies (ULIRGs), that would evolve into AGNs. This hypothesis is supported by several studies which found a large fraction of merger/interacting systems in the hosts of high-luminosity AGNs \\citep[e.g.][]{bahc95b, bahc97, hutc97, cana01, sanc03b}. Yet, \\cite{dunl03} have not found merger/interactions in their sample of low-$z$ AGNs. An important limitation of these previous studies is the lack of a well-defined comparison sample of inactive galaxies, whose data were obtained and analyzed in similar way to the active ones.  A severe practical problem for the study of the hosts of AGNs is the presence of the bright nucleus, which has to be carefully removed. This strong point-like source, as bright or even brighter than the galaxy, contaminates the host flux in all the areas covered by the point spread function (PSF).  It also reduces the SNR (signal-to-noise ratio) within the host. This affects strongly the reliability of any morphological classification of the host. Consecuently, ground-based studies are normally limited to low-redshift ($z<0.5$), where the typical size of the host is larger than the seeing disc \\citep[e.g.][]{dunl93, tayl96, mcle94a, mcle95a, jahn03e}, or to wavelength ranges where the host emission dominates the AGN (e.g., near-infrared), increasing the contrast \\citep[e.g.][]{dunl93, mcle94a, falo01, sanc03b}.  The use of HST imaging has clearly increased our understanding of the host properties, reducing the effects of contamination from the nucleus due to the very narrow intrinsic PSF \\citep[e.g.][]{mclu00, paga03}. However, most published HST host studies were based on WFPC2 imaging with short exposures due to the high AGN flux, with a limited determination/treatment of the PSF.  An additional limitation of previous studies has been the sample selection. Several samples have been selected {\\it ad hoc}, for testing specific hypotheses \\citep[e.g.,][]{dunl93, cana01}. Other samples are based on unclear selection criteria \\citep[e.g.,][]{hutc92, bahc95b} or are clearly incomplete and/or statistically small \\citep[e.g.,][]{hutc97}. In some cases the images were not deep enough for an accurate morphological analysis \\citep[e.g.,][]{lehn99}. A few studies are based on statistically significant, flux-limited samples \\citep[e.g.,][]{jahn02a, jahn03e}, although the results from these studies are limited by the lack of a similarly selected comparison sample of inactive galaxies.   All these difficulties can be better addressed by the {\\it Galaxy Evolution from Morphology and SEDs} project \\citep[\\gems,][]{rix03}, where AGNs have been identified by the \\combo\\ multi-color survey \\citep{wolf03a}.  We have selected the 15 AGNs with $0.5<z<1.1$ from \\gems, coincident with the best studied redshift range for the inactive galaxies in that survey. Using an accurate determination of the PSF, and a two-dimensional (2D) fitting technique, we decouple the host and nuclear components in these objects, and compare the properties of the host galaxies with those of inactive galaxies.  The sequence of this article is as follows. In Section 2 and 3 we describe the data and the data analysis. In Section 4 we present the results, and we discuss these in Section 5. The summary and conclusions are presented in Section 6. Throughout this article we assume a $\\Lambda$-cosmology, with $H_0=70~\\mathrm{km~s}^{-1}~\\mathrm{Mpc}^{-1}$, $\\Omega_{\\mathrm{m}}=0.7$ and $\\Omega_{\\Lambda}=0.3$. We have used the AB photometric system for the observed magnitudes and colors, and the Vega system for the rest-frame colors and absolute magnitudes, unless explicitly stated otherwise.     \\begin{table} \\caption{Properties of the \\gems\\  low-$z$ AGN Sample} \\label{tab1} \\begin{tabular}{cccccc} \\tableline\\tableline QSO ID&Tile&$z^1$&$R$-band&F606W-band&F850LP-band\\\\ \\tableline 34357&  94 &{\\it 0.543}&18.73&19.22$\\pm$0.06&18.40$\\pm$0.06\\\\ 41310&  47 &0.548&23.09&23.24$\\pm$0.09&22.48$\\pm$0.09\\\\ 41310&  40 &0.548&23.09&23.20$\\pm$0.09&22.44$\\pm$0.08\\\\ 52963&  50 &0.548&22.69&22.72$\\pm$0.09&22.10$\\pm$0.08\\\\ 36361&  94 &0.549&22.50&23.39$\\pm$0.09&22.17$\\pm$0.13\\\\ 47615&  46 &0.649&22.28&22.08$\\pm$0.08&21.36$\\pm$0.08\\\\ 50415&  85 &{\\it 0.664}&22.39&22.75$\\pm$0.09&21.70$\\pm$0.08\\\\ 50415&  51 &{\\it 0.664}&22.39&22.74$\\pm$0.09&21.66$\\pm$0.08\\\\ 44126&  42 &0.729&23.03&23.04$\\pm$0.09&22.33$\\pm$0.09\\\\ 42601&  42 &{\\it 0.733}&21.02&21.32$\\pm$0.08&20.24$\\pm$0.07\\\\ 48284&  85 &{\\it 0.734}&19.05&19.33$\\pm$0.06&18.92$\\pm$0.06\\\\ 39432&  89 &{\\it 0.738}&22.02&22.16$\\pm$0.08&21.43$\\pm$0.08\\\\ 31898&  30 &0.812&22.91&23.10$\\pm$0.09&21.62$\\pm$0.08\\\\ 15731&  18 &0.835&20.11&19.77$\\pm$0.07&19.43$\\pm$0.06\\\\ 50997&  85 &{\\it 0.837}&20.43&20.60$\\pm$0.07&20.38$\\pm$0.08\\\\ 50997&  51 &{\\it 0.837}&20.43&20.43$\\pm$0.07&20.22$\\pm$0.07\\\\ 49298&  52 &{\\it 1.031}&19.92&19.89$\\pm$0.07&19.47$\\pm$0.07\\\\ 49298&  85 &{\\it 1.031}&19.92&19.90$\\pm$0.07&19.56$\\pm$0.06\\\\ 43151&  95 &{\\it 1.037}&22.18&22.25$\\pm$0.08&21.78$\\pm$0.08\\\\ 43151&  45 &{\\it 1.037}&22.18&22.38$\\pm$0.09&21.92$\\pm$0.08\\\\ \\tableline \\end{tabular}  $^1$ values in italics are spectroscopic redshifts obtained from \\cite{szok04}.  \\end{table} ",
        "conclusions": "We have analyzed a flux-limited sample of 15 intermediate luminosity type 1 AGNs from \\combo, at $0.5<z<1.1$ using the \\gems\\ deep F606W and F850LP-band ACS images. We decomposed he host and nuclear components and detected the host in all the F850LP-band images (and in 13 of the F606W-band images). A morphological analysis has been performed to determine if the hosts were bulge or disk-dominated objects. $\\sim$80~\\% of the hosts are spheroidal galaxies, and only 1 object can be clearly classified as a disk dominated galaxy. We found evidence for on-going mergers in $\\sim$25~\\% of the early types and $\\sim$30~\\% of the late type hosts.  The host galaxies show a wide range of $\\V-\\Z$ colors, indicating the presence of different stellar populations. We found that a large fraction of the early-type hosts, $\\sim 50-70$\\%, show distinct blue colors when comparing with the red sequence inactive early-type galaxies studied by \\gems\\ at the same redshift range. The color and color-magnitude distributions of the AGN hosts are significantly different from those of red sequence galaxies. On the other hand, the absolute mangitudes, effective radii and stellar masses of these objects are remarkably similar to normal elliptical galaxies. Therefore, a significant fraction of the AGN in our sample seem to be located in \\emph{young spheroids}.  We considered two different simplified scenarios that can explain the observed color distributions: (a) a single stellar population with a mean age of close to 1~Gyr, or (b) a mix of an underlying old stellar population plus a small mass fraction in a young component.  Both can describe the observed colors well. On the other hand, the early-type host galaxies are structurally similar to well-assembled elliptical galaxies, following the luminosity-size relation of early-type galaxies. This indicates, most probably, that the process that has generated the blue colors has not strongly altered neither the luminosity ($\\sim$mass) and the effective radius of these galaxies. Taking into account this result, it seems that a mixed population, dominated by an underlying population of old stars, fits the observations better. This underlying population would reflect the population of the parent galaxy.  Our results are in excellent agreement with the results from the study of a large sample of type 1 and type 2 AGNs from the SDSS at lower redshift \\citep{kauf03}. They noted the existence of a population of \\emph{young bulge} galaxies ($\\sim$30\\% of their sample), in a state of post-starburst, and a population of even younger mergers ($\\sim$40~\\% of their sample). Recent results by \\cite{jahn02a} and \\cite{koti04} also found that low-$z$ host galaxies are bluer than field galaxies. While on the other hand, \\cite{dunl03} and \\cite{floy03} did not find evidence for distinctly blue colors in their sample of high-luminosity low-$z$ AGNs ($z<0.3$), their color range is less sensitive to young stellar populations at this redshift ($R-K$), especially if there is an underlying population of old stars.  A possible explanation to the connection between nuclear and star formation activity in apparently normal elliptical galaxies could be an enrichment of gas produced by a merging process or the infall and concentration of gas produced by a minor interaction. Since the fraction of merger galaxies in our sample seems not to be remarkably larger than in inactive galaxies, the merger/interaction should have happened in a timescale long enough to enable the galaxy to relax morphologically before the current AGN activity."
    },
    "0403/astro-ph0403159_arXiv.txt": {
        "abstract": "We re-analyze all archival {\\it Chandra/ACIS} observations of the Orion Nebula Cluster (ONC) to study the X-ray properties of a large sample of pre--main-sequence (PMS) stars with optically determined rotation periods. Our goal is to elucidate the origins of X-rays in PMS stars by seeking out connections between the X-rays and the mechanisms most likely driving their production---rotation and accretion. Stars in our sample have $L_X / L_{bol}$ near, but below, the ``saturation\" value of $10^{-3}$.  In addition, in this sample X-ray luminosity is significantly correlated with stellar rotation, in the sense of {\\it decreasing} $L_X / L_{bol}$ with more rapid rotation. These findings suggest that stars with optical rotation periods are in the ``super-saturated\" regime of the rotation-activity relationship, consistent with their Rossby numbers. However, we also find that stars with optical rotation periods are significantly biased to high $L_X$. This is not the result of magnitude bias in the optical rotation-period sample but rather to the diminishingly small amplitude of optical variations in stars with low $L_X$. Evidently, there exists in the ONC a population of stars whose rotation periods are unknown and that possess lower average X-ray luminosities than those of stars with known rotation periods. These stars may sample the linear regime of the rotation-activity relationship. Accretion also manifests itself in X-rays, though in a somewhat counterintuitive fashion: While stars with spectroscopic signatures of accretion show {\\it harder} X-ray spectra than non-accretors, they show {\\it lower} X-ray luminosities and no enhancement of X-ray variability. We interpret these findings in terms of a common origin for the X-ray emission observed from both accreting and non-accreting stars, with the X-rays from accreting stars simply being attenuated by magnetospheric accretion columns. This suggests that X-rays from PMS stars have their origins primarily in chromospheres, not accretion. ",
        "introduction": "} X-rays serve as one of our primary probes of magnetic activity in solar-type stars. Thus, much of our understanding of key physical processes thought to be connected to magnetic fields---stellar winds, for example---derives from the study of stellar X-rays.  One of the most compelling stories in stellar astrophysics told through X-rays is that of the intimate relationship between stellar rotation and magnetic field generation. Indeed, among late-type main-sequence stars, rotation is the strongest correlate of X-ray luminosity, and the observed rotation/X-ray relationship \\citep{pallavicini81,caillault96,randich97,jeffries99,randich00} has become central to the current paradigm of dynamo-generated magnetic fields, of magnetically driven stellar winds, and ultimately of the evolution of stellar angular momentum.  The observed relationship between rotation and X-ray emission on the main sequence is remarkably clean, and clearly separates stars into three regimes (cf.\\ \\citet{randich00}) typically described phenomenologically as the linear, saturated, and super-saturated regimes, in order of increasing stellar rotation \\citep{randich97}. For slowly rotating stars the X-ray luminosity, $\\log L_X$, scales linearly with the stellar angular velocity, $\\log\\Omega$, consistent with the theoretical idea that more rapid stellar rotation produces a stronger magnetic field through an $\\alpha - \\Omega$-type dynamo (for stars with radiative cores and convective envelopes) or through a distributed turbulent dynamo (for fully convective stars). For stars rotating more rapidly than a certain threshold, the X-ray luminosity is observed to ``saturate\" at a fixed value relative to the stellar bolometric luminosity: $\\log L_X / L_{bol} \\approx -3$. While the reasons for saturation have not been well understood, this observation has had important ramifications for efforts to model the angular momentum evolution of young solar-type stars \\citep{krishnamurthi97,bouvier97}. The models now routinely include saturation as a key ingredient in their parametrizations of angular momentum evolution through winds. Finally, the most rapidly rotating stars exhibit X-ray emission at levels roughly a factor of 2 below the saturation value. This ``super-saturation\" effect \\citep{james00} has been poorly understood. \\citet{barnes03a,barnes03b} has recently re-interpreted super-saturation in terms of a new paradigm for the angular momentum evolution of solar-type stars; whether this new paradigm will survive detailed scrutiny is not yet clear.  What is clear is that the rotation/X-ray relationship serves as a key observational touchstone for developing and evaluating our theoretical understanding of the generation and evolution of stellar magnetic fields, the generation and evolution of stellar winds, and the evolution of stellar angular momentum. Considerable observational effort has been invested, therefore, in trying to establish the presence of a rotation/X-ray relationship among pre--main-sequence (PMS) stars, where the questions of magnetic field generation and evolution, winds, and angular momentum evolution remain largely unanswered. Unfortunately, to date these efforts have not borne much fruit. In one recent study of T~Tauri stars (TTS) in Taurus-Auriga \\citep{stelzer01}, a rotation/X-ray correlation has been reported, but the sample size is small ($N=39$) and there are lingering concerns with respect to completeness/reliability of the rotation periods and biases in the sample both astrophysical and observational. In particular, as discussed by \\citet{feig03}, the Taurus-Auriga PMS population appears to be deficient in both high-mass stars and faint, low-mass weak-lined TTS. In addition, rotation periods in this region were determined from photometric monitoring campaigns some 10 years ago that were relatively sparsely sampled and of relatively short duration, and therefore potentially biased against both very fast and slow rotators.  Two new studies based on deep {\\it Chandra} observations of the Orion Nebula Cluster (ONC) provide the most comprehensive analyses yet of X-rays and rotation in a large, coeval ($\\sim 1$ Myr) sample of PMS stars. Neither the study of Flaccomio and collaborators \\citep{flac-basic,flac-acc,flac-time} with the High Resolution Camera (HRC), nor that of Feigelson and collaborators \\citep{feig02,feig03} with the Advanced CCD Imaging Spectrometer (ACIS), found evidence for a rotation/X-ray relationship such as that observed on the main sequence or that reported for Tau-Aur TTS by \\citet{stelzer01}. Indeed, these studies find that it is stellar mass that is by far the dominant correlate of PMS X-ray luminosity, with $\\log L_X/L_{bol}$ correlating with rotation either not at all, or perhaps slightly in the opposite sense from the main-sequence rotation/X-ray relationship.  Both studies found that stars with known rotation periods have X-ray luminosities near the main-sequence saturation value of $\\log L_X / L_{bol} \\approx -3$. Moreover, both studies suggest that PMS stars might in fact be {\\it expected} to reside in the super-saturated regime, considering that typical PMS Rossby numbers \\citep{kim-demarque,ventura} are small due to the long convective turnover times ($\\tau_c \\sim 800$ days) of these very young and fully convective stars. Thus, while direct observation of PMS stars in the linear regime of the rotation/X-ray relationship remains elusive, these studies seem to confirm, if indirectly, the basic picture of the rotation/X-ray relationship by suggesting that all ONC stars are in the super-saturated regime, and that this is where they ought to be.  But not all ONC stars detected by {\\it Chandra} have known rotation periods. To what extent is the sample with periods representative of the entire ONC population in terms of X-ray properties? If they are not representative, how does this group differ in other salient characteristics, such as accretion, and how might these differences affect our interpretation of the origin of X-rays in PMS stars? In addition, while it is clear from previous analyses that X-ray emission from PMS stars is not temporally static---X-ray flaring is ubiquitous in the ONC---it is not yet clear whether or to what extent X-ray variability may be affecting our ability to measure reliable X-ray luminosities. Might X-ray flaring be scrambling the signal of an underlying rotation/X-ray relationship? That some stars flare while others do not is interesting in its own right: Do the flaring characteristics of stars with known rotation periods represent those of all stars?  Again, how might differences here affect our interpretation of the origin of PMS X-rays?  Motivated by these questions, we have re-analyzed all archival {\\it Chandra}/ACIS observations of the ONC that include stars with known rotation periods. Our aim is to derive X-ray luminosities for as large a sample of known rotators as possible, employing a consistent analysis scheme throughout, including filtering of flares in the hopes of minimizing the effects of X-ray variability.  In \\S 2 we describe the data used and our processing/analysis procedures. We then present our basic results in \\S 3, first focusing on the X-ray nature of the rotator sample as compared to the entire ONC population. We show that stars with known rotation periods are significantly more X-ray luminous, and more likely to be X-ray variable, than stars for which rotation periods have not been measured. We then explore the relationship between X-rays and rotation. We find that most stars with known rotation periods appear to be in the super-saturated regime, having $\\log L_X / L_{bol} \\lesssim -3$, with a statistically significant correlation in which faster rotators have lower X-ray luminosities. But we also find that stars without rotation periods---being less X-ray luminous on average---show a range of $L_X/L_{bol}$ comparable to that observed on the main sequence. These stars may represent the beginnings of the linear regime of the rotation/X-ray relationship. Finally, we explore the relationship between X-rays and accretion. We find that while stars with spectroscopic signatures of accretion show harder X-ray spectra than non-accretors, they also show lower X-ray luminosities and no enhancement of X-ray flaring.  We discuss the implications of our findings in \\S 4, where we (a) emphasize that current rotation-period measurements in the ONC have not probed the full range of underlying stellar X-ray properties, (b) suggest that a main-sequence type relationship between X-rays and rotation may in fact be present in the ONC, and (c) argue that the data imply a chromospheric---not accretion---origin for X-rays from PMS stars.  We summarize our conclusions in \\S 5. ",
        "conclusions": "} We have re-analyzed all archival {\\it Chandra/ACIS} observations of pre--main-sequence (PMS) stars with optically determined rotation periods in the Orion Nebula Cluster (ONC). Our aim is to investigate the relationship between X-rays and the physical mechanisms most likely related to their production in PMS stars: rotation and accretion. Our analysis procedures include filtering of flare events in the X-ray data in an attempt to determine X-ray luminosities that are free of the stochasticity introduced by such events.  The primary findings of this study are as follows: \\begin{enumerate} \\item Stars with optically determined rotation periods are more X-ray luminous, and are more likely to be X-ray variable, than are stars without optical rotation periods. We show that the bias to high $L_X$ is not due to a magnitude bias in optical rotation-period studies of the ONC; rather, it is due to the diminishingly small amplitude of optical variability among stars with smaller $L_X$, precluding detection of their rotation periods.  \\item Stars with optically determined rotation periods have a mean $L_X/L_{bol}$ near, but lower than, the ``saturation\" value of $\\sim 10^{-3}$, implying that these stars are in the saturated or super-saturated regimes of the X-ray/rotation relationship, consistent with their Rossby numbers. There is a marginally significant ($\\sim 3\\sigma$) correlation between $L_X/L_{bol}$ and $P_{\\rm rot}$, with the more rapidly rotating stars showing lower $L_X/L_{bol}$, as is seen among super-saturated stars on the main sequence.  \\item Compared to these stars, stars {\\it without} rotation periods show a larger range of $L_X/L_{bol}$---comparable, in fact, to that found among main sequence stars. We consider the possibility that, among these, some stars may lie at the beginnings of the ``linear\" regime of the X-ray/rotation relationship. Using $v\\sin i$ data from the literature we find that, among these stars lacking known rotation periods, slower rotators do indeed show lower X-ray luminosities than do rapid rotators. This relationship is not one-to-one, however. It is also possible that the lower $L_X$ among stars lacking rotation periods is instead due to the higher incidence of active accretion among these stars, a possibility for which we also find weak evidence. The statistical significance of these two effects---$v\\sin i$ and accretion---are comparable. Measurements of $v\\sin i$ sensitive to very slow rotators ($\\lesssim 5$ km/s) would be of great value in furthering our understanding of X-ray production at the slow extremes of PMS rotation. PMS stars in the linear regime should have $P_{\\rm rot} \\gtrsim 100$ days, assuming a typical convective turnover timescale of $\\tau_c \\sim 800$ days. Such long rotation periods have yet to be observed among PMS stars.  \\item Stars in the ONC with spectroscopic signatures of active accretion show significantly harder X-ray spectra and lower X-ray luminosities than their non-accreting counterparts. These observations can be explained quantitatively by a model in which accretors and non-accretors have intrinsically similar X-ray emission properties, with the differences in $L_X$ and hardness ratio being due to absorption of soft X-rays by magnetospheric accretion columns. \\end{enumerate}  Taken together, these findings hint that there in fact exists a rotation-activity relationship among PMS stars in the ONC, and suggest that rotation---not accretion---is the primary driver of X-ray emission in low-mass ($M \\lesssim 3$ M$_\\odot$) PMS stars at 1 Myr. Indeed, our finding that stars with rotation periods show elevated levels of X-ray variability, despite showing little stochastic variability in the optical, further implies that X-ray variability has its origins in processes that are more or less independent of the processes responsible for stochastic variability in the optical (i.e.\\ accretion).  Finally, our findings raise questions about the true physical meaning of ``saturation\" in PMS stars. It is intriguing that stars with optically determined rotation periods all appear to lie in the super-saturated regime yet show diminishingly small amplitudes of optical variability at low $L_X$.  It is possible that spots on the surfaces of these stars become non-existent below a certain $L_X$ threshold. On the other hand, we speculate that the low amplitude of optical variability may be due to magnetic topologies in which the stellar surface is indeed largely covered by spots, but spots that are more-or-less randomly distributed over the stellar surface, thereby producing only very small photometric signals in the optical. More organized magnetic topologies may be present in stars with higher $L_X$, such that larger spots asymmetrically distributed on the stellar surface are possible. In this picture, these latter stars might be those whose global fields have become sufficiently organized and deeply rooted so as to begin effecting magnetic braking of the stellar rotation, a picture similar to that recently put forward by \\citet{barnes03a,barnes03b}."
    },
    "0403/astro-ph0403190_arXiv.txt": {
        "abstract": "We present new multi-configuration \\HI\\ aperture synthesis imaging of the nearby dwarf starburst galaxy NGC\\,625 obtained with the Australia Telescope Compact Array.  Total \\HI\\ column density images show gas well-aligned with the optical major axis, and low-column density \\HI\\ extending to $>$ 6 optical scale lengths.  The \\HI\\ velocity field, on the other hand, is highly disturbed, with neutral gas at nearly all detected velocities within the central region.  After considering various interpretations, we find that a blowout scenario most accurately describes the data. Since at our resolution we do not detect any large evacuated holes in the \\HI\\ disk, we interpret this blowout to be the result of the extended (both spatially and temporally) star formation event which NGC\\,625 has undergone in the last 100 Myr.  This is one of the clearest examples of \\HI\\ outflow detected in a dwarf galaxy.  We find no obvious external trigger for this extended star formation event.  We detect strong radio continuum emission from the largest \\HII\\ regions; comparing to our HST and ground-based \\halpha\\ fluxes suggests either appreciable amounts of extinction toward the star formation regions, or the contribution of non-thermal sources to the radio continuum luminosity. ",
        "introduction": "\\label{S1}  Nearby dwarf galaxies offer the opportunity to study the dynamics of galaxy evolution at relatively high spatial resolution.  In these systems we can probe the nature of star formation and observe its direct impacts on the ISM.  Furthermore, the smaller scales of dwarf systems allow us to address the global impacts of starbursts on the ISM, and in relatively high-luminosity cases, the effects on the surrounding IGM as well.  In these systems we can attempt to ascertain the triggering mechanisms of starbursts in low-mass galaxies, which will be an important component of models of the chemical and dynamical evolution of galaxies.  As a result of the relative paucity of strong starbursting dwarf galaxies in the local volume, there are few detailed studies of the neutral gas kinematics of such systems.  \\citet{kobulnicky95} found the surprising result that in the well-studied dwarf starburst NGC\\,5253, the large-scale rotation of the galaxy appears to align along the galaxy's optical major axis.  This rare kinematic alignment may be the signature of recent tidal interaction and may provide evidence for a triggered starburst in this system.  A more clear case of a triggered starburst is found in NGC\\,1569.  \\citet{stil02c} find that a low-mass companion seems to be interacting with the system, and that large amounts of the total \\HI\\ mass are present at unusually high velocities with respect to systemic.  These data suggest that infall/outflow and/or tidal effects may be important for this system, and that the current starburst may present a case of triggered star formation. This burst is having a dramatic effect on the ISM, as turbulent motions appear to support the inner disk, and the mean velocity dispersion is high ($\\sigma_V \\sim$ 20\\,\\kms) compared to other starbursts.  Not all starbursts in low-mass galaxies appear to be tidally induced, however. In NGC\\,1705, well-known because of its spectacular galactic wind, the global \\HI\\ kinematics appear to be quite regular {(Meurer, Staveley-Smith, \\& Killeen 1998)}\\nocite{meurer98}.  While high-velocity \\HI\\ gas does exist, the kinematics suggest simple rotation and a relatively quiescent velocity field. These few local blue compact dwarfs (BCD's) support the conclusions of \\citet{taylor97}, wherein it was found that $\\sim$\\,60\\%\\ of such systems harbor nearby \\HI\\ companions that may have triggered the current starbursts.  However, the local sample also suggests that there must exist another mechanism with which to initiate powerful starbursts in low-mass galaxies.  In this investigation we will examine the neutral gas properties of the nearby dwarf starburst galaxy NGC\\,625, with one of the primary goals being to discern the nature of the triggering mechanism of the current starburst.  The ejection of material from galaxies (i.e., ``blowout'') is a topic of great importance for their overall evolution.  In systems of decreasing total mass, temporally and spatially concentrated massive star formation becomes increasingly important.  For the lowest mass galaxies, ``starbursts'' can dominate the energetics and luminosity of the system \\citep{dekel86,deyoung94}.  Such events can also affect their immediate surroundings by injecting energy into the ISM.  If the energies are high enough, they can eventually break out of the disk and into the halo, releasing hot gas and processed material which may or may not escape the system entirely to join the surrounding IGM. Numerical simulations of mass loss in low-mass systems \\citep[e.g.,][]{maclow99} suggest that it is quite difficult to remove substantial fractions (i.e., \\gsim\\ 10\\%) of the ISM in galaxies with masses $\\gsim$ 10$^7$ \\msun.  Thus it is useful to examine the kinematics of gas in low-mass systems where this model can be empirically tested.  As will be discussed below, these observations add NGC\\,625 to the short list of dwarf systems where neutral gas outflow has been detected.  NGC\\,625 is an intriguing starburst system located on the far side of the nearby Sculptor Group (D = 3.89$\\pm$ 0.22 Mpc; {Cannon \\etal\\ 2003}\\nocite{cannon03}).  In that investigation, HST/WFPC2 photometry was used to construct dereddened V vs. (V$-$I) color-magnitude diagrams.  It was found that NGC\\,625 has been undergoing a prolonged ($\\sim$ 100 Myr) starburst episode.  The spatially resolved nature of those data also allowed us to show that the recent star formation has been widespread throughout the central regions. Comparing the locations of young stars with the smooth distribution of older red giant stars, it is likely that star formation has been active for an even longer period.  This is one of only a few cases where an extended burst of star formation has been observed in a nearby dwarf galaxy.  The contributing factors to such an evolutionary scenario are not yet fully understood, but the neutral gas dynamics offer us a second avenue with which to investigate the recent evolution of this system.  Previous \\HI\\ observations of NGC\\,625 were presented by {C{\\^ o}t{\\' e}, Carignan, \\& Freeman (2000)}\\nocite{cote00}.  They identified multiple emission peaks in the central regions of the galaxy and rotation about the galaxy's optical major axis (as opposed to the typical rotation about the minor axis).  The complex kinematics precluded the derivation of a mass model or a rotation curve for this system and prompted the present investigation. As the cause for the complex neutral gas kinematics, \\citet{cote00} suggest a merger scenario - perhaps with a large \\HI\\ cloud \\citep[which are numerous in this region of the sky, although mostly at differing velocities; see][]{putman02}.  In a study of the ionized gas kinematics, however, \\citet{marlowe97} also find complex orbits and motions, suggestive of a major merger as the cause of the burst.  The possibility does exist, interestingly, that the kinematics of the ionized and neutral gas are caused by the current burst itself; i.e., perhaps the disk of this system is being disrupted by the strong and extended star formation episode.  With these \\HI\\ data, we seek to further constrain the nature of the recent activity in this system.  Other investigations also point toward the importance of NGC\\,625 as a nearby example of the dwarf starburst phenomenon.  The spectroscopic investigation of {Skillman, C{\\^ o}t{\\' e}, \\& Miller (2003b)}\\nocite{skillman03b} finds strong, broad (5.5\\,\\AA\\ EW) $\\lambda$\\,4686 \\AA\\ emission arising from the largest \\HII\\ region, suggesting a Wolf$-$Rayet (W$-$R) classification for the current burst \\citep{conti91} and an age of 4$-$6 Myr.  This result is somewhat surprising, given the extended nature of the burst derived in \\citet{cannon03}, and suggests that the presence of W$-$R features does not necessarily imply a young ($<$ 6 Myr) burst. {Skillman, C{\\^ o}t{\\' e}, \\& Miller (2003a)}\\nocite{skillman03a} find a current star formation rate of 0.05 \\msun\\,yr$^{-1}$, comparable to the values found in other high-luminosity dwarf starburst systems.  \\citet{bomans98} find diffuse 0.1$-$2.4 keV x-ray emission coincident with the large \\HII\\ region and extended above the northern side of the plane, suggesting that a galactic wind has been active or has recently terminated.  The above properties imply that the recent star formation has been intense in NGC\\,625; the evolution of this system has been complex and offers an ideal opportunity to study many important questions about the starburst phase of dwarf galaxy evolution. We therefore have begun a multiwavelength study of this system, of which these \\HI\\ observations form an important part.  For reference, in Table~\\ref{t1} we summarize basic observational properties of this system.  We organize this paper as follows.  In \\S~\\ref{S2} we discuss the data obtained and our reductions.  The \\HI\\ data and its analysis, our interpretations, and a simple model are discussed in \\S~\\ref{S3}.  In \\S~\\ref{S4} we discuss the strong radio continuum emission detected from the three highest-surface brightness \\HII\\ regions, and \\S~\\ref{S5} summarizes our conclusions.  \\placetable{t1} ",
        "conclusions": "\\label{S5}  We have presented new ATCA multi-configuration \\HI\\ synthesis imaging of the nearby dwarf starburst galaxy NGC\\,625.  The \\HI\\ morphology appears smooth on global scales; low-resolution column density images show one strong central \\HI\\ peak surrounded by a moderate-size low-density halo. However, examining the morphology at higher resolution reveals multiple peaks in the column density, suggesting that the dynamics of the neutral gas are quite complex.  Examination of the velocity field shows a complicated structure, with \\HI\\ emission detected near the disk across the full velocity range. We find velocity structure in the \\HI\\ data cubes that appears consistent with the signature of neutral gas blowout from the main disk.  While at our resolution we are not sensitive to small-scale structure (which may show holes and other signatures of the influence of massive stars), we are sensitive to the global effects that this star formation is having on the ISM of NGC\\,625. Because we do not detect large \\HI\\ holes, and given the results of our HST/WFPC2 star formation history analysis \\citep{cannon03}, we prefer the interpretation that the current outflow episode is not the result of one major star formation region, but rather the collective influence of the extended (both temporally and spatially) star formation event which NGC\\,625 has undergone.  Active blowout is the only interpretation that is capable of reproducing all of the characteristics of the position-velocity diagrams along the major axis of the system.  Other potential explanations of the complex velocity structure, including superposed disks, counter-rotating disks, infall, ongoing assimilation, high-velocity clouds, and ISM turbulence, are discussed and rejected, each failing to reproduce all of the observed velocity structure.  While a detailed hydrodynamical simulation is beyond the scope of this paper, we encourage further such investigations to better understand the mechanisms involved in creating blowout in low-mass systems.  At the current sensitivity, we find no obvious external trigger for the current burst of star formation.  As NGC\\,625 is a comparatively isolated galaxy, we thus conclude that it has been able to sustain a heightened star formation rate due to the relatively large reservoir of high-column density gas in the central disk regions.  The initial triggering mechanism for this extended burst remains unknown.  We detect strong radio continuum emission from the three largest star formation regions in NGC\\,625.  Comparing to our HST/WFPC2 \\halpha\\ fluxes, we find that either appreciable amounts of extinction exist toward these regions, or that there is a strong nonthermal component to the radio continuum emission. We have undertaken multi-frequency continuum observations of this system to further differentiate between these two scenarios.  NGC\\,625 is a remarkable system.  It appears to be one of only a few dwarf galaxies that have shown neutral gas blowout from the active star formation regions.  It is the only dwarf galaxy known to show the kinematic signature of blowout, but not to show a large evacuated region near the active star formation region (although such a hole(s) may have gone undetected in the present experiment, at our resolution).  Given the detection of diffuse soft x-ray gas above the plane of this system \\citep{bomans98}, it appears that the extended massive star formation event in NGC\\,625 \\citep{cannon03} has achieved a similar integrated effect.  This has important implications for the evolution of dwarf galaxies.  Given the correct conditions, they are able to sustain extended star formation events, which will have dramatic effects on the ISM and surrounding IGM."
    },
    "0403/astro-ph0403473_arXiv.txt": {
        "abstract": "The discovery of quasi-periodic oscillations (QPOs) in low-mass X-ray binaries (LMXBs) has been reported and discussed in recent studies in theoretical and observational astrophysics. The {\\it Rossi X-Ray Timing Explorer} has observed oscillations in the X-ray flux of about 20 accreting neutron stars. These oscillations are very strong and remarkably coherent. Their frequencies range from $\\sim 10$ Hz to $\\sim 1200$ Hz and correspond to dynamical time scales at radii of a few tens of kilometers, and are possibly closely related to the Keplerian orbital frequency of matter at the inner disk. Almost all sources have also shown twin peaks QPOs in the kHz part of the X-ray spectrum, and both peaks change frequencies together.  Bearing in mind these facts, several models explain the observed frequencies by the oscillation modes of an accretion disk. In the beat-frequency interpretation, the upper kHz QPOs are associated with the Keplerian frequency of the accreting gas flowing around a weakly magnetized neutron star ($B\\sim 10^7-10^{10}$ G) at some preferred orbital radius around the star, while the lower kHz QPOs result from the beat between the upper kHz QPO and the spin frequency of the neutron star. In the relativistic precession model, while the upper kHz QPO is still the orbital frequency of matter at some radius in the disk, the lower kHz QPOs are considered as general relativistic precession modes of a free particle orbit at that radius.  In this study we propose a new generic model for QPOs based on oscillation modes of neutron star magnetospheres. We argue that the interaction of the accretion disk with the magnetosphere can excite resonant shear Alfv\\'en waves in a region of enhanced density gradients. We demonstrate that depending on the distance of this enhanced density region from the star and the magnetic field strength, the frequency of the field line resonance can range from several Hz (weaker field, farther from star), to approximately kHz frequencies (stronger field, $\\sim 6-10$ star radii from the star).  We show that such oscillations are able to significantly modulate inflow of matter from the high density region toward the star surface, and possibly produce the observed X-ray spectrum. In addition, we show that the observed $2:3$ frequency ratio of QPOs is a natural result of our model. ",
        "introduction": "Recent years have seen a surge in interest in studies of accreting neutron stars and black holes.  One reason is that these objects provide a very unique window to a deeper understanding of the physics of strong gravity and dense matter. The complicated structure of these stellar objects, including strong gravity, gravitationally-captured plasma accretion, rapid rotation, bursting in accreted materials, thermal X-ray emission, and several other features makes a most fascinating object for both theoretical and observational investigations. One of the most interesting features of these objects is the interaction of accreting plasma with the stellar magnetic field. Matter is transferred from a normal donor star to the compact object. MHD models of the infalling plasma and the neutron star magnetosphere allow a relatively simple first approach to an understanding of plasma flow and the configuration of the magnetosphere around a neutron star.  Research on plasma accretion by magnetic neutron stars began in the early 1970s with the discovery of periodic oscillations in X-ray fluxes of the stars Cen X-3 \\citep{Gia71,Sch72} and Her X-1 \\citep{Tan72}.   Interest in these systems rose about two decades ago, after the first observation of quasi-periodic oscillations (QPOs) in the persistent X-ray fluxes of the luminous X-ray sources \\citep{Van85,Has86,MP86}.  These oscillations are detected both in low and high frequencies with frequencies ranging from $\\sim 1$ to $\\sim 100$ Hz and from $\\sim 300 $ to $\\sim 1200$ Hz, respectively.  The {\\it Rossi X-Ray Timing Explorer} (RXTE) has observed persistent X-ray emissions about 20 accreting neutron stars in low-mass X-ray binaries (LMXBs). Generally, three different millisecond phenomena have now been observed in X-ray binaries: the twin kilohertz quasi-periodic oscillations (kHz QPOs), the burst oscillations, and the spin frequency of the accreting, low-magnetic field, neutron star.  These kilohertz frequencies are the highest frequency oscillations ever seen in any astrophysical object, and are now widely interpreted as due to orbital motion in the inner accretion flow.  The oscillations sometimes have one spectral peak.  In most sources, however, the X-ray spectrum shows two simultaneous kHz peaks with a roughly constant peak separation, moving up and down together in time as a function of photon count rate.  In this work, we focus on the Hz-kHz QPOs.  The burst oscillations with frequencies in the $300-600$ Hz range have been observed in a number of the X-rays sources, usually during a type I X-ray burst due to a thermonuclear runaway in the accreted matter on the neutron star surface.    Furthermore, the frequency of the burst oscillation increases by $1-2$ Hz during the burst tail, converging in time to a stable frequency at the neutron star spin frequency.   It is believed that the burst oscillations arise due to a hot spot or spots in an atmospheric layer of the neutron star. This has been explained by a decoupling of the atmospheric layer from the star and expansion by $5-50$\\,m during the X-ray burst, while its angular momentum remains constant. Therefore, during the burst tail, the frequency increases due to spin-up of the atmosphere as it re-contracts. The asymptotic frequency corresponds to a fully re-contracted atmosphere and is the closest to the true neutron star spin frequency (see the recent review by \\cite{BS03}).  In a few sources, however, a rather large frequency drift, $\\sim 5-10$ Hz, is observed. This drift is not consistent with the above model ( see \\cite{RS03} for more details).  Although there are some detailed differences in different types of X-ray sources, the observed QPOs in all sources are remarkably similar, both in frequency and peak separation.  (Six of the 20 known sources are originally identified as ``Z'' sources and the rest are known as ``atoll'' sources, see table 1. For more information about the atoll and Z sources, see \\citet{Van00}.) Such a similarity shows that QPOs should depend on general characteristics of the X-ray sources which are common to all systems.  In other words, the QPO can be regarded as a generic feature of the accreting neutron star.  In the accretion disk of accreting binaries the material from the donor star moves around the compact star in near-Keplerian orbits. Depending on the binary separation, the disk radius varies from $10^5$ to $10^7$ km. Although there is still no certainty about the flow geometry in the inner emitting region, the extension of a part of the flow down into the emitting region in the form of a Keplerian disk is the basic assumption of most proposed models for accretion of matter onto low magnetic field neutron stars (see for example \\cite{Mil98a}). Due to the interaction with the magnetic field of the star, radiation drag, and relativistic effects, the flow will be terminated at a radius $r_{\\rm in}$, the innermost stable orbital radius that is larger than radius of the compact star $R_s$. Within $r_{in}$ matter may leave the disk and so the flow is no longer Keplerian. Remarkably, the characteristic dynamical time scale for materials moving near the compact object is comparable with the observed millisecond X-ray variability, ie. $\\tau_{\\rm dyn}=(r^3/G M)^{1/2}\\sim 2~ {\\rm ms}~ (r/100 {\\rm km})^{3/2} (1.4 M_\\odot/M)^{1/2}$.   Such a natural periodicity is the foundation of most models for the observed QPOs. We will discuss these models in the next section.  Another common feature in accreting binaries is the interaction of the accretion disk with neutron star's magnetosphere.   In all sources, matters transfer from a normal donor star to the compact object. This matter is accelerated by the gravitational pull of the compact object and hits the magnetosphere of the star with a sonic/supersonic speed. The MHD interaction of the infalling plasma with the neutron star magnetosphere, will alter not only the plasma flow toward the surface of the star, as assumed by current QPO models, but also the structure of the star's magnetosphere.  The magnetic field of the neutron star is distorted inward by the infalling plasma of the Keplerian accretion flow. Since the gravity of the star confines the inward flow to a small solid angle $\\sim 10^{-2}$ ster, the magnetic field of the star will be more compressed in the disk plane than in other areas, see Fig. \\ref{star-fig1}. Furthermore, in a more realistic picture, one would expect that the highly accelerated plasma due to the infall process would be able to penetrate through the magnetic field lines. Such materials will be trapped by the magnetic field lines and produce enhanced density regions within the magnetosphere. However, due to their initial velocities, they move along the field lines, finally hit the star surface at the magnetic poles and produce the observed X-ray fluxes (Fig. \\ref{star-fig1}). See \\cite{GLP77} for more detail.  Besides modifying the geometry of magnetosphere, the compressional action of the accretion flow can excite some perturbations in the enhanced density region.  This can be understood by noting this fact that the inward motion of the accretion flow will be halted by the outward magnetic pressure at a certain distance from the star, the Alfv\\'en radius \\citep{GLP77}. In other words, the accretion flow pushes the stellar magnetic field toward the star until the inward pressure of the infalling plasma $\\rho v_r^2/2$ balances with the outward magnetic pressure $B_p^2/8\\pi$.  Here $\\rho$ and $v_r$ are the density and radial velocity of the infalling matter and $B_p$ is the poloidal magnetic field at the disk plane. Therefore, any instability in the disk at the Alfv\\'en radius would disturb such quasi-equilibrium configuration as well as the structure of magnetosphere.  As an example, the interaction of the solar wind with earth magnetosphere excites resonant shear Alfv\\'en waves, or field line resonances (FLRs), along the magnetic field lines \\citep{Sam91}.  As a result, one might expect that such Alfv\\'en waves can be excited by the compressive accreting plasma in the magnetosphere of an accreting neutron star with the accretion flow playing the role of the solar wind.  The plan of this paper is as follow: In section \\ref{sect:qpo} we discuss some of the current models for the observed QPOs. Specifically, we will review the two models, the sonic-point beat frequency model and the relativistic precession model, in more detail. We outline an occurrence of the resonant coupling in a magnetized plasma in section \\ref{sect:mag}.  Specifically, we review the excitation of shear Alfv\\'en waves by studying linear perturbations in MHD.  Next, the occurrence of field line resonances (FLRs), as a result of a resonant coupling between the compressional and shear Alfv\\'en waves, is discussed.  To illustrate the basic features of FLRs, the excitation of these resonances in a rectilinear magnetic field and in a dipolar field are considered. A dipolar topology is likely appropriate for the magnetosphere of a neutron star, near the surface.   Depending on the field line where the resonance occurs the eigenfrequency of the FLR is in range of several hundred Hz to kHz. In section \\ref{amb}, in order to consider a more realistic model for an accreting neutron star, we study an influence of the ambient flow along the field line on the excitation of FLRs. We demonstrate that in this case the eigenfrequency of the Alfv\\'en modes is modulated by the velocity of the plasma flow. Furthermore, in the presence of the field aligned flow, the plasma displacement parallel to the magnetic field lines is non-zero. These displacements, which are missing in case of zero ambient flow, might be responsible for modulating the motion of infalling materials toward the magnetic poles and producing the observed X-ray fluxes. A possible occurrence of more than one peak in the power spectrum is also discussed. In addition, we show that the observed $2:3$ frequency ratio of QPOs is a natural result of our model. Section \\ref{conc} is devoted to summarizing our results and further discussions. ",
        "conclusions": "\\label{conc}  In the present work, we have studied the interaction of an accretion disk with a neutron star magnetosphere in the LMXBs. The recent extensive observations reveal the existence of quasi-periodic oscillations in the X-ray fluxes of such stars. These oscillations, with frequencies range from $10$ Hz to $1200$ Hz, have been the subject of several theoretical and observational investigations. Based on theoretical models for the observed aurora in the earth magnetosphere, we have introduced a generic magnetospheric model for accretion disk-neutron star systems to address the occurrence and the behavior of the observed QPOs in those systems. In order to explain those QPOs consistently, we consider the interaction of the accreting plasma with the neutron star' magnetosphere.   Due to the strong gravity of the star, a very steep and supersonic flow hits the magnetosphere boundary and deforms its structure drastically. Such a plasma flow can readily excite different MHD waves in the magnetosphere, such as the shear Alfv\\'en waves.  In the Earth's magnetosphere, occurrence of aurora is a result of resonant coupling between the shear Alfv\\'en waves and compressional waves (produced by the solar wind). These resonances are known as FLRs that are reviewed in detail in section \\ref{sect:mag} of this paper. We argue that such resonant coupling is likely to occur in neutron star magnetospheres during its interaction with accreting plasma.  We formulated an improved FLR by considering a plasma flow moving with velocity $v_p$ along the magnetic field lines. Such flow is likely to occur in neutron star magnetosphere \\citep{GLP77}. For a simple geometry, the rectilinear magnetic field, and in the presence of a plasma flow we found: (a) two resonant hydrodynamical-MHD modes with frequencies $\\omega^{\\pm} = k_z(v_p \\pm v_{\\rm A})$. (b) the resulting frequencies for $k_z\\sim \\pi/R_s\\sim 3\\times 10^{-6}$ cm$^{-1}$ and typical flow and/or Alfv\\'en velocity $\\sim .1 c$ will be in kHz range. Our results match the kHz oscillations observed in the X-ray fluxes in LMXBs. As shown in figures \\ref{freq_vs_r} and \\ref{freq_vs_alfv}, the closer to the star and/or the larger the Alfv\\'en velocity the higher the frequency. (c) the quasi-periodicity of the observed oscillations can be understood by noting that due to several processes such as magneto-turbulence at boundaries and the time varying accretion rate, the FLR frequencies may vary with time. (d) a non-zero plasma displacement along the magnetic field lines $\\xi_{||}$. Such a displacement that oscillates with resulting frequencies $\\omega^\\pm$ modulates the flow of the plasma toward the surface of neutron star.  As a result, the X-ray flux from the star will show these frequencies as well. (e) setting $\\Delta\\nu=\\omega^-$ and $\\nu_2=\\omega^+$, one can well explain the behavior of the peak separation frequency $\\Delta\\nu$ relative to the upper QPO frequency $\\nu_2$.  As observed, the value of $\\Delta\\nu$ decreases as the magnitude of $\\nu_2$ decreases and/or increases. Figure \\ref{peak_freq} clearly shows such behavior. (f) for $v_p\\sim 2 v_{\\rm A}$ at $r\\sim 3 R_s$, the frequency ratio $\\nu_1/\\nu_2$ is comparable with the observed frequency ratio $2:3$.  Interestingly, using the observed values of QPO frequencies, one can determine the average mass density and the magnetic field density of the star as \\st \\bea \\stq &&k_z v_p(r)=(1/2)(\\nu_2+\\Delta\\nu)=(\\nu_2-\\nu_1/2),\\\\ \\stq &&k_z v_{\\rm A}(r)=(1/2)(\\nu_2-\\Delta\\nu)=\\nu_1/2\\,. \\eea Therefore, \\st \\bea \\label{den} \\stq &&M_s/R_s^3 \\simeq (1/8\\pi^2 G) x^3 (\\nu_2+\\Delta\\nu)^2 ,\\\\ \\stq &&\\mu_s\\simeq  (32\\pi^4 G)^{-1/4} R_s^{5/2} \\dot{M}^{1/2} (M_s/R_s^3)^{-1/4} x^{13/4} \\nu_1, \\eea where $k_z\\simeq \\pi/r$, $\\mu_s=B_0R_s^3$,$x=r/R_s$, and $B_0$ is the magnetic field strength at the surface of the star.   In table 1, we calculate the average density $M_s/R_s^3$ and magnetic dipole moment $\\mu_s$ of the star for a fixed value $x=10$. Note that larger frequencies may occur at smaller distances, ie. $x<10$.  Our results are comparable with realistic neutron star parameters.  Furthermore, our model is able to explain those low frequency ($\\sim 10$ Hz) quasi-periodic oscillations observed in the rapid burster such as MXB 1730-335 and GRO J1744-28 \\citep{Mas00}.  Such low frequencies can be extracted from the model by considering smaller inflow/Alfv\\'en velocity and/or further distances from the star.  Nevertheless, in order to avoid a number of complexities in our calculations, we used approximations such as slow rotation and cold plasma.  These assumptions may put some restriction on the validity of our model and results.  Future studies will be devoted to overcoming those restrictions.  VR wishes to thank Mariano Mendez for kindly providing QPO data. This research was supported by the National Sciences and Engineering Research Council of Canada.  \\def\\aj{{AJ}}                   % \\def\\araa{{ARA\\&A\\ }}             % \\def\\apj{{ApJ\\ }}                 % \\def\\apjl{{ApJ\\ }}                % \\def\\apjs{{ApJS\\ }}               % \\def\\apss{{Ap\\&SS}}             % \\def\\aap{{A\\&A\\ }}                % \\def\\aapr{{A\\&A~Rev.}}          % \\def\\aaps{{A\\&AS}}              % \\def\\azh{{AZh}}                 % \\def\\baas{{BAAS}}               % \\def\\jrasc{{JRASC}}             % \\def\\memras{{MmRAS}}            % \\def\\mnras{{MNRAS\\ }}             % \\def\\pra{{Phys.~Rev.~A}}        % \\def\\prb{{Phys.~Rev.~B}}        % \\def\\prc{{Phys.~Rev.~C\\ }}        % \\def\\prd{{Phys.~Rev.~D\\ }}        % \\def\\pre{{Phys.~Rev.~E}}        % \\def\\prl{{Phys.~Rev.~Lett.\\ }}    % \\def\\pasp{{PASP}}               % \\def\\pasj{{PASJ\\ }}               % \\def\\qjras{{QJRAS}}             % \\def\\skytel{{S\\&T}}             % \\def\\solphys{{Sol.~Phys.}}      % \\def\\sovast{{Soviet~Ast.\\ }}      % \\def\\ssr{{Space~Sci.~Rev.\\ }}     % \\def\\zap{{ZAp}}                 % \\def\\nat{{Nature\\ }}              % \\def\\iaucirc{{IAU~Circ. No.}}       % \\def\\aplett{{Astrophys.~Lett.}} % \\def\\apspr{{Astrophys.~Space~Phys.~Res.}} \\def\\bain{{Bull.~Astron.~Inst.~Netherlands}} \\def\\fcp{{Fund.~Cosmic~Phys.}}  % \\def\\gca{{Geochim.~Cosmochim.~Acta}}   % \\def\\grl{{Geophys.~Res.~Lett.}} % \\def\\jcp{{J.~Chem.~Phys.}}      % \\def\\jgr{{J.~Geophys.~Res.}}    % \\def\\jqsrt{{J.~Quant.~Spec.~Radiat.~Transf.}} \\def\\memsai{{Mem.~Soc.~Astron.~Italiana}} \\def\\nphysa{{Nucl.~Phys.~A}}   % \\def\\nphysb{{Nucl.~Phys.~B\\ }}   % \\def\\physrep{{Phys.~Rep.}}   % \\def\\physscr{{Phys.~Scr}}   % \\def\\planss{{Planet.~Space~Sci.}}   % \\def\\procspie{{Proc.~SPIE}}   %   ----------------------------------------------------------------------------------------"
    },
    "0403/astro-ph0403229_arXiv.txt": {
        "abstract": "Weak lensing provides a direct measure of the distribution of mass in the universe, and is therefore a uniquely powerful probe of dark matter. Weak lensing can also be used to measure the twin phenomenon of dark energy, via its effect upon the cosmological growth rate of structures. Essential for this technique are well-resolved images of background galaxies out to large distances. As a concrete example of the surveys that will become available by the end of the decade, we consider the planned {\\it Supernova/Acceleration Probe} (SNAP) space telescope. Detailed simulations of space-based images, manufactured using the shapelets formalism, enable us to quantitatively predict the future sensitivity to weak lensing shear. The high number density of galaxies resolved from space will enable maps of dark matter to be produced in two and three dimensions, with a resolution superior to that from the ground. Such observations will also afford reduced systematics for high-precision measurements of weak lensing statistics. These will be used to set tight constraints on cosmological parameters. In particular, the parameter for equation of state of dark energy, $w$, will be measured using weak lensing with a precision comparable to and somewhat orthogonal to constraints from other methods. ",
        "introduction": "\\label{intro}  During its journey to our telescopes, the light from background galaxies is deflected by foreground mass concentrations, which act as gravitational lenses along the line-of-sight. The observed distortions in the shapes of distant galaxies are directly related to the gravitational potential of foreground structures, and therefore to their mass. This effect is independent of the nature or state of the foreground mass and therefore traces the distribution of otherwise invisible dark matter. In recent years, many groups have measured coherent distortions in the shapes of galaxies to measure the mass distribution of clusters ({\\it e.g.}~Joffre {\\it et al.}~2000; Clowe, Trentham \\& Tonry 2001; Dahle {\\it et al.}~2002). Other groups have measured the distribution of large-scale structure in random regions in the sky to set constraints on cosmological parameters ({\\it e.g.}~van Waerbeke {\\it et al.}~2002; Bacon {\\it et al.}~2002; Hoekstra {\\it et al.}~2002; Jarvis {\\it et al.}~2003). The second moment of the resulting ``cosmic shear'' field already provides a constraint on the amplitude of the mass power spectrum $\\Omega_m \\sigma_8^{0.5}$ which is comparable to the constraints set by more traditional methods such as the abundance of x-ray selected clusters ({\\it e.g.}~Viana \\& Liddle~1999; Pierpaoli, Scott \\& White~2001). The higher order moments of the shear field promise to break the $\\Omega_m$-$\\sigma_8$ degeneracy (Bernardeau et al. \\etal1997) in the next generation of large cosmic shear surveys, {\\it e.g.}~CFHT Legacy Survey (CFHTLS, Mellier \\etal2000), LSST (Tyson \\etal 2002a) and Pan-STARRS (Kaiser, Tonry \\& Luppino 2000). For reviews of the techniques and current status of weak lensing measurements, other than this volume, see Mellier et al. (2001), Hoekstra et al. (2002), Wittman (2002) and Refregier (2003).  The rapid growth of dedicated surveys and improvements in shear measurement methods has meant that errors on cosmological parameters from cosmic shear are now limited by systematic rather than statistical errors. Weak lensing has the advantage that its systematics arise from imperfect instruments and image analysis, rather than unknown physics (such as the mass-temperature relation which dominates $x$-ray selected cluster samples (Huterer \\& White, 2002)). Image analysis tools are currently making important advances, from the well-tested method of Kaiser, Squires \\& Broadhurst (1995; KSB), with work by Rhodes, Refregier \\& Groth (2000; RRG), Bernstein \\& Jarvis (2003) and ``shapelets'' by Refregier (2003a). Meeting the engineering challenge, however, will inevitably require rising above the atmosphere. Convolution with a large Point-Spread Function (PSF) that is larger than the majority of distant (and therefore small) objects irretrievably destroys any information their shapes had contained. This limits the number density and redshift of galaxies that can be used to measure cosmic shear with ground-based lensing surveys. Convolution with the time-varying atmospheric PSF further limits the recovery of the shape information of the remaining galaxies to the accuracy with which the PSF can be modelled from one exposure to the next.  In these proceedings, we discuss the potential advances that a wide-field imager from space offers for weak lensing. Adopting the planned {\\it Supernova/ Acceleration Probe} (SNAP) mission as a concrete example, we present detailed shapelet-based simulations used to estimate the lensing sensitivity of such observations. Most interestingly for cosmology, SNAP's enhanced spatial resolution will capture the shapes of background galaxies up to a much greater distance. Space-based surveys will therefore probe the evolution of the mass distribution, and the growth of structures, over a large fraction of time in the evolution of the universe. This evolution provides strong constraints on both $\\Omega_m$ and the equation of state parameter of dark energy $w$.  As discovered from observations of type Ia superov\\ae~({\\it e.g.}~Perlmutter \\etal1999), this dark energy or ``quintessence'' is accelerating the expansion of the universe and hence retarding the growth of the mass power spectrum. Earlier studies of the constraints on dark energy from weak lensing surveys can be found in Hu (1999), Huterer (2001), Benabed \\& Bernardeau (2001), Hu (2001), Weinberg \\& Kamionkowski (2002), Munshi \\& Wang (2002), Benabed \\& van Waerbeke (2003). We also compare the quality of dark matter maps that will be possible from SNAP with those possible from current ground-based surveys. High resolution maps from space will enable the production of a fully mass-selected cluster catalogue, useful for a further constraint on cosmological parameters (Miyazaki \\etal 2002) but also for investigations into astrophysical effects during the non-linear infall and growth on smaller scales. Full details of our calculations are given in a series of papers by Rhodes \\etal(2004), Massey \\etal(2004b) and Refregier \\etal(2004b).  \\begin{figure} \\centerline{\\psfig{figure=fig01_lowres.ps,width=110mm}} \\caption{Illustration of the effect of gravitational lensing. \\hspace{2in}Foreground dark matter haloes differentially shear the light paths from background galaxies, distorting their shapes into an ellipse (when averaged over many galaxy shapes) or arcs near the centres of clusters. \\label{fig:cartoon}} \\end{figure} ",
        "conclusions": "A space-based, wide-field imager is ideal for weak lensing measurements. Indeed, a systematic floor due to atmospheric seeing will be reached in the next generation of lensing surveys. This is due both to the isotropic PSF smearing that reduces the number of resolved galaxies for which shape measurements are possible; and anisotropic smearing which cannot be perfectly corrected when it varies from one exposure to the next. The proposed SNAP satellite will have both the wide field needed to survey large, representative cosmic volumes and the low level of instrumental systematics affecting galaxy shapes.  Through a dual wide and deep survey strategy, weak lensing with SNAP will be able to produce unique maps of the dark matter distribution, on small scales and in both 2D and 3D. Weak lensing with SNAP will be able to produce mass-selected cluster catalogues to a $1\\sigma$ detection threshold of $10^{13}M_{\\sun}$ at $z=0.25$.  Weak lensing primarily measures the distribution of the dark matter, but it is also a powerful probe of dark energy. A change in the equation of state $w$ of dark energy modifies the growth rate of structures and the angular-diameter distance. The resulting changes in the cosmic shear statistics will be easily detectable with future weak lensing surveys. In particular, weak lensing measurements with SNAP will be able to independently constrain cosmological parameters $\\Omega_M$, $\\sigma_8$ and $w$ at a level comparable with and somewhat orthogonal to those from future supernova searches and CMB experiments.  \\newpage"
    },
    "0403/astro-ph0403535_arXiv.txt": {
        "abstract": "{ We use the very deep and homogeneous I-band selected dataset of the FORS Deep Field (FDF) to trace the evolution of the luminosity function over the redshift range \\mbox{$0.5 < z < 5.0$}. We show that the FDF I-band selection down to $I_{AB}=26.8$ misses of the order of 10~\\% of the galaxies that would be detected in a K-band selected survey with magnitude limit $K_{AB}=26.3$ (like FIRES). Photometric redshifts for 5558 galaxies are estimated based on the photometry in 9 filters (U, B, Gunn g, R, I, SDSS z, J, K and a special filter centered at 834~nm).  A comparison with 362 spectroscopic redshifts shows that the achieved accuracy of the photometric redshifts is \\mbox{$\\Delta z / (z_{spec}+1) \\le 0.03$} with only $\\sim 1$\\% outliers. This allows us to derive luminosity functions with a reliability similar to spectroscopic surveys. In addition, the luminosity functions can be traced to objects of lower luminosity which generally are not accessible to spectroscopy. We investigate the evolution of the luminosity functions evaluated in the restframe UV (1500~\\AA\\ and 2800~\\AA), u', B, and g' bands. Comparison with results from the literature shows the reliability of the derived luminosity functions. Out to redshifts of $z\\sim 2.5$ the data are consistent with a slope of the luminosity function approximately constant with redshift, at a value of $-1.07 \\pm 0.04$ in the UV (1500~\\AA\\ , 2800~\\AA) as well as u', and $-1.25 \\pm 0.03$ in the blue (g', B). We do not see evidence for a very steep slope (\\mbox{$\\alpha \\le -1.6$}) in the UV at \\mbox{$\\langle z \\rangle\\sim 3.0$} and \\mbox{$\\langle z \\rangle\\sim 4.0$} favoured by other authors. There  may be a tendency for the faint-end slope to become shallower with increasing redshift but the effect is marginal. We find a brightening of M$^\\ast$ and a decrease of $\\phi^\\ast$ with redshift for all analyzed wavelengths. The effect is systematic and much stronger than what can be expected to be caused by cosmic variance seen in the FDF. The evolution of M$^\\ast$ and $\\phi^\\ast$ from $z=0$ to $z=5$ is well described by the simple approximations \\mbox{$M^\\ast(z)= M^\\ast_0 + {a} \\ln(1+z)$} and \\mbox{$\\phi^\\ast(z)= \\phi^\\ast_0 (1+z)^{b}$} for $M^\\ast$ and $\\phi^\\ast$.  The evolution is very pronounced at shorter wavelengths (\\mbox{$a=-2.19$, and $b=-1.76$ for 1500~\\AA\\ rest wavelength}) and decreases systematically with increasing wavelength, but is also clearly visible at the longest wavelength investigated here (\\mbox{$a=-1.08$, and $b=-1.29$ for g'}).  Finally we show a comparison with semi-analytical galaxy formation models. ",
        "introduction": "\\label{sec:intro}    Observational constraints on galaxy formation have improved significantly over the last years and it has become possible to study the evolution of global galaxy properties up to very high redshifts. A crucial step to probe the properties of galaxies up to the highest redshifts was the work of \\citet{steidel_lbg:1} and \\citet{steidel_lbg:2} who used color selection to discriminate between low redshift and high redshift galaxies.  Although the Lyman-break technique is very efficient in selecting high redshift galaxies (see \\citealt{blaizot} for a detailed discussion) with a minimum of photometric data, it has the disadvantage that it does not sample galaxies homogeneously in redshift space and may select against certain types of objects.  With the advent of deep multi-band photometric surveys (Hubble Deep Field North (HDFN; \\citealt{HDF96}), NTT SUSI deep Field (NDF; \\citealt{arnouts_ntt}), Hubble Deep Field South (HDFS; \\citealt{HDFS00,HDFS00a}), Chandra Deep Field South (CDFS; \\citealt{arnouts_cdfs}), William Herschel Deep Field (WHDF; \\citealt{mccracken:1,metcalf:1}), Subaru Deep Field/Survey (SDF; \\citealt{maihara:1,ouchi:2}),  The Great Observatories Origins Deep Survey (GOODS; \\citealt{giavalisco:1})) the photometric redshift technique (essentially a generalization of the drop-out technique) has increasingly been used to identify high-redshift galaxies. Several methods have been described in the literature to derive photometric redshifts \\citep{baum:1, koo:1, brunner:1, Soto:1, benitez:1, borgne:1, firth:1}.  Based on either spectroscopic redshifts, drop-out techniques, or photometric redshifts, it has been possible to derive luminosity functions at different redshifts in the ultraviolet (UV) \\citep{treyer:98,steidel:1,cowie:1, adelberger:1, Cohen:1, sullivan:1, ouchi:1, poli:1, wilson:1, combo17:1, rowan:1, kashikawa:1, ouchi:2, iwata:1} and in the blue bands \\citep{lilly:2,heyl:1, lin:1, sawicki:1, small:1, zucca:1, loveday:1, marinoni:1, fried:1, cross:1, im:1, marinoni:2, norberg:1, bell:1, lapparent:1, liske, poli:3, prezgonz:1}.  Within the uncertainties given by IMF and dust content, the flux in the UV allows to trace the star formation rate (SFR; \\citealt{mad_poz_dick1}) in the galaxies, while the optical luminosities provide constraints on more evolved stellar populations \\citep{franx:1}.  Locally, the 2dF Galaxy Redshift Survey (2dFGRS; \\citealt{colless:1}), the Sloan Digital Sky Survey (SDSS; \\citealt{stoughton:1}) and the 2MASS survey \\citep{2MASS} have provided superb reference points for galaxy luminosity functions over a large wavelength range (see \\citealt{norberg:1} for 2dFGRS, \\citealt{blanton:1, blanton:2} for the SDSS, and \\citealt{Kochaneketal01, Coleetal2001} for  2MASS).   In parallel to the observational effort, theoretical models have been developed within the framework of the cold dark matter cosmology. Most notably, semi-analytic models (SAMs) \\citep{kauffmann:1, cole:1, sommerville:1, kauffmann:2, poli:2, wu:1, cole:2, menci:1, menci:2} and simulations based on smoothed-particle hydrodynamics (SPH) \\citep{dave:1, weinberg:1, nagamine:1, nagamine:2} have made testable predictions.  Starting with the mass function of dark matter halos and their merging history, SAMs use simplified recipes to describe the baryonic physics (gas cooling, photoionization, star formation, feedback processes, etc., see \\citealt{benson:1}) to derive stellar mass and luminosity functions.  Ideally, a comparison between observations and models should be done with deep multiwavelengths datasets that also cover a large area.  The dataset has to be sufficiently deep in order to be able to derive the faint-end slope of the luminosity function. On the other hand, one also needs as large an area as possible to overcome cosmic variance and to quantify the density of rare bright galaxies, which define the cut-off of the luminosity function.  The FORS Deep Field \\citep{fdf_data} has a depth close to the HDFs but an area of \\mbox{8 -- 10} times the area of the HDFN.  This depth allows us to detect galaxies at $z > 2$ which would be missed by Lyman-break studies which usually reach only $R_{AB} < 25.5$ (see also \\citealt{franx:1} and \\citealt{dokkum:1}).   Very reliable photometric redshifts are crucial for the analysis of the evolution of the luminosity functions in the FDF. Photometric redshifts have been determined with a template matching algorithm described in \\citet{bender:1} that applies Bayesian statistics and uses semi-empirical template spectra matched to broad band photometry. We achieved an accuracy of \\mbox{$\\Delta z / (z_{spec}+1) \\le 0.03$} with only $\\sim 1$\\% extreme outliers (numbers based on a comparison with 362 spectroscopic redshifts).  Redshifts of galaxies that are several magnitudes fainter than typical spectroscopic limits could be determined reliably and thus allowed better constraints on the faint-end slope of the luminosity functions.   In this paper we present the redshift evolution of the luminosity function evaluated in the restframe UV-range (1500~\\AA, 2800~\\AA), u' (SDSS), B, and g' (SDSS) bands in the redshift range \\mbox{$0.5 < z < 5.0$}.  Luminosity functions at longer wavelengths as well as the evolution of the luminosity density and the star formation rate will be presented in future papers (Gabasch et al., in preparation). We provide a short description of the FDF in Sect.~\\ref{sec:fdf} where we also present the selection criteria of our galaxies. In Sect.~\\ref{sec:lumfkt:selection} we investigate possible selection effects due to our purely I-band selected catalogue. In Sect.~\\ref{sec:photred} we discuss the accuracy of the photometric redshifts as well as the redshift distribution of the selected galaxies.  In Sect.~\\ref{sec:lumfkt} and in the Appendix we show luminosity functions at different wavelengths and redshifts. In Sect.~\\ref{sec:evol_parameter}, a parameterization of the redshift evolution of the Schechter \\citep{schechter:1} parameters M$^\\ast$ and $\\phi^\\ast$ is given.  We compare our results with previous observational results in Sect.~\\ref{sec:lit}, and with model predictions in Sect.~\\ref{sec:model}, before we summarize this work in Sect.~\\ref{sec:summary_conclusion}.  We use AB magnitudes and adopt a $\\Lambda$ cosmology throughout the paper with \\mbox{$\\Omega_M=0.3$}, \\mbox{$\\Omega_\\Lambda=0.7$}, and \\mbox{$H_0=70 \\, \\mathrm{km} \\, \\mathrm{s}^{-1} \\, \\mathrm{Mpc}^{-1}$}. ",
        "conclusions": "\\label{sec:summary_conclusion}   We analyzed a sample of about 5600 I-band selected galaxies in the FORS Deep Field down to a limiting magnitude of I $ = 26.8$ mag.  A comparison with the very deep K-selected catalogue of \\citet{labbe:1} shows, that more than 90~\\% of their objects are brighter than our limiting I-band magnitude.  Therefore our scientific conclusions are not affected by this color bias.   Based on 9 filters we derived accurate photometric redshifts with \\mbox{$\\Delta z / (z_{spec}+1) \\approx 0.03 $} if compared with the spectroscopic sample \\citep{noll:1, boehm:1} of 362 objects. We calculated and presented the luminosity functions in the UV (1500~\\AA\\ and 2800~\\AA), u', B, and g' bands in the redshift range \\mbox{$0.5 < z < 5.0$}.  The error budget of the luminosity functions includes both the photometric redshift error as well as the Poissonian error.  The faint-end slope of the luminosity function does not show a large redshift evolution and is compatible within $2\\sigma$ with a constant slope in most of the redshift bins and wavelengths considered here. Furthermore, the slope in the 1500~\\AA, 2800~\\AA, and u' band is very similar but differs from the slope in the g' and B band.  We derive a best fitting slope of $\\alpha=-1.07 \\pm 0.04$ for the combined 1500 \\AA , 2800~\\AA\\ and u' bands and $\\alpha=-1.25 \\pm 0.03$ for the combined g' and B bands.  We find no evidence for a very steep slope (\\mbox{$\\alpha \\le -1.6$}) at $z \\sim 3$ and 1700~\\AA\\ rest wavelength as reported by other authors (e.g., \\citealt{steidel:1}, \\citealt{ouchi:3}). From our data we can exclude a slope of \\mbox{$\\alpha \\le -1.6$} at redshift \\mbox{$\\langle z \\rangle\\sim 3.0$} and \\mbox{$\\langle z \\rangle\\sim 4.0$} at the $2\\sigma$ level.  We investigate the evolution of M$^\\ast$ and $\\phi^\\ast$ by means of a redshift parameterization of the form \\mbox{$M^\\ast(z)= M^\\ast_0 + a \\ln(1+z)$} and $\\phi^\\ast(z)=\\phi^\\ast_0(1+z)^b$. We find a substantial brightening of M$^\\ast$ and a decrease of $\\phi^\\ast$ with redshift in all analyzed wavelengths. If we follow the evolution of the characteristic luminosity from \\mbox{$\\langle z \\rangle\\sim 0.5$} to \\mbox{$\\langle z \\rangle\\sim 5$}, we find an increase of $\\sim$~3.1 magnitudes in the UV, of $\\sim$~2.6 magnitudes in the u' and of $\\sim$~1.6 magnitudes in the g' and B band.  Simultaneously the characteristic density decreases by about 80 \\% -- 90 \\% in all analyzed wavebands.  Moreover, we compare the luminosity function derived in the FDF with previous observational datasets, mostly based on photometric results, and discuss discrepancies. In general, we find good agreement at the bright end, where their samples are complete.  Differences in the faint-end slope in some cases can be attributed to the shallower limiting magnitudes of most of the other surveys.  The only observations which reach the same limiting magnitudes as the FDF observations are those of \\citet{poli:1, poli:3} and the K-selected sample of \\citet{kashikawa:1}. The FDF results for the faint-end slope are in excellent agreement with those of \\citet{kashikawa:1} but the slope of the Schechter function favored by \\citet{poli:1, poli:3} is steeper than we would expect from the FDF.    The semi-analytical models predict luminosity functions which describe (by construction) the data at low redshift quite well, but show growing disagreement with increasing redshifts."
    },
    "0403/hep-th0403044_arXiv.txt": {
        "abstract": "We argue that a brane world with a warped, infinite extra dimension allows for the inflaton to decay into the bulk so that after inflation, the effective dark energy disappears from our brane. This is achieved by the redshifting of the decay products into infinity of the 5th dimension. As a consequence, all matter and CMB density perturbations could have their origin in the decay of a MSSM flat direction rather than the inflaton. We also discuss a string theoretical model where reheating after inflation may not affect the observable brane. ",
        "introduction": "Recent observations~\\cite{WMAP} strongly support a period of primordial inflation. Besides making the universe flat and homogeneous, inflation is the only known dynamical mechanism which can stretch small quantum fluctuations outside the Hubble horizon.  These perturbations act as seeds for the large scale structures in the Universe. However, despite of the success of the inflationary paradigm (for a review, see ~\\cite{Riotto}), we know very little about the inflationary sector. The inflaton potential must nevertheless be very flat with a very small self coupling; likewise, its coupling to other fields must be extremely weak.  Such small couplings are hard to come by without fine tuning, which renders inflaton to a gauge singlet whose couplings can be adjusted at our will. Given this, the immediate question is, what is the inflaton decaying into?   Eventually the cold inflationary Universe must reheat with the Standard Model (SM) degrees of freedom, or, as the current theoretical prejudice dictates, with the Minimally Supersymmetric Standard Model (MSSM) degrees of freedom. MSSM contains a number of flat directions along which the renormalizable part of the potential vanishes (for a review, see~\\cite{Enqvist}). During inflation the massless fields corresponding to the flat directions receive scale-invariant perturbations. It would be tempting to associate them with reheating and the generation of CMB perturbations \\cite{Kasuya2,Kasuya3,Kasuya1,Others}.  This would require that the flat directions will eventually dominate over the radiation produced by the inflaton decay. In other words, the flat direction would have to act as an MSSM curvaton~\\cite{Sloth,Lyth,Moroi}. However, it turns out that such domination is not possible unless the inflaton decays into some hidden degrees of freedom, rather than into MSSM radiation \\cite{Kasuya2,Kasuya3}.   In this respect brane world scenarios, where the Universe is regarded as a three dimensional hypersurface embedded in a higher dimensional bulk, bring along new, interesting possibilities. The SM degrees of freedom are assumed to be stuck on a brane while gravity is propagating in the entire bulk. The bulk could also have a non-trivial background geometry which allows for the zero mode of graviton to be trapped at the brane location, such as in the case of an anti de-Sitter (adS) bulk in the Randall-Sundrum type models~\\cite{RS1,RS2}. The value of the Newton's constant requires the fundamental scale to be fairly large, $10^{18}~{\\rm GeV}\\geq M_{s}\\geq 10^{3}$~GeV. In these models the inflaton, again treated as a gauge singlet, could either live on the brane or in the entire bulk. There is also the exciting possibility that the inflaton energy density does not reheat the Universe, but gets deposited in the infinite bulk or onto the adS horizon. In this respect the inflaton potential could be treated as a kind of dark energy for the MSSM brane. It has been shown \\cite{Kasuya1} that if the dark energy of the inflaton field can be transferred into the bulk so that it no longer is visible on our brane, even the simplest MSSM flat directions can act as curvatons and provide both all the matter and the density perturbations in the Universe.   The brane world scenarios are supported by the string theory, where there is a natural explanation for the construction of MSSM like brane with a help of coincident stack of $Dp$ branes ($p$ is the number of spatial dimensions along the brane, $p$ is odd (even) in type IIA (IIB) string theories) attached to some orbifold point~\\cite{Quevedo1}~\\footnote{For different constructions, see~\\cite{Quevedo3}.}. The open strings attached to $Dp$ branes act as sources for gauge fields and gravity again propagates in the entire bulk. In this kind of framework the inflaton may be regarded as a moduli, for example the physical separation between a $Dp$ and anti-brane $\\bar Dp$ brane, or the angular separation between $Dp-Dp$ or $Dp-\\bar Dp$ branes (for a review see~\\cite{Quevedo2}).  Inflation may end by virtue of tachyon condensation when branes approach close to the string scale~\\cite{Sen}, or with a help of many tachyons as in the case of assisted inflation~\\cite{Panda}. It is however not guaranteed that inflaton will reheat the Universe with the MS(SM) degrees of freedom. One could rather argue that it is more likely that the inflaton will reheat the bulk.   The purpose of this paper is to construct a brane world model with a warped bulk so that it is possible to localize the inflaton energy density away from the MS(SM) like branes. We will argue that the inflaton energy can be redshifted away so that after inflation there is effectively no energy density other than that of the excited MSSM flat directions on our brane. For our purposes, the inflaton potential is a form of dark energy which is only responsible for making the Universe, parallel to the brane directions, large, homogeneous and isotropic.    The paper is organized as follows. In Section II we recapitulate some known results of the brane world models and discuss infinite extra dimensions. In Section III we  present a brane-world where the dark energy can be dumped into the bulk instead onto the brane and estimate the escape rate of the inflaton decay products from the brane. In Sect. IV we motivate the model by considerations relating to string theoretical inflation. In Sect. V we give our conclusions and discuss how the present results can be combined with curvaton-like scenarios involving MSSM flat directions to yield all matter and an adiabatic, scale-invariant spectrum of perturbations. ",
        "conclusions": "During inflation, massless MSSM fields, or the (almost) flat directions corresponding to certain combinations of squarks, sleptons and Higgses, will be subject to fluctuations and form condensates. Like the ordinary inflaton, the condensates will receive scale-invariant spatial perturbations. Once the inflaton energy has disappeared into the bulk, the potential terms along some flat direction will eventually start to dominate the energy density of the MSSM brane. A potential arises because of non-renormalizable interactions and the soft supersymmetry breaking mass terms~(for a review, see \\cite{Enqvist}). (In the cosmological context the Higgs coupling $\\mu H_{u}H_{d}$ does not spoil flatness as $\\mu$ is much smaller than the relevant field amplitudes).  Once the condensate decays, it will reheat the universe with MSSM degrees of freedom and imprint on the MSSM gas the inflationary perturbations. The flat direction condensate acts as an MSSM curvaton~\\cite{Sloth,Lyth,Moroi}. The simplest possibility is the flat direction that consists of the Higgses $H_u$ and $H_d$, which has been discussed in detail in \\cite{Kasuya1}. The reheat temperature can be estimated to be less than $10^9$ GeV. The amplitude of the fluctuations along the $H_{u}H_{d}$ direction can match the observed density perturbations in the CMB radiation and the spectrum with a spectral tilt very close to 1, with some weak dependence on the Higgs potential~\\cite{Kasuya1}.  We have argued that a brane cosmology with a warped, infinite extra dimension allows for the inflaton to decay into the bulk and for the subsequent redshifting of the decay products. The inflaton decays efficiently into a continuum of Kaluza Klein modes which carry non-zero momentum along the extra dimension and move away from our four dimensional world, taking inflaton energy with them. In effect, the effective dark energy present on the MSSM brane will disappear into the bulk and be hidden behind the particle horizon at the infinity of the 5th dimension.   Although we have sketched a possible inflationary scenario involving $\\bar D3$ and $D3$ branes that annihilate at an adS throat, a string theoretical model for reheating remains a challenge. Nevertheless, it seems that at least within brane world cosmologies it is possible to have the inflaton decay products disappearing into the bulk so that all matter could have its origin in the decay of the MSSM condensate rather than in the inflaton energy density."
    },
    "0403/astro-ph0403620_arXiv.txt": {
        "abstract": "{ Since the H$\\alpha$ and UV fluxes from galaxies are sensitive to stellar populations of ages $< 10^{7}$ and $\\approx 10^{8}$~yr respectively, their ratio $f$(H$\\alpha$)/$f$(UV) provides us with a tool to study the recent $t \\leq 10^{8}$~yr star formation history of galaxies, an exercise that we present here applied to 98 galaxies in 4 nearby clusters (Virgo, Coma, Abell~1367 and Cancer). The observed $f$(H$\\alpha$)/$f$(UV) ratio is $\\sim$ a factor of two smaller than the expected one as determined from population synthesis models assuming a realistic delayed, exponentially declining star formation history. We discuss various mechanisms that may have affected the observed $f$(H$\\alpha$)/$f$(UV) ratio and we propose that the above discrepancy arises from either the absorption of Lyman continuum photons by dust within the star formation regions or from the occurrence of star formation episodes. After splitting our sample into different subsamples according to evolutionary criteria we find that our reference sample of galaxies unaffected by the cluster environment show an average value of $f$(H$\\alpha$)/$f$(UV) two times lower than the expected one. We argue that this difference must be mostly due to absorption of $\\approx 45$\\% of the Lyman continuum photons within star forming regions. Galaxies with clear signs of an ongoing interaction show average values of $f$(H$\\alpha$)/$f$(UV) slightly higher than the reference value, as expected if those objects had SFR increased by a factor of $\\simeq 4$. The accuracy of the current UV and H$\\alpha$ photometry is not yet sufficient to clearly disentangle the effect of interactions on the $f$(H$\\alpha$)/$f$(UV) ratio, but significant observational improvements are shortly expected to result from the GALEX mission. ",
        "introduction": "A number of environmental mechanisms able of affecting significantly the evolution of galaxies in rich clusters have been proposed in the literature: gas stripping by ram pressure (Gunn \\& Gott 1972; Abadi, Moore \\& Bower 1999), galaxy-galaxy harassment in close encounters (Moore et al. 1996), tidal stirring by the cluster potential (Byrd \\& Valtonen 1990; Fujita 1998).    These mechanisms should produce morphological disturbances, gas removal and, on long timescales, significant quenching of the star formation rates (SFRs) of galaxies due to ``fuel'' exhaustion (see Gavazzi et al. 2002a). However galaxy-galaxy interactions might also enhance the star formation in gas-rich systems, both in their nuclei and disks, as it has been observed at several wavelengths (Larson \\& Tinsley 1978; Kennicutt et al. 1987; Hummel et al. 1987; Soifer et al. 1984), although this {\\bf enhancement} might be mild (Bergvall et al. 2003). The dynamical interaction of galaxies with the IGM can also produce an enhancement in the galaxies SFR by ram pressure (Fujita \\& Nagashima 1999).  A conclusive evidence for a separate evolution of galaxies in clusters is offered by the Butcher-Oemler effect (Butcher \\& Oemler 1978), i.e. distant ($z \\sim 0.3$) clusters show a larger fraction of blue galaxies than nearby clusters. Several follow-up studies (Couch \\& Sharples 1987; Barger et al. 1996; Couch et al. 1994; Poggianti et al. 1999; Balogh et al. 1999) brought to today's accepted scenario that clusters are continuously accreting galaxies from their neighborhood, with the accretion rate increasing with look-back time.   Several observable quantities have been proposed as reliable estimators of the SFRs of galaxies (Kennicutt 1998; Rosa-Gonz\\'{a}lez et al. 2002): H$\\alpha$, UV, radio continuum and Far-IR luminosities. Among these, we focus our analysis on the H$\\alpha$ and UV luminosities. The H$\\alpha$ luminosity comes from stars more massive than 10M$_{\\odot}$ and it traces the SFR in the last $\\leq 10^{7}$~yr. The UV luminosity at 2000~\\AA\\ comes from stars more massive than 1.5M$_{\\odot}$ and it can be used as an indicator of the SFR in the last $\\approx 10^{8}$~yr, under the condition that it stayed approximately constant during this period. The two quantities combined, in other words the ratio $f$(H$\\alpha$)/$f$(UV), should give us a ``clock'' suitable for telling if the SFR was constant over the last $10^{8}$~yr.  The present paper is aimed at studying the role of the cluster environment on the star formation histories of cluster galaxies by using the $f$(H$\\alpha$)/$f$(UV) ratio for a sample of galaxies in four nearby clusters: Virgo, Coma, Abell~1367 and Cancer. This analysis relies on the multifrequency database that we collected so far for a large sample of galaxies in nearby cluster and we made available to the community through the GOLDmine WEB site (Gavazzi et al. 2003). Beside the H$\\alpha$ and UV data which are directly used for computing the two SFR indicators, other corollary data (e.g. Near-IR, Far-IR, H{\\sc i} fluxes and optical spectroscopy) are used throughout this paper. These corollary data play a fundamental role in the determination of the dust extinction at UV wavelengths (through the FIR/UV ratio, e.g. Buat \\& Xu 1996) and at H$\\alpha$ (from the Balmer decrement, e.g. Lequeux et al. 1981).  The galaxy sample is presented in Section~2. The observed vs. expected $f$(H$\\alpha$)/$f$(UV) ratio for cluster galaxies is discussed in Section~3. In Section~4 we discuss the limitations and the potentiality of the method applied in this preliminary analysis. A brief summary of the results is presented in Section~5. Details about the estimate of the birthrate parameter $b$ are given in Appendix~A. A second appendix contains a detailed analysis of the observational uncertainty affecting the $f$(H$\\alpha$)/$f$(UV) ratio. ",
        "conclusions": "The $f$(H$\\alpha$)/$f$(UV) ratio of cluster galaxies is analyzed in this paper as a promising tool to estimate if the the star formation history of galaxies has remained constant on timescales of $\\simeq 10^{8}$~yr. The observed $f$(H$\\alpha$)/$f$(UV) distribution is compared to the one predicted by models of galaxies, assuming a continuum SFH. The dispersion of the observed $f$(H$\\alpha$)/$f$(UV) distribution is consistent with the one expected from the observational uncertainties. We find a systematic negative difference between the average observed value and the model predictions. We discuss some mechanisms that could possibly produce such an observed difference and we highlight the two most likely ones: the absorption and, in a minor way, the escape of Lyman continuum photons and the occurrence of star formation bursts overimposed to a smooth SFH.  The $f$(H$\\alpha$)/$f$(UV) distribution is considered for different galaxy subsamples, each of them comprising galaxies in different evolutionary stages, possibly induced by the cluster environment. The {\\em ``reference''} unperturbed galaxies have $f$(H$\\alpha$)/$f$(UV) lower by 0.34 dex on average than the one predicted by the models. We suggest that absorption (and to a lesser extent escape) of Lyman continuum photons causes the observed discrepancy. We estimate that about 45\\% of the Lyman continuum photons are absorbed by dust in the star forming regions before ionization, consistently with the estimate of Hirashita et al. (2003) on similar objects.  When galaxies with signatures of recent or past interactions with the cluster environment ({\\em ``interacting''} and {\\em ``asymmetric''}) are considered, we find that their $f$(H$\\alpha$)/$f$(UV) ratio is slightly higher than the one of {\\em ``reference''} galaxies. Even though the absorption of Lyman continuum photons is taken into account, the observed $f$(H$\\alpha$)/$f$(UV) ratio can be reconciled to the predicted one only assuming that these objects underwent bursts of star formation of intensity  $\\sim$ 100 times larger than normal, as intense as Arp~220. Objects of this kind are however not presently observed in nearby clusters.  The present observational uncertainties on both the H$\\alpha$ and UV fluxes are still too large to allow disentangling the effects of recent star formation bursts from those of absorption of Lyman continuum photons. However we stress the potentiality of the proposed H$\\alpha$/UV method for studying the recent history of star formation in late type galaxies, once improvements in modeling the radiation transfer through the dust in star forming regions will be achieved and more precise UV and Far-IR photometry will be available. This will soon become a reality after the GALEX and ASTRO-F experiments will perform their all sky surveys, providing $\\Delta f(\\mbox{UV})$ and $\\Delta f(\\mbox{Far-IR}) \\approx 10$\\%."
    },
    "0403/astro-ph0403416_arXiv.txt": {
        "abstract": "The Hanle effect has been proposed as a new diagnostic of circumstellar magnetic fields for early-type stars, for which it is sensitive to field strengths in the 1-300~G range. In this paper we compute the polarized P-Cygni line profiles that result from the Hanle effect. For modeling the polarization, we employ a variant of the ``last scattering approximation''. For cases in which the Sobolev optical depths are greater than unity, the emergent line intensity is assumed to be unpolarized; while for smaller optical depths, the Stokes source functions for the Hanle effect with optically thin line scattering are used. For a typical P~Cygni line, the polarized emission forms in the outer wind, because the Sobolev optical depth is large at the inner wind.  For low surface field strengths, weak P~Cygni lines are needed to measure the circumstellar field. For high values of the surface fields, both the Zeeman and Hanle diagnostics can be used, with the Zeeman effect probing the photospheric magnetic fields, and the Hanle effect measuring the magnetic field in the wind flow.  Polarized line profiles are calculated for a self-consistent structure of the flow and the magnetic geometry based on the WCFields model, which is applicable to slowly rotating stellar winds with magnetic fields drawn out by the gas flow. For surface fields of a few hundred Gauss, we find that the Hanle effect can produce line polarizations in the range of a few tenths of a percent up to about 2 percent. ",
        "introduction": "The Hanle effect is a diagnostic that refers to the modification of resonance line scattering polarization in the presence of a magnetic field. The effect begins to have an influence on the polarization  at fairly small field strengths of just a few Gauss with sensitivity to fields up to around 300~Gauss.  Experiments to describe the polarization from resonance line scattering date back primarily back to the first third of the 20th century (see Mitchell \\& Zemansky 1934).  The influence of a magnetic field on the line polarization was explained first by a young physicist named Wilhelm Hanle.  Hanle (1924) described the change of linear polarization by the magnetic field in semi-classical terms as arising from the precession of an atomic, damped, harmonic oscillator. From a quantum mechanical point-of-view, the effect is understood in terms of interferences that occur when the degeneracy of the magnetic sublevels in the excited state is partially lifted. The effect has come to have applications to many topics in atomic physics as described in Moruzzi \\& Strumia (1991).  Astrophysically, the Hanle effect has been used only in the Sun. Applications include magnetic field measurements in the chromosphere, corona, and particularly in prominences and filaments (e.g., see Lin, Penn, \\& Kuhn 1998).  A detailed description of the physics of the Hanle effect and polarized radiation transport for solar studies appears in Stenflo (1994).  However, for stellar astronomy, the Hanle effect remains relatively unknown.  This paper is the fourth in a series to explore applications of the effect for other stars, especially hot stars with winds (Ignace, Nordsieck, \\& Cassinelli 1997 [Paper~I]; Ignace, Cassinelli, \\& Nordsieck 1999 [Paper~II]; Ignace 2001 [Paper~III]). It has been our goal to develop model line profiles with the Hanle effect that can be used to interpret spectropolarimetric data for inferring the properties of circumstellar magnetic fields.  The Zeeman effect has had some success in measuring {\\it photospheric} surface fields for a few hot stars that are not part of the extreme Bp class (e.g., Henrichs 2003; Donati \\etal\\ 2001, 2002; and Neiner 2002). However, Ignace \\& Gayley (2003) have shown that the Zeeman effect suffers serious technical challenges for use in diagnosing circumstellar magnetic fields from wind emission lines.  For example, typical circular polarizations for such lines have peak values of only 0.01\\% for a field with a surface value of around 100~G. Fields of this magnitude can have a significant influence on the wind flow from hot stars (e.g., Maheswaran \\& Cassinelli 1992; Babel \\& Montmerle 1997ab; Cassinelli \\etal\\ 2002; ud-Doula \\& Owocki 2002).  Small polarizations result for the Zeeman effect because the Zeeman-split components that produce the circular polarization are incoherent and oppositely signed. The wavelength separation of the Zeeman split lines is small compared with typical astrophysical broadening effects.  Consequently, the net circular polarization of the line is severely reduced owing to polarimetric cancellation from blending of the Zeeman components. The Hanle effect is a {\\em modification} of the linear polarization that arises from resonance line scattering that does not suffer from polarimetric cancellation due to line blending.  In a semi-classical description of the Hanle effect, the line scattering can be split into two parts: an isotropic scattering component and a dipole scattering component like free electrons (e.g., Hamilton 1947; Chandrasekhar 1960).  The fraction of the scattering that is dipole-like is referred to as the ``polarizability'', $E_1$; the isotropic scattering contribution is given fractionally as $(1-E_1)$.  The introduction of a magnetic field induces a precession of the classical oscillator motion of the atomic scatterers.  At the same time, the oscillations are damped, emitting their absorbed radiation at a rate given by the Einstein $A$-value for the transition of interest.  Thus there is a competition between the radiative rate versus the Larmor precessional rate that has the frequency, $\\omega_L = e B /m_e c$.  Using these same arguments, Hanle (1924) deduced that a significant effect results when these two rates are somewhat comparable.  We thus define a ``Hanle ratio'' given by  \\begin{equation} \\frac{B}{B_{\\rm Han}} = \\frac{2g_L \\omega_L}{A_{ul}} = g_L\\,\\frac{(B/5~G)}{(A_{ul}/10^8~s^{-1})}, \\end{equation}  \\noindent where $g_L$ is the Lande factor, and $A_{ul}$ is the Einstein $A$-value from upper level $u$ to lower level $l$.  The Hanle field strength $B_{\\rm Han}$ is a conveniently defined parameter given by  \\begin{equation} B_{\\rm Han} = \\frac{m_e\\,c\\,A_{ul}}{2\\,g_L\\,e}, \\end{equation}  \\noindent that characterizes the magnetic field scale at which the the Hanle effect is important.  Values of $B_{\\rm Han}$ for typical strong resonance lines common to astrophysics are displayed in Figure~\\ref{fig1}. Note that the lines at shorter wavelengths tend to have larger Hanle fields, because $A$-values increase roughly as $\\lambda^{-2}$.  Recognizing that the Hanle effect is sensitive to modest field values of 10--100~G, we explored in Paper~I its use for determining stellar wind magnetic field values and geometries from the total line polarization of optically thin emission lines. In Paper~II we considered the Hanle effect in optically thin and resolved line profiles that from in equatorial disks.  In Paper~III expressions were determined for the depolarizing effect of having stellar radiation arising from a finite disk instead of from a point star. For thin lines the correction was found to be the same as that for electron scattering polarization as derived by Cassinelli, Nordsieck, \\& Murison (1987). It was pointed out by Cassinelli, Nordsieck, \\& Ignace (2001) that in contrast with the solar case, the application of the Hanle effect to hot stars is somewhat simpler and less ambiguous. This is for two reasons. (a) The hot star wind lines are highly NLTE so that virtually all bound electrons of metal species are in the ground state, whereas detailed NLTE considerations are required in solar applications, and (b) the stellar continuum radiation that is being scattered is relatively ``flat'' over the frequency width of any given line, while solar applications often involve scattering of chromospheric emission lines for which the flux changes strongly with wavelength across the resonance line profiles.  The goal of these papers has been to move from an introduction of the Hanle effect for the broader stellar community to an investigation of its use by way of model calculations of increasing physical realism and numerical sophistication. In this paper we treat more realistic and more dynamically self-consistent magnetic field distributions in the wind, and we present results for the distribution of polarization across P~Cygni line profiles.  Section~2 outlines the line polarization calculations, and introduces a ``single scattering'' approximation that is used together with traditional Sobolev line profile analyses.  We find that line optical depth effects have a significant influence on the polarized line profiles.  Heuristic results are described in \\S 3, including simple spherical shells with axial fields, and spherical winds with dipole fields.  Results for a more realistic scenario for the field and flow are presented \\S 4. The model chosen for this is the slow magnetic rotator field configuration based on the WCFields theory by Ignace, Cassinelli, \\& Bjorkman (1998). Significant line polarizations of around 1\\% are found to result.  Lastly, \\S 5 offers a discussion of our modeling, with implications for observations of hot star winds at ultraviolet wavelengths and future directions for modeling efforts. ",
        "conclusions": "\\label{sec:disc}  This contribution has focussed on the Hanle effect in resonance scattering lines common to hot star winds.  We have stressed the fact that the Hanle effect is non-linear, in that it is the non-radial field components that give rise a change in the line polarization, yet at the same time, it is the total field strength (including the radial component) that governs the amount of the change.  Consequently, the Hanle effect can be used to infer the full 3D magnetic geometry in the wind.  For a given set of scattering lines, stronger fields will lead to higher line polarizations. Field topology also has an influence, as evidenced by the double-peaked polarized profiles for a dipole magnetic field, versus the single-peaked profiles for slow magnetic rotators.  The Far Ultraviolet Spectro-Polarimeter (FUSP) is a sounding rocket payload with the goal of making the first spectro-polarimetric measurements in the wavelength range of 1050--1500~\\AA\\ for several stellar targets (Nordsieck 2003).  With a 50~cm primary, the instrument will have a spectral resolution of 0.65~\\AA, corresponding to a resolving power of $\\lambda/\\Delta \\lambda = 1800$ at a wavelength of 1170~\\AA.  This in turn corresponds to a velocity resolution of about 170~km~s$^{-1}$, which is sufficient to resolve wind-broadened lines with typical full widths of a few thousand km~s$^{-1}$.  Targeting the star $\\zeta$~Ori, which was observed in the FUV with Copernicus (Snow \\& Morton 1976), it is anticipated that FUSP will produce the first detection of the Hanle effect in a star other than the Sun.  To correctly interpret the polarization of scattering lines that will be measurable with FUSP, care must be taken to identify all sources of polarized radiation.  For example, interstellar polarization is endemic to all polarimetric studies of distant objects, and there are standard techniques for its correction, based on wavelength dependence and lack of variability.  Extracting the influence of a non-spherical wind geometry for the line scattering polarization can be approached in two different ways.  (a) If the interstellar polarization can be determined, the polarization outside line frequencies due to Thomson scattering can be used to model the underlying wind geometry.  However, this is difficult in part because it may not be clear that the line forms in the same region that gives rise to the electron scattering polarization.  (b) Alternatively, since ${\\bf B}({\\bf r})$ is fixed for any given star, one can use a multi-line approach to infer the polarization arising from an aspherical distribution of scatterers.  Different lines have different Hanle field values, so that lines with large Hanle values can be used to estimate the non-magnetic contribution.  Once again, care must be taken since different lines may not form in the same spatial locations; however, the P~Cygni profiles themselves could be modelled to obtain information about the ionization distribution in radius and latitude.  The study of the Hanle effect for circumstellar environs and the development of radiative transport techniques is ongoing, especially in terms of relaxing the single scattering approximation for the line polarization.  In this regard one can use expressions in Jeffery (1990) for ``Sobolev-P'' theory to show that the single scattering approximation captures the flavor of the polarimetric behavior of scattering lines, but that in fact there is non-zero polarization from line scattering even at rather high Sobolev optical depths of a few, and even at optical depths of just a few tenths, the line polarization is not truly described by single scattering.  We anticipate using Monte Carlo radiation transport techniques in the future to model the line polarization more accurately. The proven versatility of these techniques will be useful in handling more realistic wind models.  Another effect ignored here needs to be explored. While our analysis applies directly to resonance transitions in ions with no nuclear spin (e.g., \\ion{O}{6}, \\ion{Si}{4}, \\ion{C}{4}, \\ion{Mg}{2}, and \\ion{Ca}{2}), in other Lithium-like ions such as \\ion{P}{5}, \\ion{N}{5}, \\ion{Na}{1} and \\ion{K}{1}, the presence of nuclear spin breaks the degeneracy of the spin states, and this alters the value of the polarizability $E_1$ making it sensitive to optical pumping.  In fact in the presence of optical pumping, the magnetic field has another effect on the line polarization through ``magnetic re-alignment'' (Nordsieck 2001), or in Solar Physics parlance, the ``second Hanle effect''.  We are investigating the regime in which the Hanle effect and the magnetic re-alignment physics both apply."
    },
    "0403/astro-ph0403285_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\indent \\m  The general paradigm in cosmology is that we are living in a flat universe, which is dominated by a nearly homogeneous component with negative pressure. This component is often called Dark Energy (DE) and causes the expansion rate of the universe to accelerate.  The recent measurements of type Ia supernovae (thereafter denoted SN) are the most direct evidence of the presence of this component \\cite{SCP,Riess,Tonry,newSCP}. It is also confirmed by the combination of results from the large-scale distribution of galaxies \\cite{Efst2DF} and  the most precise data on the cosmic microwave background (CMB) from the Wilkinson Microwave Anisotropie Probe (WMAP) \\cite{Spergel}. A  recent combined analysis is presented in \\cite{SDSSWMAP}.\\\\ Recently the detection of the late  Integrated Sachs-Wolfe (ISW) effect has reinforced the case for DE \\cite{Scranton}. All this is most frequently interpreted in the framework of the so-called \"concordance Cosmology\" \\cite{Bahcall}. \\m  A fundamental problem is the identification of the underlying nature of DE : cosmological constant, quintessence (for a review see \\cite{Peebles}) or something else. The most common way is to measure its equation of state (EoS) defined as $w \\eg p_X/\\rho_X\\,$ where $p_X\\,$ is the pressure and $\\rho_X\\,$ the energy density of the DE. The ratio of the DE density to the critical density will be denoted $\\Omega_X$ in a general model and $\\Omega_{\\Lambda}$ in the simplest case of a Cosmological Constant ($w = -1$). $\\Omega_M$ is the corresponding parameter for (baryonic+cold dark) matter.  An ambitious SN program  is now on the way. Important pieces of information will be provided by large ground-based surveys such as the Supernova Legacy Survey (SNLS) \\cite{SNLS} and these investigations will culminate with a space mission as the Supernova Acceleration Probe (SNAP) instrument, part of the Joint Dark Energy Mission (JDEM), which aims at the discovery and follow-up of some 2000 SNIa per year in the redshift range $z = 0.2-1.7$ with very precise magnitude measurements \\cite{SNAPweb}. Of course, the validity of the obtained precision will not depend on the size of the SN sample only, but also on the ability to control the systematic uncertainties at the same level \\cite{Kim}. These systematic errors are coming on one side from the instrument itself (calibration, etc, ...) and also from the SN astrophysical environment (evolution, lensing effects, ...). The SNAP collaboration proposes a global strategy to control such effects at the percent level \\cite{SNAPweb}. \\m In this context, it is necessary to analyze very carefully at which precision level it will be possible to draw any conclusions  from the expected rich amount of data.  Authors have used the present SN data together with simulated sets to  : \\begin{itemize} \\item evaluate the accuracy on the $\\Omega_{\\Lambda}$ parameter to validate the acceleration. \\item measure the DE through the equation of state with $w$ constant in redshift. \\item evidence a possible redshift dependence $w(z)$ with various parametrizations. \\item look at the model dependences and degeneracies and study the possible strategies to break the latter. \\end{itemize}  There is a large consensus that the future SN data {\\em alone} will have difficulties to constrain an evolving equation of state and to break model degeneracies. It is mandatory to have a prior knowledge of the values of some parameters. In particular, a precise knowledge of $\\Omega_M$ will be essential if one hopes to pick out the $z$ dependence of $w$ (see e.g. \\cite{Maor1,Maor2,WA,HT,Goliath,Gerke,LH}) even in the simplest flat cosmology. However, to use some pieces of information from other sources than SN, it will be essential to combine the data in a coherent way, that is under the same hypothesis, in particular on the DE properties. \\m On the other hand, some potential difficulties have been pointed out by various authors. Indeed, most papers have been mainly interested in predicting, for each cosmological parameter, the errors around some fiducial value, close to the ones obtained from the present data as $\\Omega_M = 0.3, \\Omega_X = 0.7, w = -1, dw/dz = 0$. In this procedure, the framework is a particular fiducial model and the chosen fitting function is usually  the one used to generate the data. This strategy is valuable for a first estimate but is too restrictive to pin down the underlying physics which is so far unknown.  For example, Maor et al. \\cite{Maor2} and Gerke and Efstathiou \\cite{Gerke} have stressed the problems which arise if the redshift dependence of the equation of state is neglected whereas this dependence is present. The multiple integral relation between the luminosity distance and $w$ smears out information about $w$ and its time-variation. Assuming in the fitting function that $w$ is a constant, whereas it is not in the fiducial model, can lead to a very wrong estimate of the \"constant\" or even \"effective\" $w$ value. At the same time, the central value of  $\\Omega_M$ (or $\\Omega_X$) is badly reconstructed. \\\\  It is unavoidable to get some ambiguities when trying to fit a particular fiducial cosmology with the \"wrong\" model ! In the following, we will call this problem the \"bias problem\". For present SN data, as will be shown, this question is not a concern as the statistical errors on model parameters stay large due to the limited statistics. In the perspective of much richer SN samples, this bias problem cannot be ignored. We address the problem in a more quantitative way than the authors of \\cite{DiPietro} which have only considered some specific classes of models. In this paper, we investigate the full parameter range, in a model independent way, to quantify where the 'bias' creates potential difficulties and where it can be ignored. We also evaluate the impact of priors when only SN measurements are used. We  get a handle on the possible bias introduced in the analysis by looking at the different hypothesis which are done at the fitting function level. Results are presented in the parameter space of the fiducial models we have simulated.  \\m This paper is organized as follows. In Section 2 we recall the theoretical description of  the magnitude-redshift diagram for supernovae and we present the experimental situation. We identify the relevant parameters which enter as the parameters of the fiducial models and we define the bias which could be introduced by the fitting procedure.  In section 3, we focus on the extraction of the EoS parameter $w_0$. We illustrate the confusion created by the bias on examples. Then we analyze the impact of bias on simulated data : we consider future statistical samples corresponding to SNAP, SNLS and also a sample corresponding to present data. We try to find the safest way for extracting an effective $w$ from the data with minimum bias and minimum priors. Our summary and conclusions are given in Section 4. ",
        "conclusions": "In this paper, we have studied the best strategy to extract the cosmological parameters and in particular the dark energy component from the future supernovae data within the $\\Omega_T=1$ paradigm. We have compared results from the present statistics with the ones we can expect from future large ground-based surveys and from a future space experiment like SNAP.\\\\ In order to be as model independent as possible we have looked at a large range of fiducial models by varying the parameters of the equation of state of the dark energy $w_0^F$ and $w_1^F$ where $w^F(z)=w_0^F+w_1^Fz$ \\cite{WA}.  With present statistics or the one expected from a survey like SNLS, precision on $w_1$ is very weak and the physical question concerns the extraction of $w_0$ and $\\Omega_M$. For a SNAP like experiment a good precision on $w_1$ is only possible with a strong prior on $\\Omega_M$.  To avoid the use of strong priors on $\\Omega_M$, we have focused on the extraction of $w_0$ only and neglected the $w_1$ contribution by fixing it in the fitting procedure. We have to face in this case the problem of introducing a bias if the fiducial value of $w_1^F$ is far away from the fixed one. We compare the bias to the statistical error which is expected for each survey.  To study the validity of such a strategy, we have presented an extensive study of fiducial models in a flat universe in which we vary $w_0^F$ and $w_1^F$ in a complete range $-2<w_0^F<0$ and $-2<w_1^F<2$ corresponding to a large variety of DE theoretical models. On the other hand $M_s^F$ and $\\Omega_M^F$ are fixed. In the fit we have used in general a weak prior on $\\Omega_M (\\sigma(\\Omega_M)< 0.1)$  and we have looked at the bias introduced by fixing $w_1=0$. We call this procedure the 3-fit ($M_s, \\Omega_M, w_0$) procedure.   For SNAP (and SNLS) like statistics, we obtain the following conclusions : \\begin{itemize} \\item $w_0$  is not biased for an important part of the scanned plane ($w_0^F, w_1^F$) which corresponds to  models describing a Universe which is today in acceleration. \\item for $M_S$, we get the same conclusion but we emphasize that this parameter is so much constrained in the fitting procedure that a small bias has no visible effect. \\item $\\Omega_M$ is strongly biased if the redshift dependence is large ($|w_1^F|>0.2$(0.6)). \\item If a strong prior on $\\Omega_M$ is used, $w_0$ is biased for most of the plane. Namely, the 3-fit procedure is relevant to extract $w_0$ if only a weak $\\Omega_M$ prior is used. \\end{itemize}   We have compared the accuracy of the results on the $w_0$ parameter to the ones obtained if $w_1$ is left free as in a 4-fit ($M_s, \\Omega_M, w_0,w_1$) procedure.  We found that, in spite of the bias, the 3-fit procedure, where $w_1$ is fixed, can be used for SNAP and SNLS to test accelerating models and give better or equivalent results than the simplest 4-fit where $w_1$ is completely free. This is no longer true if $\\mid w_1^F\\mid$ is large. For decelerating models the 4-fit is mandatory. Note that, concerning the choice of the number of parameters to be fitted, the bayesian method presented in \\cite{Bridle} seems promising.\\\\  These conclusions have been set using only the value $\\Omega_M^F=0.3$ for the fiducial models. If the true value is different the areas of the different Validity Zones increase if the contribution of DE is smaller, namely if $\\Omega_M^F$ is larger. Conversely, if $\\Omega_M^F<0.3$, the validity of the 3-fit is reduced.\\\\  Using the presently available statistics of the SCP collaboration \\cite{newSCP}, it appears that the procedure where $w_1$ is fixed can be trusted for almost all the fiducial models considered here. Only fiducial models with strong recent variations of the EoS ($\\mid w_1^F \\mid>0.7$) may lead to biased results.\\\\  Let us point out that other bias would be introduced, by a departure from the $\\Omega_T=1$ paradigm in the fiducial model, or if the central value of the $\\Omega_M$ prior itself (used in the fitting procedure) is not the same as the fiducial value. Indeed a small bias in the prior could induce some strong bias for the other cosmological parameters. These bias will be studied in a future paper \\cite{enpreparation}. \\\\   It has been advocated recently \\cite{Wang} that it should be easier to constrain the DE density $\\rho_X(z)$ and its time derivative $\\rho^{\\,\\prime}_X(z)$, instead of constraining the equation of state $w(z)$.  These authors have simulated a large variety of fiducial models in the range $-1.2 \\leq w^F_0 \\leq -0.5$ and $-1.5 \\leq w^F_1 \\leq 0.5$, which is almost completely contained in our Validity Zone for SNAP (see Fig.~\\ref{fig:statplussyst}). They interpret physically parts of this plane in term of increasing, decreasing or non-monotonic DE densities $\\rho_X(z)$.  Their approach is interesting to answer if the DE energy density is a constant or not. Nevertheless, focusing on the equation of state itself remains mandatory if, for instance, one wants to validate precisely the acceleration since, in this case, a precise knowledge of $w_0$ (and also on $\\Omega_M$ and $\\Omega_X$) is necessary. Finally, we would like to point out that trying to answer any particular question on the nature of DE requires an investigation on the choice of the most relevant fitting procedures.\\\\        \\m  \\no {\\bf Acknowledgments}\\\\   We wish to thank A. Mazure and the members of the Laboratoire d'Astrophysique de Marseille, for their support, and also E. Linder, R. Miquel and F. Steiner, for helpful discussions. ``Centre de Physique Th\\'eorique'' is UMR 6207 - ``Unit\\'e Mixte de Recherche'' of CNRS and of the Universities ``de Provence'', ``de la M\\'editerran\\'ee'' and ``du Sud Toulon-Var''- Laboratory affiliated to FRUMAM (FR 2291). ``Centre de Physique des Particules de Marseille'' is ``Unit\\'e Mixte de Recherche'' of CNRS/IN2P3 and of the University ``de la M\\'editerran\\'ee''.\\\\"
    },
    "0403/astro-ph0403550_arXiv.txt": {
        "abstract": "{ The upper bound on the value of the surface gravity, $g_{\\rm s}$, for neutron stars with equations of state respecting $v_{\\rm sound}\\le c$, is derived. This bound is inversely proportional to the  maximum allowable mass $M_{\\rm max}$, and it reads $g_{\\rm s}\\le 1.411 \\times 10^{15}\\;({M_{\\rm max} /{ M}_\\odot})^{-1}~{\\rm cm~s^{-2}}$. It  implies an {\\it absolute upper bound} $7.4\\times 10^{14}~{\\rm cm~ s^{-2}}$ if one uses the $2\\sigma$ lower bound on the neutron mass measured recently in 4U1700-37, $1.9 M_{\\odot}$. A correlation between $g_{\\rm s}$ and the compactness parameter $2GM/Rc^2$ for baryonic stars is analyzed. The properties of $g_{\\rm s}$  of strange quark stars and its upper bounds are discussed using the scaling properties of the strange-star models. ",
        "introduction": "\\label{sect:introduction} Gravitational acceleration on the stellar surface, usually called the surface gravity, is an important parameter of the theory of stellar atmospheres, in particular the atmospheres of neutron stars (see, e.g., Zavlin \\& Pavlov 2002). For  neutron stars, $g_{\\rm s}$ should be defined taking into account the space-time curvature. The value of $g_{\\rm s}\\sim {\\rm few}\\times 10^{14}~{\\rm cm~s^{-2}}$  is involved the relation between internal and surface temperature of neutron stars (Potekhin et al. 2003 and references therein).  The surface gravity of neutron stars is by many orders of magnitude larger than for other stars; it is $\\sim 10^5$ times larger than for white dwarfs, and $10^8$ times stronger than the gravity at the solar surface. As we show in the present paper, at a given neutron-star mass the value of $g_{\\rm s}$ depends very strongly on the largely unknown equation of state (EOS) of dense matter at supra-nuclear densities. It is therefore of interest to derive upper bounds on $g_{\\rm s}$. In Sect.\\ \\ref{sect:gs.bounds} we derive such upper bounds on $g_{\\rm s}$ resulting from the condition of subluminality of the EOS (speed of sound not exceeding $c$). We show how an upper bound on $g_{\\rm s}$ can be obtained from measured neutron star masses. In Sect.\\ \\ref{sect:gs.EOS} we analyze the EOS-dependence of  $g_{\\rm s}$ for a set of 31 models of dense matter and compare the maximal surface gravity, reached at the maximum allowable mass, with subluminal upper bounds. In Sect.\\ \\ref{sect:gs.SS.scaling} we consider surface gravity of strange quark stars. We derive a scaling formula for $g_{\\rm s}$ of strange stars and use it to relate the maximum surface gravity for various models of such stars. Concluding remarks are  presented in  Sect.\\ \\ref{sect:conclusion}. ",
        "conclusions": "\\label{sect:conclusion} We derived an upper bound on the surface gravity of static neutron stars with subluminal EOSs.  Maximum $g_{\\rm s}$ is reached at the maximum allowable mass $M_{\\rm max}$.  Even at the shortest observed pulsar period of $1.56$ ms the  effect of rotation on the $M_{\\rm max}$ is very small, so that the static approximation is justified. The upper bound on $g_{\\rm s}$ is inversely proportional to the highest measured neutron-star mass $M_{\\rm obs}^{\\rm max}$ and is $7.4\\times 10^{14}~{\\rm cm~s^{-2}}$ if $M_{\\rm obs}^{\\rm max}$ is replaced by $1.9~{\\rm M}_\\odot $, which is the  $95\\%$ confidence-level lower bound on mass of the compact star in 4U1700-37 measured by Clark et al. (2002). However, this upper bound on $g_{\\rm s}$ should be taken as highly unreliable. Firstly, Clark et al. (2002) used a specific model of the companion star and of the binary to evaluate  the compact-star mass. Secondly,  one cannot exclude, unfortunately,  that the compact object in 4U1700-37 is actually a low-mass black hole and not a neutron star (Clark et al. 2002 give arguments for and against the black-hole presence).  We have studied $g_{\\rm s}$ for a set of 31 EOSs of baryonic matter. The dependence of $g_{\\rm s}$ on the stellar mass $M$ is very sensitive to the EOS. On the contrary, the dependence of $g_{\\rm s}$ on the stellar compactness $r_{\\rm g}/R$ has a generic character for baryonic EOSs. The maximum surface gravity $g_{\\rm s,max}$ is sensitive to the EOS of dense matter and ranges from about $2\\times 10^{14}~{\\rm cm~s^{-2}}$ to $5\\times 10^{14}~{\\rm cm~s^{-2}}$ for subluminal baryonic EOSs.  The dependence of $g_{\\rm s}$ on the mass and compactness of strange stars is very different from that of baryonic stars. However, the range of $g_{\\rm s,max}=(3-6)\\times 10^{14}~{\\rm cm~s^{-2}}$ is quite similar to baryonic stars.  We hope that the EOS-sensitive features of $g_{\\rm s}$ will be useful in extracting information about the EOS of dense matter, for instance, by combining the values  of $g_{\\rm s}$ obtained  fitting the thermal component of the observed photon spectra with atmospheric models and the surface redshift measured for the identified spectral lines.  \\acknowledgement{This work was motivated by a question of  A. Majczyna concerning the possible range of surface gravity in neutron star atmospheres, which is gratefully acknowledged. We are grateful to D.G. Yakovlev for the reading  of the manuscript  and for helpful remarks. This work was partially supported by the KBN grants no.  2P03D.019.24 and  5P03D.020.20. }"
    },
    "0403/astro-ph0403236_arXiv.txt": {
        "abstract": "{Motivated by the first detections of rotation signatures in the DG~Tau jet (Bacciotti et al. \\cite{bacciotti2002}), we examine possible biases affecting the relation between detected rotation signatures and true azimuthal velocity for self-similar MHD disc winds, taking into account projection, convolution as well as excitation gradients effects. We find that computed velocity shifts are systematically smaller than the true underlying rotation curve.  When outer slower streamlines dominate the emission, we predict observed shifts increasing with transverse distance to the jet axis, opposite to the true rotation profile. Determination of the full transverse rotation profile thus requires high angular resolution observations ($< 5$~AU) on an object with dominant inner faster streamlines. Comparison of our predictions with HST/STIS observations of DG~Tau clearly shows that self-similar, \\textit{warm} MHD disc wind models with $\\lambda = 13$ and an outer radius of the disc $\\simeq 3$~AU are able to reproduce detected velocity shifts, while \\textit{cold} disc wind models ($\\lambda > 50$) are ruled out for the medium-velocity component in the DG~Tau jet. ",
        "introduction": "\\label{intro}  It is now widely accepted that a large scale magnetic field is responsible for both the acceleration and collimation of jets around young accreting stars. However, the exact launching zone (stellar surface, disc truncation radius or wide range of disc radii) remains subject to debate. So far, studies aimed at constraining proposed ejection models have concentrated on the jet collimation, poloidal velocities, and excitation conditions (Dougados et al. \\cite{dougados2000}, Lavalley-Fouquet et al. \\cite{lavalley2000}, Bacciotti et al. \\cite{bacciotti2000}, Woitas et al. \\cite{woitas2002}).  The detection of rotation signatures in the DG~Tau jet has recently opened new prospects to constrain MHD ejection models (Bacciotti et al. \\cite{bacciotti2002}). In particular, approximate launching radii $\\simeq 0.3$--$3$~AU were inferred in DG~Tau, indicative of a disc wind (Anderson et al. \\cite{anderson2003}). In this Letter, we investigate two classes of self-similar disc wind models (Sect. 2) and analyse in detail the biases introduced in their observed rotation signatures by projection, beam dilution and emissivity gradients within the jet (Sect. 3).  Implications for a proper interpretation of the DG~Tau observations are discussed in Sect. 4. ",
        "conclusions": ""
    },
    "0403/astro-ph0403693_arXiv.txt": {
        "abstract": "The spectrum of the hard X-ray background records the history of accretion processes integrated over the cosmic time. Several pieces of observational and theoretical evidence indicate that a significant fraction of the energy density is obscured by large columns of gas and dust. The absorbing matter is often very thick, with column densities exceeding $N_H\\simeq 1.5\\times 10^{24}$~cm$^{-2}$, the value corresponding to unity optical depth for Compton scattering. These sources are called ``Compton thick'' and appear to be very numerous, at least in the nearby universe. Although Compton thick Active Galactic Nuclei (AGN) are thought to provide an important contribution to the overall cosmic energy budget, their space density and cosmological evolution are poorly known. The properties of Compton thick AGN are reviewed here, with particular emphasis on their contributions to the extragalactic background light in the hard X-ray and infrared bands. ",
        "introduction": "Most of the active galactic nuclei (AGN) in the local universe are obscured in the X-ray band by large amounts of gas and dust, which prevent the observation of their nuclear emission up to energies that depend on the amount of intrinsic absorption. If the X-ray obscuring matter has a column density which is equal to or larger than the inverse of the Thomson cross-section ($N_H\\ge \\sigma_T^{-1} \\simeq 1.5 \\times 10^{24}$~cm$^{-2}$), then the source is called, by definition, ``Compton thick''. The cross-sections for Compton scattering and photoelectric absorption have approximately the same value for energies of order 10~keV, which can be considered as the low energy threshold for probing the Compton thick absorption regime. Indeed, if the column density does not exceed a value of order 10$^{25}$~cm$^{-2}$, then the nuclear radiation is visible above 10~keV, and the source is called mildly Compton thick. For higher column densities (heavily Compton thick), the entire high energy spectrum is down-scattered by Compton recoil and hence depressed over the entire X-ray energy range. The presence of Compton thick matter may be inferred through indirect arguments, such as the presence of a strong iron K$\\alpha$ line complex at $6.4-7$~keV and the characteristic reflection spectrum.  The study of Compton thick sources is relevant for several reasons: (1) there is observational evidence that suggests that a large fraction of AGN in the local universe are obscured by Compton thick gas (Maiolino et al.\\ 1998; Risaliti, Maiolino, \\& Salvati 1999a; Matt et al.\\ 2000); (2) a sizable population of mildly Compton thick sources is postulated in all the AGN synthesis models for the X-ray background (XRB) in order to match the intensity peak of the XRB spectrum at about 30~keV. The absorbed luminosity eventually will be reemitted in the far-infrared (far-IR), making Compton thick sources potential contributors to the long wavelength background. Finally, accretion in the Compton thick AGN may contribute to the local black hole mass density.  Unfortunately, the most efficient energy range to search for mildly Compton thick sources is just above the highest energy accessible to the past and present generation of satellites with imaging capabilities for faint limiting fluxes. As a consequence, the search for Compton thick sources has been limited, so far, to the relatively bright fluxes accessible to the high energy detectors onboard {\\em BeppoSAX} and {\\em RXTE}.  In this chapter, we review the evidence for obscured AGN, with a special emphasis on the evidence for Compton thick sources. We then examine in some detail the contributions of Compton thick sources to the XRB. ",
        "conclusions": "Different sets of X-ray observations and theoretical arguments indicate that the reprocessing of high energy radiation by Compton thick matter in the circumnuclear regions of AGN is very common, if not ubiquitous. Unfortunately, even the most sensitive hard X-ray surveys carried out with the {\\em Chandra} and {\\em XMM-Newton} observatories turn out to be rather inefficient for searching for the typical signatures of thick matter. Further observations, though useful in several respects, will not provide the breakthrough in this direction. In fact, according to what can be considered a solid model for the XRB, only a small fraction of the Compton thick population could be detected by present X-ray surveys below $8-10$~keV.  A mission capable of exploring the hard X-ray sky in the $10-70$~keV band with focusing/imaging instruments able to reach fluxes of order $10^{-14}$~ergs~cm$^{-2}$~s$^{-1}$ would provide a quantum leap forward. As an order of magnitude estimate, the three dex jump in the limiting flux would be similar to that achieved by {\\em BeppoSAX} and {\\em ASCA} in the $2-10$~keV band with respect to the {\\em HEAO1-A2} survey of Piccinotti et al.\\ (1982). In other words, such a gain in sensitivity would push the resolved fraction of the hard ($10-70$~keV) XRB from a negligible 1\\% to about $30-40$\\%. Sensitive, high energy surveys will open up a large volume of discovery space that so far only has been touched by {\\em BeppoSAX} observations. It is important to stress that the scientific output of these observations would be even more rewarding than a mere improvement in the limiting flux, as the $10-70$~keV band fully encompasses the peak of the XRB energy density at 30~keV.  The X-ray Evolving Universe Spectrometer {\\em (XEUS)} mission, currently under study to be the next ESA cornerstone mission for X-ray astronomy, will be capable of surveying the hard X-ray sky at, or even below, $10^{-14}$~ergs~cm$^{-2}$~s$^{-1}$ to energies of $40-50$~keV. Obviously, the scientific objectives of such a mission would be much broader than a better sampling of the Compton thick population. The mission would be devoted to probing a significant fraction of the electromagnetic radiation from accretion onto supermassive black holes."
    },
    "0403/astro-ph0403370_arXiv.txt": {
        "abstract": "{We present an analysis of CCD photometric observations of the eclipsing novalike cataclysmic variable \\object{DW UMa} obtained in two different luminosity states: high and intermediate. The star presents eclipses with very different depth: $\\sim1.2$ mag in the high and $\\sim3.4$ mag in the intermediate state. Eclipse mapping reveals that this difference is almost entirely due to the changes in the accretion disc radius: from $\\sim0.5R_{\\rm L_1}$ in the intermediate state to $\\sim0.75R_{\\rm L_1}$ in the high state ($R_{\\rm L_1}$ is the distance from the white dwarf to the first Lagrangian point). In the intermediate state, the entire disc is eclipsed while in the high state, its outer part remains visible. We also find that the central intensity of the disc is nearly the same in the two luminosity states and that it is the increase of the disc radius that is responsible for the final rise from the 1999/2000 low state. We find that the intensity profile of the disc is rather flat and suggest a possible explanation. We also discuss the effect of using a more realistic limb-darkening law on the disc temperatures inferred from eclipse mapping experiments. Periodogram analysis of the high state data reveals \"positive superhumps\" with a period of $0\\fd1455$ in 2002 and $0\\fd1461$ in 2003, in accord with the results of Patterson et al. However, we cannot confirm the quasi-periodic oscillations reported by these authors. We obtain an updated orbital ephemeris of \\object{DW UMa}: $T_{\\rm min}[HJD]=2446229.00687(9)+0\\fd136606527(3)E$. ",
        "introduction": "\\object{DW Ursae Majoris} is an eclipsing nova-like (NL) cataclysmic variable (CV) with an orbital period of $\\sim3.27$ hours. The star is one of the NLs which show deep ($\\sim4$ mag) low states (Hessman \\cite{hess1}; Honeycutt et al. \\cite{hon, hon1}), and \"negative\" and \"positive superhumps\" (Patterson et al. \\cite{patt02}). Shafter et al. (\\cite{shaf1}) carried out the first detailed photometric and spectroscopic study of \\object{DW UMa} in a high state. The photometry revealed $\\sim1.5$ mag deep eclipses, with a peculiar V-like shape. The spectral observations showed single-peaked H, \\ion{He}{i} and \\ion{He}{ii} emission lines with superimposed transient narrow absorption components. These peculiarities led Thorstensen et al. (\\cite{th91}) to group \\object{DW UMa} along with \\object{SW Sex}, \\object{V1315 Aql} and \\object{PX And}, and to define so-called SW Sex stars.  Biro (\\cite{biro}) presented $V$ and $R$ eclipse maps of \\object{DW UMa} in a high state. The eclipse maps show that the radial profile of the accretion disc (AD) temperature distribution is rather flat, similarly to with what is found in other \\object{SW Sex} novalikes (Rutten et al. \\cite{rutt}; Baptista et al. \\cite{bsh}). Biro (\\cite{biro}) also estimated the distance to the system to be $270\\,\\pm50$ pc.  Spectral observations of Dhillon et al. (\\cite{dhil}) in low state revealed narrow Balmer emission lines originating from the irradiated face of the secondary star. Marsh \\& Dhillon (\\cite{mdh}) could not detect the secondary star in their $I$-band spectra obtained in low state and derived a lower limit of the distance 450 pc or 850 pc depending on the assumed spectral type of the secondary. Ultraviolet spectroscopy during a recent low state was presented by Knigge et al. (\\cite{knigge}). Surprisingly, they found that during the low state the continuum shortward of 1450 \\AA\\ was higher and bluer than in the high state. This led Knigge et al. (\\cite{knigge}) to suggest that in the high state of \\object{DW UMa} the white dwarf (WD) is permanently hidden from our sight by the rim of a flared accretion disc. Knigge et al. (\\cite{knigge}) also modeled the ultraviolet spectrum of the WD and estimated the distance to \\object{DW UMa} to be $830\\,\\pm150$ pc.  \\begin{figure*}[t] \\includegraphics*[width=18cm]{ms0145f1.eps} \\caption{Photometric observations of \\object{DW UMa}. Note the much larger range of the $y$-axis in the upper row panel.} \\label{pats} \\end{figure*}    Using the same observations as Knigge et al. (\\cite{knigge}), Araujo-Betancor et al. (\\cite{syspar}) could measure the eclipse phases of the WD. This allowed them to determine the system parameters of \\object{DW UMa} with very good accuracy. Repeating the analysis of Knigge et al. (\\cite{knigge}), and using $I$- and $K$-band photometry, Araujo-Betancor et al. (\\cite{syspar}) obtained two estimates of the distance to \\object{DW UMa}: $930\\,\\pm150$ pc and $590\\,\\pm100$ pc, respectively.    In this article we present  photometric observations of \\object{DW UMa} in two different brightness states -- intermediate and high. We performed periodogram analysis in the high state and detected \"positive superhumps\". We also present eclipse maps of \\object{DW UMa}'s accretion disc in the intermediate and high states. In the last part of the article we discuss the results and present our conclusions. ",
        "conclusions": "Perhaps the most intriguing results of this study are the AD radius variations and the nearly equal central intensities of the high and intermediate state maps. As high and low states are caused by strong variations of mass accretion rate, the later is difficult to explain within the standard accretion disc theory. However, what we formally call \"intermediate state\" is in fact a transition from the 1999/2000 low state to the high state and during this transition the disc might not be in equilibrium with the actual accretion rate. However, this is most likely not the case. Figure\\,\\ref{lt}b shows that the recovery from the low state is rather slow and takes $\\sim4$ months, i.e. the disc has enough time to adjust its structure to changes of mass accretion rate. This suggests that at given value of the accretion rate some kind of saturation occurs and further increase of the accretion rate does not lead to increase of the emission from the inner part of the disc. Instead, the rise of the system luminosity is due to an increase of the disc radius from $\\sim0.5R_{\\rm L_1}$ to $\\sim0.75R_{\\rm L_1}$\\footnote{We stress that the fact that we see this behavior on the rising branch of the low state does not necessarily mean that one should expect the opposite on the falling branch.}. These values are in accord with the AD radii in NLs, including \\object{DW UMa} itself, determined by Harrop-Allin \\& Warner (\\cite{haw}) from high- and low-state observations. Note that the ratio of the area of circles with these radii is $\\sim2.25$. If the AD brightness distribution does not depend on radius, this will give $\\sim1$ mag increase of the brightness, as observed. The reason of this behavior is not clear. It could be that at certain value of the accretion rate some mechanism starts removing energy from the inner disc, thus preventing further increase of the emission of this part. Thus, it appears that this issue could be tightly related to the mechanism responsible for the flat intensity profile of the disc, which we discuss next.    The flat AD intensity profile found in \\object{SW Sex} type NLs is still a puzzle. This problem has been extensively discussed in the literature, but there is no widely accepted model. Currently, two of the proposed models seems to be most viable. Accretion disc wind (or something else) can remove energy from the inner part of the disc thus causing the profile to be flatter than the steady-state $T_{\\rm eff}\\propto r^{-3/4}$ dependence. The second model assumes that \\object{SW Sex} stars are intermediate polars and that the inner part of their discs is missing. In order to account for the other properties of the \\object{SW Sex} stars it is also assumed that after the first impact with the AD's edge, part of the gas in the accretion stream continues moving above the disc and hits its surface again close to the WD. In the intermediate polar model the stream hits the WD's magnetosphere (for a review see Hellier \\cite{helrev}). In this later model the flat profiles found in the eclipse mapping experiments are simply a result of the assumption that the AD extends down to the WD's surface. The later model has recently gained additional support. Rodriguez-Gill et al. (\\cite{gil1, gil2}) reported the discovery of periodic circular polarization in two \\object{SW Sex} stars. Patterson et al. (\\cite{patt02}) reported possible detections of periodic signals ($\\sim1000$ s) in the light curves of several \\object{SW Sex} and suggested that this could be a manifestation of the rotation of a highly magnetic WD. Note, however, that in the intermediate polars the radius of the WD's magnetosphere or equivalently the inner radius of the AD $r_{\\rm in}$ depends on the mass accretion rate $\\dot{M}$: $r_{\\rm in}\\propto\\dot{M}^{-2/7}$. Thus, the higher $\\dot{M}$ the lower $r_{\\rm in}$. Therefore, it can be expected that in the high-state eclipse mapping analysis the AD brightness distribution will appear closer to the steady-state law than that in intermediate state. We do not see this behavior and the AD brightness distribution is rather flat in both state eclipse maps.   Recently, Knigge et al. (\\cite{knigge}) found that during the 1999 low state, the UV continuum of \\object{DW UMa} shortward of 1450 \\AA\\ was higher and bluer than in the high state.  The authors suggest that in the high state the WD and the inner hot part of the disc are permanently hidden from our sight by the rim of a flared AD. This certainly can alter the disc intensity profile of \\object{DW UMa} and cause it to appear flat in the eclipse maps. However, the weakness of this model is the need of a very high AD rim; its height should be at least $\\sim0.15R_{\\rm disc}$ if the system inclination is $82^\\circ$. With the present knowledge, it is difficult to understand how such a high rim can be formed. Moreover, it should have a constant height because Knigge et al. (\\cite{knigge}) found that in the low state the UV spectrum is {\\it at all} orbital phases bluer and the flux higher than in high state. In the classical steady-state AD model this would require unrealistically high mass accretion rate. The impact of the accretion stream with the disc can form a thick bulge but it could not persist along the whole disc rim.  Here, we suggest another possibility. Bobinger et al. (\\cite{bob}) found that the intensity profile of \\object{IP Peg}'s AD in decline from outburst is flat. The authors explored many possible causes of this and concluded that the presence of an optically thick and geometrically flat layer on top of the AD's surface masks its true intensity profile. Bobinger et al. (\\cite{bob}) suggested that this layer may be formed by the wind from the disc. This is even more relevant in NLs, which generally have hotter discs than dwarf novae.  Knigge \\& Drew (\\cite{knigge97}) modeled the wind in the NL \\object{UX UMa} and found that between the disc's photosphere and fast-moving part of the wind there is a transition zone of a slowly outflowing and relatively dense layer. We suggest that such a layer may also be present in \\object{DW UMa}, and may shade the WD and the inner hot part of the disc, or what is more likely, attenuate the radiation of the inner disc and the WD. Although this may seem somewhat speculative, it can account for (1) the flat disc profile; (2) the red UV spectrum in high state and (3) nearly the same temperature of the inner disc in the high and intermediate states. However, it should be borne in mind that in \\object{UX UMa} itself this layer does not flatten the AD's temperature profile; \\object{UX UMa} is often given as an example where the results of eclipse mapping are consistent with the steady-state $T_{\\rm eff}\\propto r^{-3/4}$ law. The reason for the difference may be different parameters of the wind.    \\begin{figure*}[t] \\includegraphics*[width=12cm]{ms0145f7.eps} \\caption{Radial profile of the AD brightness temperature obtained  by using (i) temperature dependent limb-darkening coefficient -- symbols connected with solid lines and (ii) with constant limb-darkening coefficient of 0.6 -- symbols connected with dash-dotted lines. The dotted curves show the effective temperature of steady-state ADs for mass transfer rates $\\dot{M}=10^{-10}-10^{-7}\\,M_{\\sun}$\\,yr$^{-1}$ onto 0.77$M_{\\sun}$ WD.} \\label{rad} \\end{figure*}   Having the eclipse map calculated one can estimate the brightness temperature $T_{\\rm BR}$ of the accretion disc assuming blackbody emission. This is done for the $V$-band eclipse maps by solving for $T_{\\rm BR}$ the following equation: \\begin{equation} f_j=\\frac{C\\delta_j\\cos i}{D^2}\\,\\frac{\\int v(\\lambda)B_\\lambda(T_{BR,j})\\lambda\\,d\\lambda} {\\int v(\\lambda)\\lambda\\,d\\lambda}. \\label{temp} \\end{equation} Here $f_j$ is the flux emitted by the $j$-th element of the AD, $T_{\\rm BR,j}$ and $\\delta_j$ are its brightness temperature and area, respectively, $B_\\lambda(T_{\\rm BR,j})$ is the blackbody spectral distribution, $v(\\lambda)$ is the response of the $V$ bandpass and $D$ is the distance to the system. As noted by Baptista et al. (\\cite{bap98}) the relation between the brightness temperature and the effective temperature is not trivial and to obtain such a relation the vertical structure of the disc has to be known. However, the emission of the hot optically thick ADs in novalikes probably closely follows the blackbody law  and $T_{\\rm BR}$ should be close to the effective temperature. Optically thick ADs are however subject to an effect analogous to the limb-darkening in stars and the factor $C$ takes this into account. This later effect greatly complicates the situation because it is well known that 1) the real limb-darkening differs significantly from often used linear law and 2) $C$ is a function of surface gravity, temperature and wavelength (Wade \\& Rucinski \\cite{ld}; Diaz et al. \\cite{diaz}).   To investigate the importance of limb-darkening further, we have calculated $C$ by using the available synthetic stellar spectra with solar abundances and $\\log g=5$ calculated by R.L. Kurucz\\footnote{available at {\\tt http://cfaku5.harvard.edu}}. We calculated $C$ for $V$-band wavelengths, $i=82^\\circ$ and temperatures from 5000 to 50\\,000\\,K. Not surprisingly $C$ turns out to depend strongly on the temperature (non-linearly) and changes from $\\sim0.35$ at $T_{\\rm eff}=5000\\,K$ to $\\sim0.84$ at $T_{\\rm eff}=50\\,000\\,K$. Thus, it seems that the dependence of limb-darkening on the temperature can affect significantly the estimation of AD brightness temperature obtained from eclipse mapping. This dependence would also affect the obtained radial profile of the temperature in ADs with {\\it temperature varying in a wide range}. This would have a direct consequence on the inferred accretion rates and distances to the CVs studied by eclipse mapping.   Although ADs structure is different from that of the normal stars, Diaz et al. (\\cite{diaz}) have shown that the difference between limb-darkening laws obtained with Kurucz stellar spectra and AD model calculations is small (see Fig.\\,1 in Diaz et al. \\cite{diaz}). Thus, the main source of error when calculating $C$ is the assumption of constant surface gravity $\\log g=5$ ($\\log g=5$ is the maximal gravity in Kurucz' models). In real discs the gravity changes from $\\log g\\sim4.5$ at the outer edge to $\\log g\\sim6.5-7.0$ at the inner parts. The work of Wade \\& Rucinski (\\cite{ld}) shows that in stars for given effective temperature the limb-darkening coefficient decreases with increasing $\\log g$, which means $C$ increases with $\\log g$. If this holds true for ADs and $\\log g>5$, then the temperatures one obtains with $\\log g=5$ overestimate the real temperature of the AD, particularly in its inner part, where $\\log g>5$; the degree of overestimation is however difficult to estimate without model spectra with $\\log g>5$.  Since the distance to \\object{DW UMa} is rather uncertain (see the Introduction for details) we have solved Eq.\\,\\ref{temp} for three distances: 300, 600 and 900 pc. Figure\\,\\ref{rad} shows the results: symbols connected with solid lines show solutions with limb-darkening as determined from Kurucz' model spectra, and those connected with dash-dotted lines show the brightness temperature obtained using a linear limb-darkening law with a coefficient of 0.6. Also shown with dotted curves is the expected effective temperature radial distribution of steady-state ADs for mass transfer rates $\\dot{M}=10^{-10}-10^{-7}\\,M_{\\sun}$\\,yr$^{-1}$ onto $0.77M_{\\sun}$ WD (Araujo-Betancor et al. \\cite{syspar}). It is clearly seen that the temperature of the AD in \\object{DW UMa} in both high and intermediate state is very far from following the steady-state radial dependence: it is much flatter that expected. This finding is in accord with the results of Biro (\\cite{biro}). Also seen is that the usage of more realistic limb-darkening law results in lower AD temperatures. The effect on the radial dependence is however rather small, which is a direct consequence from the flat intensity profile.  \\begin{figure*}[t] \\includegraphics*[width=18cm]{ms0145f8.eps} \\caption{Reconstruction of the radial $V$-band intensity profile of simulated eclipses with S/N ratio of 100. Dots -- the reconstruction, solid line -- the disc profile used to simulate the eclipse. {\\bf a--c} -- flat discs with different radial dependence of the temperature, {\\bf d--f} 3D steady-state discs with emitting outer rims of different height (see text for details). In all reconstructions the eclipses were fitted to reduced $\\chi^2=1$ using a flat 1D disc geometry and full azimuthal streaming default images.} \\label{simul} \\end{figure*}    Although the recent studies point to a rather large distance to \\object{DW UMa}, $\\geq500$ pc, we note that this seems inconsistent with the results of eclipse mapping (Biro \\cite{biro}; this study). The main source of concern is the relatively high temperature ($\\sim20\\,000$\\,K) at the outer disc regions obtained when distance larger than 500 pc is assumed. To achieve this, a very high mass accretion rate would be needed. It seems not very likely that the reason for this is the usage of a simplified limb-darkening law. The cause of the discrepancy is a puzzle. Although the dependence of limb-darkening on both the temperature and surface gravity is in favor of obtaining flatter radial temperature profiles (compared to if constant limb-darkening coefficient is used), this could hardly explain the very flat profile in \\object{DW UMa}. The tests of the algorithm over simulated eclipses showed a good recovery of the radial profile. Results of some of the simulations of flat discs are shown in Figs.\\,\\ref{simul}{\\it a-c} -- {\\it (a)}  the $V$-band intensity profile of a steady-state accretion disc with $T_{\\rm eff}\\propto r^{-0.75}$; {\\it (b)} $T_{\\rm eff}\\propto r^{-0.375}$ and {\\it (c)} a combination of both. The disc elements are assumed to emits as a black body at given temperature. The emergent radiation is passed through a $V$-band filter, foreshortened and limb-darkened. We used a linear limb-darkening law with a coefficient 0.6. Gaussian noise of S/N=100 was added and so simulated eclipses were fitted to reduced $\\chi^2=1$ using a full azimuthal streaming default image. It is seen that the recovery in the middle part of the disc is good. Figure\\,\\ref{simul} also clearly shows one well-known problem of the ME technique, namely that it tends to smear out any sharp structures in the disc. As a result the outer edge of the disc could not be accurately reconstructed and the radial profile is steeper at the outer part. On the other hand, in the inner disc only the very innermost pixels are smoothed and this is not a serious problem for the recovery of the radial profile. We notice that decreasing the noise level to S/N=300 does not help much to obtain a better reconstruction and the above situation does not improve significantly.  Of more serious concern are the various systematic effects that may become important in the nearly edge-on systems. As already discussed, the mass accretion rate in NLs is probably not high enough for the disc rim to obscure the disc surface if the inclination is $\\sim80^\\circ$. However, if $\\dot{M}\\sim10^{-8}\\,M_{\\sun}$\\,yr$^{-1}$ the disc rim may still be as high as $\\sim0.07-0.1R_{\\rm disc}$\\footnote{This refers to the height above the disc mid-plane. The actual size of the rim is twice larger.}. Recently, Smak (\\cite{smak}) suggested that in this case if the temperature of the rim is comparable to that of the outer disc, in the high-inclination NLs the rim may contribute more than 50\\% of the system light. Clearly, this would affect the eclipse profile significantly and when not accounted for may bias the eclipse mapping results.   Smak (\\cite{smak94}) and Bobinger et al. (\\cite{bob}) have studied some aspects of the reliability of the reconstruction of 3D discs. Of particular interest is the work of Smak (\\cite{smak94}) who tested the reliability of the reconstruction of 3D discs if a simple flat 1D disc geometry is used in the reconstruction. Smak's work is however not based on the ME technique. As a part of the testing of our eclipse mapping code we have performed a similar study. We simulated eclipses in a binary with $i=82^\\circ$, $q=0.39$, $M_1=0.77M_\\odot$ and $\\dot{M}\\sim10^{-8}\\,M_{\\sun}$\\,yr$^{-1}$. A steady-state disc, whose (half)thickness varies as $r^{9/8}$ is assumed. The outer disc rim is modeled as a cylindrical surface with a temperature equal to that of the outermost part of the disc. The disc has a radius of $0.75R_{\\rm L_1}$. We then simulated eclipses for three values of the rim height, 0.07, 0.1 and $0.13R_{\\rm disc}$, which are appropriate for very high mass accretion rates. The $V$-band intensities were computed as in the case of the flat disc simulations\\footnote{Note that in the disc geometry assumed, the angle at  which  given disc element is seen is also a function of the orbital phase. In this case the foreshortening and limb-darkening have to be applied simultaneously and they become part of the visibility function (or projection matrix) applied to the disc elements.}. Gaussian noise of S/N=100 was added and the simulated eclipses were fitted assuming a flat disc geometry.  The reconstructed radial $V$-band intensity profiles are shown in Figs.\\,\\ref{simul}{\\it d-f}. Generally, our results show that the radial profile of a 3D accretion disc may be still well reconstructed assuming a flat geometry, provided that the inclination is not very high. However, the reconstruction clearly worsens with increasing the rim height and in Figs.\\,\\ref{simul}{\\it d-f} one can see the increasing contribution of the rim emission. Thus, for very high-inclination NLs with $i>85^\\circ$ this will definitely be a problem, especially when the rim stars shielding part of the disc. Note that the impression for a better reconstruction of the disc's edge, compared to the flat disc simulations is misleading. This is due to the additional light from the rim, which compensate for the smoothing of the disc edge. For \\object{DW UMa} with $i\\sim82^\\circ$, we may conclude that emission of the outer disc rim cannot be held responsible for the flat intensity profile, unless it is very bright or/and high. Note also that with increasing rim height the scatter of the reconstructed intensities at given radius also increases. The scatter in the maps of  \\object{DW UMa} is nearly twice less than that in the simulation with the rim height  $0.13R_{\\rm disc}$, suggesting that the contribution of the rim in \\object{DW UMa} is not very significant. We conclude then that the flat temperature profile is not caused by the operation of the algorithm and the assumption of flat disc geometry.  Finally, we would like to note that the \\object{SW Sex} stars are perhaps the most complex of all types of CVs (Warner \\cite{war}). The presence of overflowing stream, possibly a flared accretion disc, accretion disc wind and a magnetic white dwarf violate some of the assumptions of eclipse mapping. In addition, many of the \\object{SW Sex} stars show superhumps, both \"positive\" and \"negative\". The former are now well understood to arise in an elliptic, progradely precessing accretion disc and will certainly complicate the situation ever more.  More generally, the presence of hot spots and spiral shocks should also be considered. Detailed investigation of the effect of all these factors on the eclipse mapping is a very complicated task and is out of the scope of this article. We will only  briefly discuss some of these topics.  The largest contribution of the spiral shocks should be $\\sim10-15$\\% (e.g. Baptista et al. \\cite{bap00,bap02}). When spread over the disc, this will be much smaller and can hardly affect the radial profile. Pre-eclipse humps similar to those in quiescent dwarf novae are almost never observed in \\object{SW Sex} stars, suggesting little contribution of the hot spot, if any at all. This is expected though. The traces of the hot spot in the eclipses of dwarf novae disappear in outburst. This generally implies little contribution of the spot when the disc is hot. Because the discs in \\object{SW Sex} stars are hot, the spot should little affect the eclipse mapping results of these objects. Probably this is also the case with the superhumps, whose light source in NLs contributes no more than 10-20\\% of the total system light. A bigger problem with superhumps would be the different disc geometry, for example the elliptic disc shape during the \"positive\" superhumps. At least in one system, \\object{PX And} (Stanishev et al. \\cite{stpx}), the eclipse depth varies greatly with the \"negative\" superhumps phase. This effect may be due to the presence of precessing tilted accretion disc and would have serious implications if the inclination is close to $90^\\circ$. Thus, both superhumps will affect the eclipse profile. However, when averaging many eclipses observed at different superhump phases, the effect of the superhumps will be reduced and  as a result average disc maps will be obtained. Other possible effects such as the spirals being thicker than the rest of the disc and any light scattered in the disc wind or/and the overflowing stream are difficult to take into account and estimate. In conclusion, we may say that various systematic effects may play a role in the eclipse mapping reconstructions of the high-inclination NLs. Thus, the results of such experiments should be always discussed with this in mind."
    },
    "0403/astro-ph0403146_arXiv.txt": {
        "abstract": "{A new sample of northern isolated galaxies, which are defined by the physical criterion that they were not affected by other galaxies in their evolution during the last few Gyr, is presented. To find them we used the logarithmic ratio, $f$, between inner and tidal forces acting upon the candidate galaxy by a possible {\\it perturber}. The analysis of the distribution of the $f$-values for the galaxies in the Coma cluster lead us to adopt the criterion $f \\leq -4.5$ for isolated galaxies. The candidates were chosen from the CfA catalog of galaxies within the volume defined by cz $\\leq$5000 km/s, galactic latitude higher than 40\\deg~ and declination $\\ge -$2.5\\deg. The selection of the sample, based on redshift values (when available), magnitudes and sizes of the candidate galaxies and possible perturbers present in the same field is discussed. The final list of selected isolated galaxies comprises of 203 objects from the initial 1706. The list contains only truly isolated galaxies in the sense defined, but it is by no means complete, since all the galaxies with possible companions under the $f$-criterion but with unknown redshift were discarded. We also selected a sample of perturbed galaxies comprising of all the disk galaxies from the initial list with companions (with known redshift) satisfying  $f \\geq -$2 and $\\Delta(cz) \\le 500$ km/s, a total of 130 objects. The statistical comparison of both samples shows significant differences in morphology, sizes, masses, luminosities and color indices. Confirming previous results, we found that late spiral, Sc type galaxies are, in particular, more frequent among isolated galaxies, whereas Lenticular galaxies are more abundant among perturbed galaxies. Isolated systems appear to be smaller, less luminous and bluer than interacting objects. We also found that bars are twice as frequent among perturbed galaxies compared to isolated galaxies, in particular for early Spirals and Lenticulars. The perturbed galaxies have higher L$_{FIR}$/L$_B$ and M$_{mol}$/L$_B$ ratios, but the atomic gas content is similar for the two samples. The analysis of the luminosity-size and mass-luminosity relations shows similar trends for both families, the main difference being the almost total absence of big, bright and massive galaxies among the family of isolated systems, together with the almost total absence of small, faint and low mass galaxies among the perturbed systems. All those aspects indicate that the evolution induced by interactions with  neighbors, would proceed from late, small, faint and low mass Spirals to earlier, bigger, more luminous and more massive spiral and lenticular galaxies, producing at the same time a larger fraction of barred galaxies, but preserving the same relations between global parameters. The properties we found for our sample of isolated galaxies appear rather similar to those of high redshift galaxies, suggesting that the present day isolated galaxies could be quietly evolved, unused {\\sl building blocks} surviving in low density environments.  \\keywords {Isolation criteria -- Disk  Galaxies -- Interacting Galaxies}} ",
        "introduction": "Galaxies are presently most commonly found in aggregates of different density. Inside these aggregates, the observed morphology and the segregation of types \\citep{dressler, fas00} may have been determined by the local density of matter and by the subsequent interactions between galaxies. Exchanges and interaction with its neighbors can affect the galaxy global properties and their star formation rates, as shown by the members of dense groups or pairs \\citep{moles}, even if the effects, in general, are not as dramatic as they are in the strongest cases (Moles et al., 1994b; M\\'arquez \\& Moles, 1999).  Another way to look into the effects of interaction on the overall equilibrium of a galaxy is through the analysis of the relations found between the structural parameters for different families of galaxies. The E/S0 cluster (Faber et al., 1987; Davis \\& Djorgovski, 1987; J\\o rgensen, Franx, \\& Kj\\ae rgaard, 1996), and field galaxies (Treu et al., 2001) satisfy the Fundamental Plane (FP) relation, even if its form is not the same for both families. On their side, the Spiral (S) galaxies satisfy the Tully-Fisher (TF) relation. It has been indicated that the TF relation defined by isolated spiral galaxies presents a smaller scatter than non-isolated galaxies (M\\'arquez \\& Moles, 1996; M\\'arquez et al., 2002; van Kannappan et al., 2002).  To find isolated galaxies, a physically grounded and operational definition of isolation has to be given. Often one speaks of field galaxies to refer to systems that are not in dense aggregates, but frequently they are still under the influence of neighboring galaxies and cannot be considered as genuinely isolated objects (see for example M\\'arquez et al., 2002). The first systematic compilation of isolated galaxies was made by Karachentseva (1973), who used as the main criterion that they don't show companions within 20 diameters in the PSS plates. Later, \\cite{tg76}, while searching for groups of galaxies, found that some of the galaxies couldn't be assigned to any of the identified groups. It was argued that they would trace a cosmological homogeneously distributed population of galaxies. This conclusion was however challenged by \\cite{ht77} who used a deeper sample (m $\\leq$ 15.7 instead of m $\\leq$ 14). The next step was taken by \\cite{vsc86}, who adopted stricter criteria to find isolated galaxies. Their conclusion was that it was very unlikely that there were a cosmological component of galaxies without any kind of clustering (at least at small redshift). But an important observation was the detection of galaxies in such poor environments that they had very probably evolved without any external perturbation, as if they were isolated. In that perspective what is needed is an operational definition of isolation that could lead to identifying galaxies that have evolved free from external influences for most of its life, a point of view adopted by \\cite{mm96} and M\\'arquez et al. (1999). In this context, it was noticed that non-axisymmetric structures like bars, tails, or plumes, which are usually explained invoking gravitational interaction with companions or satellites, can also be present in isolated galaxies, posing the problem of their origin in absence of sizable interactions (Moles et al., 1994; Moles, M\\'arquez \\& P\\'erez, 1995).  In the present work we discuss the criteria used to find isolated galaxies and build a statistically well defined sample to analyze their properties. Due to the lack of redshift information for most of the faint, possible {\\it perturbers}, the final sample is by no means complete, our main goal being to keep only truly isolated galaxies, discarding all the suspect systems. In the last section we compare the properties of the family of the isolated galaxies with that of confirmed perturbed galaxies, to show the main differences between both families.  We also consider the similarity of some of the properties of present day isolated systems with those reported for the presumed {\\sl building blocks} in the distant Universe \\citep{ferg03,trujillo}. The point is that stellar systems situated in regions of very low density would have evolved differently from galaxies in aggregates without a significant contribution from the surroundings and, therefore they might be similar to those original building blocks. ",
        "conclusions": "We used a well defined, physically meaningful criterion to define isolated galaxies. We performed a search of isolated galaxies using as the parameter the logarithmic ratio $f$ between the internal and tidal forces at the outskirts of a given galaxy. The adopted limit, $f < -$4.5, was checked with results from numerical simulations \\citep{byrd92}, and with data from the Coma Cluster. Similarly, a sample of {\\sl perturbed} galaxies was defined comprising of galaxies with confirmed companions satisfying $f \\geq -$2.  Comparison of the properties of the galaxies in both samples was made using the Mann-Whitney U-test. The first result to notice is the significant differences in morphological types, with Sc types being more abundant among isolated galaxies, whereas S0 galaxies are significantly more abundant among perturbed systems. We also found that barred galaxies are more frequent (about twice) among perturbed than among isolated galaxies except for late type Spiral galaxies.  They also differ in all the properties we have examined, except in atomic gas content, with similar values of the ratio M$_{HI}$/L$_B$. The isolated galaxies appear to be smaller and fainter, with bluer color indices. They have less molecular gas an FIR luminosity per unit of blue luminosity as well. The dynamical parameters are also consistently smaller for isolated galaxies. Even if the differences can be appreciable for all the families (S0, early Sp and late Sp), the biggest differences are for S0 and early Sp. The exception is in molecular gas content, for which the maximum difference is found for late type Sp galaxies.  The M-L and M-size relations are consistent in showing the absence of big, luminous and massive systems among the isolated galaxies. Moreover, we also found the almost complete absence of perturbed systems with small sizes, low luminosities and low masses. We also found that barred galaxies do not occupy a particular region in the T-F diagram.  In view of all these results it is tempting to consider the differences as arising from the different evolving conditions. Our results indicate that the gravitational interaction in aggregates would produce evolution from late Spiral, relatively faint and low mass galaxies, toward earlier, more luminous and massive Spiral and Lenticular types. This would also favor the formation of bars in early type Spirals and Lenticulars. However, the relations between global parameters are similar for both families, even if they tend to occupy different regions in the corresponding plots (see Figures 2, 3 and 4)  Since the isolated galaxies wouldn't have had the opportunity to accrete other systems and grow in the way the hierarchical models predict, they would still be similar today to the high redshift systems supposed to be the {\\sl building blocks} for the formation of larger galaxies. This view is consistent with the results reported by \\citet{ferg03}, showing that the sizes of galaxies at z $\\approx$ 4 are smaller than nearby luminous galaxies. If cannibalism and accretion in the early phases of galaxy evolution are important to fix the final global properties such as size, luminosity and mass, isolated galaxies would resemble those initial pieces, the left over fragments in the early process of galaxy formation."
    },
    "0403/astro-ph0403478_arXiv.txt": {
        "abstract": "The combination of dynamical and photometric properties of galaxies offers a largely un-tapped source of information on how galaxies assembled and where stars formed. Bi-dimensional kinematic measurements have been the stumbling block. The light-gathering power of SALT coupled with the high-throughput performance of the Prime Focus Imaging Spectrograph (PFIS) yield a superb facility for measuring velocity-ellipsoids of stars and gas in galaxies out to gigaparsec distances. From these data dynamical asymmetries arising from lopsided or elliptical halos may be probed; disk-mass and mass-decompositions may be uniquely determined; mechanisms for disk heating constrained; and a zeropoint for the mass-to-light ratios of stellar populations set. A number of groups within the SALT consortium are interested in making these measurements using a variety of different, but complementary approaches. The scientific potential from their synthesis is very promising. We describe some unusual observational modes in which PFIS may be used to probe the shape of dark-matter halos and the content of galaxy disks. ",
        "introduction": "Extant information about galaxies comes primarily from broad-band optical images. These deliver a wealth of information on when and where stars have formed. Missing is an accurate understanding of the distribution of mass and the details of its assembly -- knowledge which requires dynamical arguments substantiated with kinematic measurements. Hence kinematic information is essential for a basic understanding of galaxy formation and evolution. The last decade has seen the growth of bi-dimensional optical spectroscopy -- once the purview of radio (HI) observations. SALT can continue this development. With SALT's Prime Focus Imaging Spectrograph we will be able to measure velocity-ellipsoid maps for a variety of dynamical tracers which include (collisionless) stellar populations and (`sticky') ionized gas. These measurements will be accessible via Fabry-Perot and dispersed spectroscopy over a very large (8$^\\prime$) field.  \\begin{figure} \\plotfiddle{conf_sparsepak_s.ps}{2.0in}{0}{34}{34}{-200}{-18.5} \\plotfiddle{ppak_s.ps}{0in}{0}{35.7}{35.7}{-10}{-16.5} \\caption{{\\it SparsePak} ({\\bf left}, Bershady et al. 2003, 2004) and {\\it PPAK} ({\\bf right}, Verheijen et al. 2004) IFUs, which respectively feed spectrographs on the WIYN and Calar Alto 3.5m telescopes, both capable of $\\sim$11 km/s ($\\sigma$) resolution. Units are presented at the same subtended angular scale at the telescope focal-plane; physical dimensions scale with respective telescope+fore-optic $f$-ratios. Light fibers transmit to the spectrograph; dark fibers serve as mechanical buffers.} \\vskip -0.1in \\end{figure}  As promising for SALT is the critical mass of consortium scientists with overlapping interests in galaxy dynamics (see, for example, contributions by Balona, Cecil, Cote, Crawford, Ferrarese, Sellwood, Sparke, Wilcots, Williams, and Ziegler in this volume). A few specific scientific prospects of interest to members of the consortium include (but are not limited to!) the frequency and amplitude of lopsided disks and triaxial halos as probed by dynamical and photometric asymmetries; mergers and perturbations inferred studies of kinematic irregularities; disk heating mechanisms judged from velocity dispersion ellipsoids ($\\sigma_z/\\sigma_R$, Gerssen et al. 2000); feed-back processes in disks and the ISM; bar dynamics and their pattern-speeds; the Tully-Fisher relation and its evolution; and direct dynamical mass decompositions determined by combining rotation curves with stellar velocity dispersions.  Here we focus on how well we can directly separate disk from halo mass in spirals to yield dark-halo profiles, the zeropoint for stellar M/L ratio, and plausible constraints on the faint-end of the stellar IMF. These are of fundamental interest, respectively for testing both hierarchical cold-dark-matter structure-formation scenarios and star-formation models. This particular science case allows us to illustrate a novel application of PFIS to bi-dimensional spectroscopy with wide application. The concept combines narrow-band filters with dispersed spectroscopy to form what we call {\\it m}assively {\\it m}ulti-plexed {\\it s}pectroscopy, or MMS. The approach complements other methods using fluid-dynamical modeling and kinematic measurements of barred potentials (Weiner et al. 2001; Gerssen et al. 2003; Debattista \\& Williams 2004). ",
        "conclusions": "\\begin{figure} \\plotfiddle{other_ifu_rv1_top.ps}{2.75in}{-90}{75}{75}{-210}{265} \\caption{Grasp versus spectral power for two-dimensional spectroscopic systems on 4m and 8m-class telescopes (solid and dashed lines, respectively, with Fabry-Perot instruments shown as filled circles). Total grasp is the product of area $\\times$ solid-angle (A$\\Omega$). Spectral power is the product of the spectral resolution, $\\lambda/\\Delta\\lambda$, times the number of spectral resolution elements, N$_{\\Delta\\lambda}$.  Variations in the shapes of covered parameter space depends on how a given instrument achieves a range of spectral resolution and spatial sampling, i.e., through changes in gratings, slit-widths, or both.  Note the unique location of PFIS in these diagrams is significantly extended with the use of MMS.} \\vskip -0.05in \\end{figure}  Galaxy kinematic measurements with SALT's PFIS can be used to answer many outstanding, and fundamental questions about the structure and formation of galaxies. Highlighted here is one example of unique, dynamical mass decompositions of spiral galaxies. A number of observational modes are available with PFIS. We have described one powerful, new mode, referred to as ``massively multiplexed spectroscopy,'' (MMS), which is enabled by the unique combination of VPH gratings and narrow-band filters in PFIS. With MMS and Fabry-Perot imaging, PFIS is a highly competitive survey instrument for two-dimensional studies of gas {\\it and} stars in nearby galaxies. There are several groups within the SALT consortium interested in applying these to a similar range of dynamical problems. This holds the promise of fruitful collaboration, cross-checks, and innovation.  What about higher redshift studies (see Ziegler and Crawford, this volume)? PFIS can sample a 0.6 arcsec slitwidth, but the sub-arcsec image quality enjoyed regularly by VLT and Gemini is needed unless we remain content with spatially-unresolved line-widths.  Should it prove feasible to phase the SALT primary segements, we may consider pushing aggressively for significantly improved image quality. This will have a major impact on both resolved, extragalactic as well as high-resolution stellar spectroscopic observations. But SALT is merely a 9m telescope; for extended sources observed at significant spectral resolution, the area--solid-angle product (A$\\Omega$) is what counts. As other facilities rush for the highest possible angular resolution that, with today's detectors, may be better suited for tomorrow's 30m-class telescopes, we should keep good aim on the type of ground-breaking science enabled by our large field of view and the exquisite, multi-faceted spectroscopic capabilities of PFIS."
    },
    "0403/astro-ph0403152_arXiv.txt": {
        "abstract": "{ We have undertaken a thorough dynamical investigation of five extrasolar planetary systems using extensive numerical experiments on the supercomputer of the Max Planck Institute for Gravitational Physics (Albert Einstein Institute). This work was performed as part of the Helmholtz Institute for Supercomputational Physics Summer School on ``Chaos and Stability in Planetary Systems'' (2003). The systems Gl~777~A, HD~72659, Gl~614, 47~Uma and HD~4208 were examined concerning the question of whether they could host terrestrial like planets in their habitable zones (=HZ).  First we investigated the mean motion resonances between fictitious terrestrial planets and the existing gas giants in these five extrasolar systems. Then a fine grid of initial conditions for a potential terrestrial planet within the HZ was chosen for each system, from which the stability of orbits was then assessed by direct integrations over a time interval of 1 million years. For each of the five systems the 2-dimensional grid of initial conditions contained 80 eccentricity points for the Jovian planet and up to 160 semimajor axis points for the fictitious planet. The computations were carried out using a Lie-series integration method with an adaptive step size control. This integration method achieves machine precision accuracy in a highly efficient and robust way, requiring no special adjustments when the orbits have large eccentricities.  The stability of orbits was examined with a determination of the R\\'enyi entropy, estimated from recurrence plots, and with a more straight forward method based on the maximum eccentricity achieved by the planet over the 1 million year integration. Additionally, the eccentricity is an indication of the habitability of a terrestrial planet in the HZ; any value of $e>0.2$ produces a significant temperature difference on a planet's surface between apoapse and periapse. The results for possible stable orbits for terrestrial planets in habitable zones for the five systems are summarized as follows: for Gl~777~A nearly the entire HZ is stable, for 47~Uma, HD~72659 and HD~4208 terrestrial planets can survive for a sufficiently long time, while for Gl~614 our results exclude terrestrial planets moving in stable orbits within the HZ.  Studies such as this one are of primary interest to future space missions dedicated to finding habitable terrestrial planets in other stellar systems. Assessing the likelihood of other habitable planets, and more generally the possibility of other life, is the central question of astrobiology today. Our investigation indicates that, from the dynamical point of view, habitable terrestrial planets seem to be quite compatible with many of the currently discovered extrasolar systems. \\vspace{2cm} ",
        "introduction": "In this study\\footnote{The results presented here were derived during the $3^{\\mathrm{rd}}$ Helmholtz Summerschool at the Helmholtz Institute for Supercomputational Physics in Potsdam: ''Chaos and Stability in Planetary Systems'' from September $1^{\\mathrm{st}}$ to September $26^{\\mathrm{th}}$ 2003} the extrasolar systems Gl~777~A, HD~72659, Gl~614, 47~Uma and HD~4208 were examined on the question of whether they could host additional terrestrial like planets in their habitable zones (=HZ).  Since the discovery of the first extrasolar planetary system about 10 years ago~(Mayor \\& Queloz \\cite{mayor}), a major point of dynamical investigations has been the determination of stable regions in extrasolar planetary systems, where additional planets on stable orbits could exist. Today we have evidence for 105 planetary systems with 120 planets, where 13 systems have more than one planet.  Because of the observational methods which are in use, our knowledge of extrasolar planets is highly biased: almost all planets were detected via ground based measurements of the central star's radial velocity. This method favors the detection of very large (Jupiter-sized) planets that move very close to the main star. Several space missions are planned (i.e. {\\sf COROT}, {\\sf KEPLER}, {\\sf TPF}, {\\sf DARWIN}) that will be able to find planets with a much smaller mass (comparable to that of the Earth) by using other detection methods. The first results of these missions will only be available after more than a decade, however, theoretical studies can currently make predictions on the existence of stable terrestial planets within the known systems. This serves to establish testable theories of planet formation and evolution, as well as aiding the future searches for such habitable terrestial planets.  When one looks at the distribution of semimajor-axes of all known exosolar planets, one can see that they are located from $\\sim 0.02$ AU to $\\sim 6.5$ AU. These distances cover the so-called ``habitable zone'' (=HZ)\\footnote{The region were a planet fulfills all the requirements of being able to develop life and where liquid water can be present on the surface of the planet.} of a typical main sequence star, ranging from $\\sim 0.7$~AU to $\\sim 1.5$ AU~(Kasting et al. \\cite{Kasting93}). Some of these large planets move in the region around 1 AU that overlaps with the HZ, and leave no room for possible terrestrial planets orbiting the host star within the HZ -- only a large satellite (like Titan) or a Trojan like planet could have all the properties which are necessary for orbital stability and also the development of a stable atmosphere in such cases. When the large planets move well outside (or inside) the HZ they will perturb any additional undetected planet that may move within this zone.  For extrasolar systems with only one planet the possible existence of additional planets has been investigated (e.g. Dvorak et al. \\cite{Dvorak03b}; P\\'al \\& S\\'andor~\\cite{Pal03}). In an extensive paper (Menou \\& Tabachnik~\\cite{Menou03}) the stability of possible terrestrial planets was studied for all extrasolar systems and a classification according to their stability was established, but there the role of resonances was underestimated. Resonances -- mean--motion resonances and secular resonances --  can either stabilize the motion of a planet and protect it from close encounters and collisions or, in other cases, intensify the gravitational perturbations and therefore destroy the stability of the orbits. Resonances between the orbits of the observed large planets have been studied (Callegari et al.~\\cite{Callegari02}; Hadjidemetriou~\\cite{Hadjidemetriou02}; \\'Erdi \\& P\\'al~\\cite{Erdi03}), but a detailed study of  the perturbations on possible additional planets due to resonances is still lacking. Thus, any study of the dynamics inside the HZ of an extrasolar system has also to include an investigation of the resonances.  Determinations of the dynamical stability of multiple planetary systems (Kiseleva et al. \\cite{Kiseleva02a},~\\cite{Kiseleva02b}; Beaug\\'e \\& Michtchenko~\\cite{Beauge03}) and of planets in double stars (Pilat-Lohinger \\& Dvorak \\cite{Pilat02a}; Pilat-Lohinger et al.~\\cite{Pilat02b},~\\cite{Pilat03}; Dvorak et al.~\\cite{Dvorak03a}; Lammer et al.~\\cite{Lammer03}) have been done recently.  In this article we  investigate the dynamics inside the habitable zone of five different extrasolar systems: Gl~777~A, 47~Uma, HD~72659, Gl~614 and HD~4208. Except 47~Uma, all these systems consist of the central star and a Jovian planet. According to the spectral type of the main star, the habitable zone was determined; a detailed analysis of the mean motion resonances inside the HZ was carried out and the full width of the HZ was investigated according to the stability of possible additional terrestrial planets. This article is organized as follows: Section~\\ref{sec2} describes the five extrasolar systems for which we performed a dynamical study for additional planets, in section~\\ref{sec3} we introduce the dynamical models and the methods used to analyze the time series data resulting from our numerical experiments. Section~\\ref{sec4} shows how we dealt with the terrestrial planet orbits in the major mean--motion resonances of the Jovian planet(s). Subsequent sections focus on each of the 5 systems individually, detailing our investigations and results. We conclude with a summary of these results and their implications, including comparison to earlier, less detailed studies by Menou \\& Tabachnik~(\\cite{Menou03}). The appendicies include concise descriptions of the analytical and numerical methods used, and may serve as a useful reference to the reader. ",
        "conclusions": "We have undertaken a dynamical study of five extrasolar planetary systems using extensive numerical experiments to answer the question of whether they could host terrestrial like planets in habitable zones. For the single--planet systems we used the elliptic restricted three body problem, and for the system 47 Uma the mutual perturbations of the two Jovian planets were taken into account and therefore the restricted four-body problem served as a dynamical model. Because of the dependence of the stability of an orbit in a mean motion resonance on angular position we have selected 8 different positions for the mean anomaly of the planet with a step of $45^\\mathrm{\\circ}$ in the  apoastron and the periastron position of the Jovian planet.  The characteristics of each system dictated the initial conditions, chosen in a fine 80 by 80 grid within the habitable zone, from which the orbits were computed using a robust numerical method (Lie--series integration) for 1 million years. The grid of the initial conditions of the fictitious terrestrial planets was chosen to cover the whole habitable zone of the system, and also to model the uncertainties in the elements of the observed planet(s). The stability of orbits was assessed with two methods, namely the computation of the R{\\'e}nyi entropy as measure of the chaoticity of an orbit and the determination of the maximum eccentricity of the orbit of a fictitious planet during its orbital evolution of 1 million years.  We can say that our computations for such a fine grid, taking into account also the essential role of the MMRs, lead to a deeper insight concerning the dynamics of the five systems which we studied. We also give the percentage of orbits which survived in the paper (=MT) of Menou \\& Tabachnik~(\\cite{Menou03}) where they investigated all known extrasolar planetary systems with respect to possible additional terrestrial planets. We note that a direct comparison of MT with the percentages of 'our' survivors is not useful here because of the different approaches used; we have emphasized the role of the MMRs and neglected possible inclinations. However, we know that terrestrial planets will form within a protoplanetary disk thus staying with small orbital inclinations (Richardson et al.~\\cite{Richardson00}; Lissauer~\\cite{Lissauer93}); additionally  in a recent publication (Dvorak et al.~\\cite{Dvorak03a}) it was shown that the inclinations of the fictitious planets up to 15 degrees do not change the stability of orbits in the HZ. The results for possible stable orbits may be summarized as follows:  \\begin{itemize}  \\item In the system {\\bf Gl~777~A} the stability zone for the motion of terrestrial planets is well inside of the HZ and suggests any planets residing there will survive for a sufficiently long time (in MT, 86.8 \\% of the orbits were found to be stable).  \\item  In the system {\\bf 47 Uma} we can also say that there is a good chance for planets to move inside the HZ with small eccentricities between the main resonances; this is a result which is consistent with others but not with MT where only 28 \\% remained.  \\item {\\bf HD 72659} turned out to be a very good candidate for hosting planets in the HZ; again this does not confirm the results of MT -- they found that only 40.2 \\% of the orbits were stable.  \\item    The results of the computations for {\\bf Gl 614} show that  it is very unlikely that there is an additional planet moving in the HZ -- these results are consistent with MT (9.2 \\% stable orbits)  \\item For {\\bf HD 4208}  we can say that in the HZ there is enough room left for terrestrial planets and that they could survive for sufficiently long time; these results are more or less consistent with those of MT where 50.2 \\% of the orbits were stable.  \\end{itemize}  New observational possibilities provided by missions like COROT, DARWIN or the TPF make the first search for terrestrial exoplanets seem possible in the next decades. In an ESA study the goal of the missions is summarized as follows: ''To detect and study Earth--type planets and characterize them as possible abodes of life''. In this sense dynamical studies like the one we present here should help to define promising targets for observations."
    },
    "0403/astro-ph0403364_arXiv.txt": {
        "abstract": "The Submillimeter Array (SMA) has been under construction at the 4100 m summit of Mauna Kea, Hawaii. The array is going to allow imaging of  lines and continuum at sub-arcsecond resolution in submillimeter wavelengths. The status of the array and the results from recent commissioning observations of nearby galaxies are reported. ",
        "introduction": "Molecular gas and dust in starburst galaxies are known to be warmer and denser than those in non-starburst environments. Emission lines from highly excited molecules and dust continuum in submillimeter are therefore among the best probes to study neutral ISM in starburst galaxies. Detailed study of starbursts in external galaxies needs a resolution of the order of an arcsecond. An interferometer is necessary to achieve this resolution in submillimeter. However, there has not been a submillimeter interferometer since the pioneering JCMT--CSO interferometry (c.f. Wiedner 2002). SMA is going to be the first full-time submillimeter interferometer with fast imaging capability of eight antennas. Though still in construction and testing phase,  the array has recently started its initial observations of galaxies as reported here. Another contribution by Iono et al. in these proceedings shows preliminary results of more distant starbursts. ",
        "conclusions": ""
    },
    "0403/astro-ph0403708_arXiv.txt": {
        "abstract": "We present preliminary results of spectroscopic data obtained with the La Silla 3.6\\,m TIMMI2 and the VLT-ISAAC instruments. We have obtained 23 TIMMI2 N-band and 18 Q-band spectra for our total sample of 36 stars in March 2003 to investigate the evolution of the grain properties in circumstellar disks. Our sample consists of Herbig stars, T\\,Tauri stars and Vega-like stars, some of which have recently been identified to be members of nearby ($<$100\\,pc) stellar associations.  Silicate and Polycyclic Aromatic Hydrocarbon (PAH) emission features are detected and the spatial resolution of TIMMI2 allowed to place upper limits on the spatial extent of the emission. We also obtained ISAAC L-band spectra of 4 well-studied disks. For the Herbig Be star HD\\,100546, we detect spatially extended PAH emission at 3.3\\,$\\umu$m which corresponds to a source with radius 11$\\pm$3\\,AU in this line. ",
        "introduction": "In the last decade, circumstellar disks at different evolutionary stages have been intensively studied from optical to millimeter wavelengths. Low-mass (T~Tauri) and intermediate-mass (Herbig Ae/Be) young stars have been shown to have young dusty and gaseous protoplanetary disks for which imaging can provide constraints on the structure \\cite{dar03,man97,wol03}. More evolved gas-free dust disks around main-sequence stars, studied at sub-millimeter and near-infrared wavelengths, show a ``lumpy'' structure with gaps, holes and/or asymmetries in the dust distribution, indicative of gravitational perturbations by stellar companion(s) or unseen planet(s) \\cite{aug01b,aug04,wya03}.  Near- and mid-infrared spectroscopy of circumstellar disks around Pre Main Sequence (PMS) stars at different evolutionary stages has shown a wealth of gas and dust features, enabling detailed studies of dust mineralogy and its relation with grain processing and disk evolution \\cite{mee01}. The shape of the 10\\,$\\umu$m feature can serve as a tracer of grain growth in disks and the presence of crystalline silicates, not observed in the ISM, is interpreted as evidence that the dust composition evolves as the disk is dissipating and large protoplanetary bodies are forming \\cite{prz03,boe03}. ",
        "conclusions": ""
    },
    "0403/astro-ph0403222_arXiv.txt": {
        "abstract": "Irregular (Irr) and dwarf irregular (dIrr) galaxies are gas-rich galaxies with recent or ongoing star formation.  In the absence of spiral density waves, star formation occurs largely stochastically.  The scattered star-forming regions tend to be long-lived and migrate slowly.  Older populations have a spatially more extended and regular distribution.  In fast-rotating Irrs high star formation rates with stronger concentration toward the galaxies' center are observed, and cluster formation is facilitated. In slowly or nonrotating dIrrs star formation regions are more widely distributed, star formation occurs more quiescently, and the formation of OB associations is common.  On average, Irrs and dIrrs are experiencing continuous star formation with amplitude variations and can continue to form stars for another Hubble time.  Irrs and dIrrs exhibit lower effective yields than spirals, and [$\\alpha$/Fe] ratios below the solar value.  This may be indicative of fewer Type II supernovae and lower astration rates in their past (supported by their low present-day star formation rates).  Alternatively, many metals may be lost from the shallow potential wells of these galaxies due to selective winds.  The differences in the metallicity-luminosity relation between dIrrs and dwarf spheroidals (which, despite their lower masses, tend to have too high a metallicity for their luminosity as compared to dIrrs) lends further support to the idea of slow astration and slow enrichment in dIrrs.  The current data on age-metallicity relations are still too sparse to distinguish between infall, leaky-box, and closed-box models.  The preferred location of dIrrs in the outer parts of galaxy groups and clusters and in the field as well as the positive correlation between gas content and distance from massive galaxies indicate that most of the dIrrs observed today probably have not yet experienced significant interactions or galaxy harassment. ",
        "introduction": "In this contribution, I will focus on the evolutionary histories of irregular (Irr) and dwarf irregular (dIrr) galaxies, including their chemical evolution. The name ``irregular'' refers to the irregular, amorphous appearance of these galaxies at optical wavelengths, where the light contribution tends to be dominated by scattered bright H\\,{\\sc ii} regions and their young, massive stars.  Irrs are typically gas-rich galaxies that lack spiral density waves as well as a discernible bulge or nucleus. Many Irrs are disk galaxies and appear to be an extension of late-type spirals.  The most massive disky Irrs with residual spiral structure are also called Magellanic spirals; e.g., the Large Magellanic Cloud (LMC) (Kim et al.\\ 1998) is a barred Magellanic spiral.  Looser and more amorphous Irrs like the Small Magellanic Cloud (SMC) are sometimes also referred to as Magellanic irregulars (or barred Magellanic irregulars if a bar is present); see de~Vaucouleurs (1957). A different system of subdivisions was suggested by van~den~Bergh in his DDO luminosity classification system (van~den~Bergh 1960, 1966). DIrrs are simply less massive, less luminous Irrs; the distinction between the two is a matter of definition rather than physics.  Typical characteristics of dIrrs are a central surface brightness $\\mu_V \\lesssim 23$ mag arcsec$^{-2}$, a total mass of $M_{\\rm tot} \\lesssim 10^{10} M_{\\odot}$, and an H\\,{\\sc i} mass of $M_{\\rm HI} \\lesssim 10^9 M_{\\odot}$. Solid body rotation is common among the Irrs and more massive dIrrs, while low-mass dIrrs do not show measurable rotation; here random motions dominate. A typical characteristic of Irrs and dIrrs alike is ongoing or recent star formation.  The star formation intensity may range from burst-like, strongly enhanced activity, to slow quiescent episodes. Irrs and dIrrs can continue to form stars over a Hubble time (Hunter 1997).  Substantial progress has been made in the photometric exploration of the star formation histories of Irrs and dIrrs over the past decade, largely thanks to the superior resolution of the {\\it Hubble Space Telescope}. There is still very little known about the progress of the chemical evolution in these galaxies as a function of time, but the advent of 6-m to 10-m class telescopes and their powerful spectrographs is beginning to change this situation.  The most detailed information is available for nearby Irr and dIrr galaxies in the Local Group, most notably the LMC, which is the Irr galaxy closest to the Milky Way.  \\begin{figure*} \\centering \\includegraphics[width=0.85\\columnwidth,angle=270]{grebel_fig1.ps} \\vskip 0pt \\caption{ Morphological segregation in the Local Group (filled histograms; see Grebel 2000) and in the M81 and Cen A groups (dashed histograms; input data from Karachentsev et al.\\ 2002a, b).  Note the pronounced concentration of gas-poor, early-type dwarfs around the nearest massive primary galaxy, while the gas-rich, late-type dwarfs show less concentration and are more widely distributed. This may be a signature of the impact of environmental effects, such as gas stripping.} \\label{fig1} \\end{figure*} ",
        "conclusions": ""
    },
    "0403/nlin0403032_arXiv.txt": {
        "abstract": " ",
        "introduction": "A sunspot \\cite{bray} is a relatively cool area on the solar face that appears as a dark blemish. The number of sunspots is continuously changing in time in a random fashion and constitutes a typical, a priori, random time series. Because of the symmetry of the twisted magnetic lines that originate sunspots, they are generally seen in pairs or in groups of pairs at both sides of the solar equator. As the sunspot cycle progresses, spots appear closer to the sun's equator giving rise to the so called ``butterfly diagram'' in the latitude distribution. Magnetic fields about the sunspots are very strong and keep heat out of these regions on the sun surface. They are formed when magnetic field lines are twisted and poke through the solar photosphere. The twisted magnetic fields above sunspots are sites where solar flares are observed. It has been found that chromospheric flares show a very close statistical relationship with sunspots. This is illustrated by the Waldmeir relationship \\cite{bray} between the mean number of flares per day, $E$, and the sunspot number, $R$, namely $E = 0.061\\,R$. Being the number of sunspots much larger than the corresponding to flares, it seems more appropriate for statistical analysis. Moreover, this connection between sunspots and solar magnetic field is particularly motivating for a further analysis of their frequency of appearance. In summary, being a consequence of the solar cycle, the variation of the number of sunspots is closely related to the so-called solar activity. In fact, the solar maximum activity corresponds to many sunspots present, while the minimum of activity is related to very few sunspots.  Another effect of the solar activity is the presence of fluctuations in the interplanetary magnetic field. Experimental data suggest \\cite{matt} a $1/f$ dependence for the frequency spectra of its fluctuations in the range of large times, namely, $2.7\\times10^{-6}$ Hz to $1.8\\times 10^{-5}$ Hz. In this analysis it is argued that on the basis of the observed behavior there is a superposition of signals with scale invariant distributions of correlation times. In other words, a kinematical superposition of signals due to scale-invariant reconnection of magnetic structures, near the solar surface, gives rise to the $1/f$ spectrum. It should be also noticed  that at higher frequencies the interplanetary spectrum is known to have approximatively an $1/f^{5/3}$ behavior; this time associated with magnetohydrodynamics (MHD) homogeneous turbulence.  The result of our analysis of the sunspot number distribution along almost 250 years shows a clear correlation among all these phenomena.  After introducing the data and the different methods of analysis used, we present the results with a further comment on their implications. ",
        "conclusions": ""
    },
    "0403/hep-th0403173_arXiv.txt": {
        "abstract": "We elaborate further on the evolution properties of cosmological fluctuations through a bounce. We show this evolution to be describable either by ``transmission'' and ``reflection'' coefficients or by an effective unitary $S-$matrix. We also show that they behave in a time reversal invariant way. Therefore, earlier results are now interpreted in a different perspective and put on a firmer basis. ",
        "introduction": "A (Jordan or Einstein frame-) bounce~\\cite{bounce} is predicted to have taken place in some string -- or quantum gravity -- motivated early universe models, as \\eg the pre big bang~\\cite{PBB} and ekpyrotic cases~\\cite{ekp,noekp}. Those models are often presented as plausible challengers to the otherwise well-confirmed inflation paradigm, and must therefore be able to provide an almost scale-invariant perturbation spectrum.  The difficulty resides in the propagation of the spectrum, usually produced in the contracting phase, through the bounce, as the latter often involves high curvatures and thus knowledge of at least some non-linear terms in the underlying theory.  Even assuming that the perturbations can in fact somehow propagate through, \\ie supposing the first order to remain meaningful all along, a point quite debatable in itself (see \\eg Lyth in Refs.~\\cite{noekp}), it is necessary to impose a specific model to examine the validity of the usual assumptions that are made, either explicitly or implicitly, in the abovementionned frameworks.  This has been done in Refs.~\\cite{MPD,MPL} in which we used pure general relativity, a Friedmann-Robertson-Walker metric with positively curved spatial sections and a scalar field $\\varphi$ to provide the matter content and lead the dynamics of the scale factor $a(\\eta)$ as a function of conformal time $\\eta$. We have since received inquiries about some seemingly unexpected features of the results, such as an apparently vanishing-determinant transition matrix, and we were questioned about the unitarity of the evolution.  In this brief report we clarify these issues by showing that, since the evolution equation can be understood as a time-independent Schr\\\"odinger equation, its solutions can be interpreted as transmission and reflection coefficients, which ensures the conservation of probability, or, stated differentely, the unitarity of the $S-$matrix that can be built out of the transition. Note however that these results stem from the mathematical properties of the Schr\\\"odinger equation, although the actual interpretation of the underlying classical perturbed Einstein equations does involve neither any physical wavefunction nor any probabilistic meaning. In particular, this should not be confused with the effective quantization of the physical modes, \\ie the second quantization of the semi-classical effective field theory which underpins inflationary structure formation scenarios.  In addition, we show that the transition matrix reflects the underlying invariance under time reversal. Finally, we also demonstrate that the vanishing of the determinant signaled before was to be expected and turns out in fact to provide a necessary consistency check of the calculations performed in Ref.~\\cite{MPD}.  \\begin{figure*}[th] \\includegraphics[width=7cm]{potx.ps}\\hskip1cm \\includegraphics[width=7cm]{poteta.ps} \\caption{The analogy and differences between the wavefunction across a potential in ordinary quantum mechanics and primordial perturbations through a bounce. In both cases, the total current is conserved.} \\label{fig:potentials} \\end{figure*} ",
        "conclusions": "\\vspace{-0.2cm}  We have presented a brief note on the propagation of primordial perturbations through a bouncing phase in the early universe to interpret this evolution in terms of either transmission and reflexion coefficients across the potential of a non-relativistic Schr\\\"odinger-like equation or in terms of the corresponding $S-$matrix that can be constructed out of the model.  We have demonstrated that the results obtained in Refs.~\\cite{MPD,MPL} imply that $\\RB+\\TB=1$ and that the $S-$matrix is unitary, thereby checking the mathematical consistency of those results. Finally, we have also emphasized that the structure of the evolution equation is invariant under time reversal."
    },
    "0403/hep-th0403203_arXiv.txt": {
        "abstract": " ",
        "introduction": "Inflation is an attractive idea that solves many important problems in cosmology. It is also in good agreement with current observational data. It is therefore important to understand if inflation can arise in string theory.     Despite several attempts no satisfactory model of inflation in string theory has been constructed as yet. This problem is closely tied to the issue of moduli stabilization. It is well known that a very flat potential is required for inflation. There are many light fields, called moduli, in string theory which to first approximation have an exactly flat potential. So it might seem at first that a slowly enough varying potential can be easily generated. However, any attempt to generate such a potential typically runs into difficulty. One finds that some very unstable direction develops, along which the potential descends much too rapidly. As a result, the required flatness condition is not met and inflation is not realized. Recently, there has been considerable progress in moduli stabilization. So one can be hopeful that these obstacles will be overcome.  In this paper we outline the construction of a model for inflation in string theory. Our construction is based on the recent developments in moduli stabilization. The setting is warped flux compactifications of type IIB string theory, in the presence of D3-branes and anti-D3-branes. Using fluxes we fix all the complex structure moduli of the Calabi-Yau space and the dilaton-axion \\cite{GKP}. We also use a generic superpotential of the kind which arises due to gaugino condensation to stabilize the volume modulus \\cite{KKLT} (KKLT). Our discussion of inflation is closely related to the recent attempt in \\cite{KKLMMT} (KKLMMT). In particular the inflaton in our model is a D3-brane modulus, i.e., a scalar field which corresponds to the location of the D3-brane in the internal Calabi-Yau space. In \\cite{KKLMMT} it was argued that a superpotential of the kind mentioned above leads to the D3-brane moduli acquiring a mass, $m^2 =1/6 R$, where $R$ is the curvature scalar. This mass is too big and ruins the required flatness of the potential.  We explore a small twist on this story here. Consider a Calabi-Yau space with a $Z_2$ symmetry\\footnote{This $Z_2$ symmetry need not be the one involved in the orientifold action.}. A mobile D3-brane is located in the vicinity of the $Z_2$ symmetric point and experiences an attractive Coulomb force due to two symmetrically located anti-D3-branes. In such a situation we show that by adjusting the fluxes and the brane-anti-brane separation, the Coulomb attraction can nearly cancel the effect of the curvature induced mass mentioned above. As a result the $Z_2$ symmetric point turns into a maximum of the potential. The near cancellation results in a broad maximum, with $|m^2|/H^2 \\ll 1$, where $H$, is the Hubble scale corresponding to the height of the potential at the maximum. By starting close enough to such a  maximum the required conditions for slow-roll inflation can then all be met.  The observational constraints, especially  the scale of  density perturbations, impose  stringent restrictions on the model. We show that the constraints can be met and  determine the important scales in the compactification to within an order of magnitude\\footnote{More accurately, our estimates of some of the important energy   scales involved is uncertain by factors of order unity. Up to these uncertainties we show that the constraints can be met.}. A robust feature of our model, independent of many details,  is that the  scale of inflation is low. The Hubble scale, $H$, is of order $10^{10}$ GeV, which corresponds to a cosmological constant of order $10^{14}$ GeV. Thus the production of tensor perturbations is highly suppressed. The observation of gravity waves by the Planck experiment would therefore rule out this model.  A more careful study of whether all the constraints can be met  will require  concrete constructions of the model in string theory and   is left for the future. The non-perturbative superpotential we evoke, and the assumption that the full potential can be obtained by adding the brane-anti-brane interaction to the term coming from the superpotential\\footnote{This assumption would be correct if the brane-anti-brane interaction arise from a $D$-term. Evidence in support of this has been found in \\cite{Dvali:2003zh,Binetruy:2004hh}. We thank S. Kachru for bringing these references to our attention.} also needs to be studied further. The last two issues are common to many KKLT type constructions.   As a model for inflation our construction is incomplete in three ways. We have not addressed how inflation ends, how it begins, and how the standard model can be incorporated in it. Ending inflation successfully requires adequate reheating. This  depends on how the standard model is incorporated. In terms of beginning inflation, it could be that the model does not depend sensitively on initial conditions. It has been argued that a broad maximum of the kind in this model gives rise to eternal inflation. Regions where quantum fluctuations have driven the inflaton to the top of the potential hill grow exponentially more rapidly and soon dominate the universe, regardless of initial conditions. The inflationary epoch discussed in this paper then arises when fluctuations cause the scalar field to descend far enough from the top so that the classical evolution becomes dominant. This is an appealing picture but it needs to be understood better. We leave these issues for the future.   The important features of this model are quite general. They  essentially depend only   on the existence of a broad maximum, with $|m|^2/H^2 \\ll 1$, and are independent of most details. For example, we have emphasized the role of a $Z_2$ symmetry above. But the idea works more generally, even when there is no such symmetry, for a D3-brane located between  two appropriately positioned anti-D3-branes. More generally this construction can be viewed as an existence proof for broad maxima in the landscape of string theory. It seems reasonable to believe that there are many such maxima, with the complex structure moduli or K\\\"ahler moduli also playing the role of the inflaton\\footnote{S.P.T. thanks M. Douglas for emphasizing this point.}. The inflationary parameters, like the scale of density perturbations, or the tilt in the spectrum of scalar perturbations, probably take many different values at these maxima, most of which will not agree with observation. Further progress in moduli stabilization will allow us to test this grim possibility.   This paper is organized as follows. Our  basic set-up is discussed in section 2.1. The positive mass for the brane moduli due to the curvature coupling, is reviewed in section 2.2. In section 2.3, we discuss the  potential for a D3-brane located between two  symmetrically placed anti-D3-branes, and show that a broad maximum can arise. The resulting inflationary scenario is discussed in section 2.4. The constraints on the compactification which arise are analyzed in section 2.5. We close with an extended discussion in section 3.   Before proceeding we should comment on some of the relevant literature. The idea of brane inflation was first discussed by \\cite{DT}. Other related papers are  \\cite{others}, \\cite{angles}. In particular, while not worrying about moduli stabilization, \\cite{BMNQ} showed that a flat potential could be obtained by considered symmetrically positioned branes. For a review, see \\cite{Quevedo}. Standard textbook references for inflation are \\cite{Linde} and \\cite{Liddle}, see also \\cite{LythRiotto}. Some relevant references for moduli stabilization are \\cite{GVW} and \\cite{GKP} for general framework, and \\cite{KSTetal} for some  examples. For recent progress towards meeting the conditions of the KKLT construction, see \\cite{GKTT}. \\cite{EGQ} explores the KKLT construction further, \\cite{BKQt} considers inducing the anti-D3-brane charge on D7-branes. A variant of the KKLT scenario which does not require the anti-D3-brane is \\cite{Saltman:2004sn}. An investigation of de Sitter vacua using F-term potentials and additional light moduli is in \\cite{Brustein:2004xn}. A recent attempt to overcome the problems faced in KKLMMT involves the use of a shift symmetry \\cite{shiftsymmetry}. It would be nice to see if shift symmetry is present in Calabi-Yau orientifolds or their F-theory generalizations, which are required for controlled stabilization of all moduli and preserve only ${\\cal N}=1$ supersymmetry. For an  attempt in the context of string theory to use higher derivative terms for inflation see  \\cite{ST}. Inflation  with a quadratic potential of the kind we obtain here was studied earlier in \\cite{German:2001tz}, which  arrived at similar conclusions  about the low energy scale during inflation etc\\footnote{We thank S. Sarkar for bringing this paper to our attention.}. Two related papers on brane inflation appeared  while our manuscript was being readied. \\cite{Burgess:2004kv} uses additional D-terms obtained by adding the Standard Model fields in the KKLMMT set up and argues numerically that inflation can be obtained, and \\cite{DeWolfe:2004qx} explores the possibility of inflation in the dynamics of more than one anti-brane in a K-S throat.   Our conventions are as follows. $M_{10}$  refers to  the ten-dimensional Planck scale. It is related to the string scale, $\\alpha^{'}$, and the string coupling, $g_s$, by ${1 \\over M_{10}^8}= {1\\over 2}(2\\pi)^7(\\alpha^{'})^{4}g_s^2.$ $M_{Pl}$  refers to  the four-dimensional Planck scale. It is  defined by $M_{Pl}^2={1 \\over 8 \\pi G_N}$ and satisfies the relation,  $M_{Pl}^2=M_{10}^8 L^6$, where $L^6$ is the volume of the six-dimensional internal space. Finally, the tension of the D3-brane is given by $T_3={1\\over (2\\pi)^3 (\\alpha^{'})^2 g_s}$. ",
        "conclusions": "{\\it 1. Putting the Proposal on Firmer Footing:}  Some important issues need to be studied further to put this model on a firmer footing. Two of these are common to many  KKLT type constructions.  First, non-perturbative effects, like gaugino condensation,  responsible for volume stabilization need to be understood better. Second, in deriving the full potential we  assumed that the brane-anti-brane potential can be simply added to the contribution coming from the superpotential. This would be true if the former arises through a $D$-term. For evidence in support of this see \\cite{Dvali:2003zh,Binetruy:2004hh}. Finally, as mentioned above, explicit examples of Calabi-Yau spaces need to be constructed, where all the constraints discussed above can be met simultaneously. For all these reasons the construction discussed in this paper should be thought of as a proposal rather than a concrete model at the moment.  For recent progress towards meeting some of the requirements of the KKLT construction see, \\cite{GKTT}.  {\\it 2. Building a More Complete Model:}  As a model of inflation our construction is incomplete in three important ways. We have not addressed how inflation ends, and how it begins. And we have not asked if the standard model can be satisfactorily incorporated in this construction. Let us briefly comment on some of these issues here.   {\\it 2a)  Ending Inflation:} This model is in fact an example of hybrid inflation. When the brane-anti-brane separation gets to be of order the string scale, a tachyon develops. What happens next is currently a matter of active study. See for example \\cite{Sen:2002nu,Cline:2002it,Lambert:2003zr,Gaiotto:2003rm,Sen:2003xs,Sen:2003mv}. It seems reasonable to believe that the energy in the brane-anti-brane pair is eventually converted to light closed string modes like the graviton. Whether this energy can be efficiently transferred to the standard model degrees of freedom, reheating the universe satisfactorily, is   dependent on how the standard model couples to the degrees of freedom  in the inflationary  throat. This is tied to another incomplete aspect of the model mentioned above, namely how the standard model is incorporated in it, and needs to be studied further.    Cosmic strings, both of D and F type will be produced in this model at the end of inflation due to D3 anti-D3-brane annihilation \\cite{Sen:1999mg}, \\cite{Tyestrings}. If these strings are metastable \\cite{Copeland:2003bj}, their tension must meet the condition, $G_N T_1 \\le 10^{-7}$, to avoid generating unacceptably large  anisotropies. A D1-brane at the bottom of the K-S throat has tension,  $T_{D1}={1\\over 2\\pi g_s \\alpha^{'}} Z^2$, where the $Z$ dependence arises due to the redshift. The tension of the anti-D3 brane is given by $T_3 Z^4$ and is of order the vacuum energy during inflation $(2 \\times 10^{14} \\hspace{1mm} {\\rm GeV})^4$. Using the relation $T_{D1}=\\sqrt{2\\pi T_3 Z^4 / g_s}$, $T_{F1}=\\sqrt{2\\pi T_3 Z^4  g_s}$, we get $G_N \\sqrt{T_{D1} T_{F1}} \\simeq  6.5 \\times 10^{-10}$. This is lower than the bound mentioned above. Interestingly,  future observations might be sensitive to  such low values of the tension, \\cite{Copeland:2003bj}, \\cite{gwaves}.   A small positive cosmological constant is needed after the end of inflation to account for the acceleration of the present epoch. This could arise as follows. The brane is drawn towards one of the throats and annihilates the anti-brane at the bottom of it. But the anti-brane in the second throat survives after inflation ends. The final resulting cosmological constant will get a contribution from this anti-brane's tension and the superpotential terms, eq.~(\\ref{potenr}). These can  nearly  cancel leaving behind a small positive value.   {\\it 2b) The Initial Conditions and Eternal Inflation}  The scale of inflation in this model, $\\Lambda \\sim 10^{14}$ GeV, is considerably smaller than the ten-dimensional Planck scale, and   makes the question of initial conditions all the more important.  We can only offer a few speculations about this here. It is possible  that the universe did not begin in  a big bang, with temperatures of order the Hagedorn temperature.  Instead as has been suggested in \\cite{Vilenkin,lindeeter}, it might have begun by a tunneling event in the vicinity of the maximum of the potential.  It could also be that the question of initial conditions is not very significant in this model. It has been argued that  a potential with a broad maximum, of the kind we have here, will give rise to eternal inflation. If the inflaton is close enough to the top of the potential of eq.~(\\ref{poti}) with $\\phi < \\phi_c \\sim H^3/m^2 $, quantum fluctuations can  drive it up the hill at a faster rate than the classical gradient term allows it to descend. Since the regions where the cosmological constant is bigger also grow exponentially more rapidly, soon the universe will be dominated by regions where the inflaton is at the top of the hill, making the initial conditions irrelevant. The observed universe in this picture would arise when fluctuations cause the inflaton  to descend far enough from the top so that the classical evolution discussed in section 2.4, becomes valid. This is an attractive and fairly plausible picture, but it is somewhat speculative at the moment and needs to be understood better.   {\\it 3. Generalizations of the Model:}  {\\it 3a) Asymmetric Throats:}  As was mentioned towards the end of section 2.4, the inflationary scenario implemented in this note is mainly dependent on the small curvature at the maximum and independent of many details of the model. For example, we have considered the $Z_2$ symmetric model in our discussion above, but the $Z_2$ symmetry is not essential. In general, the curvature couplings discussed in section 2.2 will result in a potential for the inflaton. Suppose this potential has a minimum at some point in the compactification. Two K-S-like throats containing anti-branes at the bottom, not necessarily symmetrically located, whose distance from the minimum is adjusted relative to the red-shift factors  and the number of anti-branes at the bottom of the throats, can then nearly cancel the resulting mass giving rise to a maximum with a small curvature. The resulting inflationary scenario will then be qualitatively unchanged, in particular the scale of inflation will continue to be quite low.  {\\it 3b) More Symmetry:}  Conversely we can have situations with higher symmetry, with more than two throats symmetrically located with respect to a point of enhanced symmetry. The term $({6r_{1m} r_{1n} \\over r_1^2} -\\delta_{mn})$ in eq.~(\\ref{eintpot}), would then be replaced by the quadrapole moment of the resulting configuration, $Q_{mn}$. As long as $Q_{mn}$ does not vanish, it will have some positive eigenvalues and a similar discussion can go through. We note here that higher symmetries, e.g., a $Z_4$ symmetry, could lead to positive  eigenvalues of $Q_{mn}$ along more than one direction. As a result additional directions would be unstable in the potential and  during inflation the brane can  roll along these additional directions as well.  {\\it 4)  The Landscape and Inflation:}  {\\it 4a) Broad Maxima:}  It seems quite likely that maximum of the kind found in this paper, which are broad with a small curvature, are more generally present in the string theory landscape \\cite{Susskind}. The unstable  directions could be complex structure moduli or K\\\"ahler moduli as well. This possibility can be examined further as our understanding of moduli stabilization progresses. For complex structure moduli in the KKLT type constructions, one needs a better understanding of the non-perturbative corrections that are involved\\footnote{In particular these contributions go like $W_{NP}=A(z_i) e^{-a\\rho}$. The prefactor,  $A(z_i)$, depends on the complex structure moduli, and one needs to  know  this dependence.}. As far as K\\\"ahler moduli are concerned, it seems difficult to get the required small mass with only one K\\\"ahler modulus,  since the canonically normalized field is logarithm of the volume. But this might be possible with more than one K\\\"ahler moduli, again this requires further developments in our understanding of moduli stabilization.  Generically broad maxima in the landscape, which agree with observational constraints, will have a  low scale of inflation\\footnote{To see this, let us write the potential as follows, \\be \\label{genrpot} V=\\Lambda^4 f(\\phi/M_{Pl}) = \\Lambda^4 - m^2 \\phi^2 + {\\cal{O}}(1) {\\Lambda^4 \\over M_{Pl}^4} \\phi^4 + \\cdots \\,, \\ee where the ellipses denote higher powers of $(\\phi/M_{Pl})$. Typically the slow roll conditions stop holding when the quartic terms come into play, i.e., when  $\\phi \\sim M_{Pl} {m \\over H} $. Imposing $\\eta ={m^2/3H^2} \\sim 10^{-2}$, and the observed value of the density perturbations, eq.~(\\ref{sdp}), then leads to $H \\sim  10^{12} $ GeV. The resulting  power in gravity waves is ${\\cal{P}}_{grav} \\sim 10^{-14}$, which is quite small even for future observations. In specific models, like the one considered in this paper, inflation can end earlier and the resulting Hubble scale can be lower.} $H \\lesssim 10^{12} {\\rm GeV}$. If inflation arises due to such a maximum in the landscape,  gravity waves are highly suppressed and will be difficult to observe in future experiments.   {\\it 4b)  The Landscape and Various Inflationary Scenarios:}  So far in our discussion we have used the observational data and deduced various constraints on the compactification. Once explicit constructions become possible, we can turn this around and ask instead whether string theory makes any predictions about the observational data. While it is premature to speculate on this of course, it seems likely that in general the landscape will have many broad maxima of various different types, and the resulting values for the Hubble scale, $H$ and the mass, $m^2/H^2$, will take various different values, resulting in many different possibilities for the scale of density perturbations and the tilt. Most of these will probably not agree with observation.      \\bigskip \\goodbreak \\centerline{\\bf Acknowledgments} \\noindent We are grateful to Prasanta Tripathy for collaboration at early stages. S.P.T. acknowledges conversations with Alan Guth and other participants during the Cambridge workshop on Brane World Cosmology. He also acknowledges discussions with participants of the String Theory and Cosmology workshop at KITP, Santa Barbara, and with the participants of the Coorg Conference on CMBR. N.I. thanks people at Harish-Chandra Research Institute in Allahabad for nice hospitality.  S.P.T. acknowledges support from the Swarnajayanti Fellowship, D.S.T., Govt. of India. This research is supported by the generosity of the people of India, so we deeply thank  them."
    },
    "0403/hep-ph0403247_arXiv.txt": {
        "abstract": "The seesaw model of neutrinos might explain the size, age, flatness and near-homogeneity of the Universe via sneutrino inflation, as well as explaining the origin of matter via leptogenesis. The sneutrino inflation hypothesis makes specific, testable predictions for cosmic microwave background observables, which are compatible with the first release of data from WMAP, and for flavour-violating charged-lepton decays. In particular, $\\mu \\to e \\gamma$ should occur with a branching ratio very close to the present experimental upper limit, whilst $\\tau \\to \\mu \\gamma$ and $e \\gamma$ should occur further below the present limits. \\vspace{1pc} ",
        "introduction": "What is the connection between the data on the cosmic microwave background (CMB) from WMAP~\\cite{WMAP} and searches for $\\mu \\to e \\gamma$ and other charged-lepton flavour-violating (LFV) processes? One may be provided by sneutrino inflation, an idea first proposed by Murayama-san, Suzuki-san, Yanagida-san and Yokoyama-san~\\cite{MSYY1,MSYY2}. This appealing idea languished in obscurity for over a decade, for reasons that are unknown to me. Personally, although I have always been attracted towards inflationary cosmology, I have been reluctant to embrace it wholeheartedly, largely because of the absence of a convincing candidate for the inflaton field. Only recently did it dawn upon me that the supersymmetric spin-zero partner of one of the heavy singlet neutrinos in the seesaw might be a suitable candidate. My interest was then stimulated by the first release of data from the WMAP satellite~\\cite{WMAP}, which were consistent with many key inflationary predictions, such as the Gaussian nature and almost scale invariance of the primordial density perturbations. Moreover, the WMAP data excluded many rival inflationary models, such as those with a simple quartic potential, for example.  Last year, Raidal, Yanagida-san and I~\\cite{ERY} revived the sneutrino inflation idea, calculated several inflationary observables such as the scalar spectral index (tilt) and the magnitude of the tensor perturbations relative to the scalar modes, showed they were consistent with the WMAP data, used the model to constrain the seesaw parameters, and finally calculated LFV processes showing, for example, that $\\mu \\to e \\gamma$ might occur not far below the present experimental limit. More recently, Chankowski, Pokorski, Raidal, Turzynski and I~\\cite{CEPRT} have explored in more detail the sneutrino inflation predictions for LFV processes. This talk is based on these two papers.  First, however, I summarize the basic features of inflation~\\cite{Inflation}. Then I recall the 18 parameters of the minimal three-generation seesaw model~\\cite{CI}, and how they appear in low-energy LFV observables as well as neutrino oscillations and leptogenesis~\\cite{EHLR,ERaidal}. Then I compare the predictions of sneutrino inflation~\\cite{ERY} with WMAP data~\\cite{WMAP} and recall that the gravitino problem~\\cite{gravitino} favours non-thermal scenarios for leptogenesis~\\cite{ERY}. Finally, I dissect~\\cite{CEPRT} the predictions of sneutrino inflation for various LFV processes~\\cite{ERY}, showing how they follow from the (near) decoupling of the sneutrino inflaton, which is motivated by the gravitino problem and some approaches to neutrino masses within the seesaw model~\\cite{Chank}. ",
        "conclusions": "The contemporary techno-folk singer Moby tells us that `We are all made of stars'. Perhaps this is a reference to astrophysical nucleosynthesis? In contrast, my main message in this talk is that `We are all made of neutrinos', not just in the sense of leptogenesis, but also in the sense that the overall size, age, flatness and entropy are all due to neutrinos, via cosmological inflation. Moreover, this hypothesis enables some testable predictions to made for the radiative charged-lepton decays $\\mu \\to e \\gamma$, $\\tau \\to \\mu \\gamma$ and $\\tau \\to e \\gamma$.  \\newpage"
    },
    "0403/astro-ph0403544_arXiv.txt": {
        "abstract": "We present out  results on the cross-correlation of Lyman break galaxies (LBGs) around three damped \\Ly\\ absorbers (DLAs) at $z_{\\rm abs}\\simeq 3$ from deep ($\\mu_{I,\\rm AB}(sky)\\simeq 27.6$mag arcsec$^{-2}$)   $UBVI$ KPNO 4m/MOSAIC images. The large area of the MOSAIC images, 0.31  $\\deg^2$ or $\\sim 65\\times 65h_{71}^{-1}$~Mpc co-moving at redshift $z=3$, allows us to probe the clustering of LBGs on scales up to 20~Mpc co-moving. Our survey covers a total of 1 deg$^2$ and contains $\\sim$3,000 LBGs with photometric redshifts between $2.8$ and $3.5$. Using the redshift likelihood distributions with $m_I$ as a prior, we selected LBGs  within a redshift slice of width $W_z=0.15$ (corresponding to $\\sigma_z$, the uncertainty in photometric redshifts)  centered on the redshift of the absorbers. Within that redshift slice, we find that the DLA-LBG cross-correlation $w_{dg}$ is $w_{dg}=(\\crossPdla\\pm \\crossPdlaerr) \\times w_{gg}$, where $w_{gg}$ is the LBG auto-correlation. This corresponds to a correlation length of $r_o=5\\pm4.5h^{-1}$ (co-moving) (or $r_o=7\\pm 6.8h_{71}^{-1}$~Mpc). The cross-correlation is most significant  on scales $5-10$~Mpc. Through Monte Carlo simulations, we find that $w_{dg}$ is significantly greater than zero at the $>95$\\% level. In three other redshift slices that do not contain a DLA, we do not find any evidence of clustering. A larger sample will enable us to discriminate between $w_{dg}/w_{gg}$ $<1$ or $w_{dg}/w_{gg}$  $>1$, i.e  to test whether DLA halos are more or less massive than LBG halos. ",
        "introduction": "QSO absorption lines,   including damped Ly-$\\alpha$ absorbers (DLAs), and Lyman break galaxies (LBGs)   are currently our two major sources of information on high redshift galaxies.  After more than two decades of study, the exact nature and detailed characteristics of  damped absorbers remain unexplained. Here, we seek to constrain the properties of DLA halos using LBGs \\citep{steidel99} as tracers of large scale structure.  DLAs contain the largest reservoir of neutral hydrogen (\\HI) at high redshifts \\citep[e.g.][]{lanzetta91,lanzetta95,ellison01b}. They contain more neutral \\HI\\ than all the absorbers in the Ly-alpha forest combined. Morever, the amount of \\HI\\ in  DLAs at high redshifts  corresponds to the amount of \\HI\\ in stars today at $z=0$: \\citet{ellison01b} find  $\\Omega_{HI}(z=3)=10^{-2.6}$, while \\citet{bell03} measure  $\\Omega_*(z=0)=10^{-2.56}$ (both numbers are for $h=0.65$). These facts led \\citet{wolfe86} to put forward the `disk hypothesis', namely that DLAs are large thick gaseous disk galaxies. Despite the numerous observations directed at DLAs in the past  decade \\citep[e.g. imaging studies such as][]{moller93,lowenthal95, steidel94,steidel95,lebrun97,bunker99,fynbo99,kulkarni00,pettini00,rao00,bouche01,moller02}, this hypothesis has been debated and the role of DLAs in galaxy formation is still not understood \\citep[see the more exhaustive summary of][]{pettini03}.  Hints to the nature of DLAs are given by numerical \\citep{katz96,haehnelt00,gardner01,nagamine03} and semi-analytical \\citep{kauffmann96,mo99,okoshi03} simulations of galaxy formation  at high-redshifts.  All, these simulations indicate that  DLAs are in majority faint (sub-L$^*$) in small dark matter halos $V_c\\ll100$\\kms. Based on cross-section arguments, \\citet{fynbo99}, \\citet{moller02}, and \\citet{schaye01} arrived to the same conclusions. From the chemical evolution point of view, \\citep{matteucci97,jimenez99,boissier03} argued that DLAs  are  caused by gas rich low surface brightness dwarf galaxies, as seen locally in at least one case \\citep{bowen01}.  However, DLAs show asymmetric profiles of their  high ionization species \\citep{prochaska97,ledoux98}. This has been used to argue  that DLAs are, in fact, due to thick massive rotating disks \\citep{wolfe86,wolfe95,prochaska97}. But others, e.g. \\citet{maller00}, \\citet{mcdonald99} and \\citet{haehnelt00}, showed that a large range of morphologies can reproduce the observed kinematics: DLAs can arise from the combined effect of a massive central galaxy and a number of smaller satellites  or filaments. In fact, cold gas accretion along filaments could be an important mechanism, especially at high redshifts \\citep{keres03}.  Whether or not DLAs are indeed massive will lead to different  clustering properties of the galaxies around them. In  hierarchical galaxy formation models \\citep[e.g.][]{mo96,mo02},  this clustering   yields a measurement of the dark matter halo mass associated with DLAs relative to that of the galaxies used as tracers of the large scale structure. In particular, if the galaxies are less (more)  correlated with the DLAs than with themselves, this will imply that the halos of DLAs are less (more) massive than those of the galaxies. Here, we use $z\\simeq3$ LBGs  \\citep{steidel99}   as tracers of the large scale structure.   In analyses similar to that presented here, \\citet{gawiser01} found no clustering of galaxies around one single DLA at $z=4$ towards BR 0951-04, and \\citet{adelberger03} found a lack of galaxies near  four DLAs (they found two within a cylinder of radius of $5.7h^{-1}$Mpc and depth  $W_z<0.025$  whereas $\\sim 6$ were expected if the cross-correlation is the same as the galaxy auto-correlation). They argued that this is evidence that DLAs and LBGs ``do not reside in the same parts of the universe''. It is important to note that both of these studies were not sensitive to scales larger than 5$h^{-1}$~Mpc because   because of the small field of view available.  Other studies, however, have pointed to an over-density of galaxies near DLAs. \\citet{wolfe93} combined several studies of \\lya\\ emitters around DLAs and found evidence for a correlation between emitters and DLAs at a mean redshift $\\overline z=2.6$, significant at $>99$\\%. \\citet{francis93} reported the discovery of super clustering of sub-DLAs at $z\\sim 2.4$ and $z\\sim2.9$: a total of four \\HI\\ clouds are seen towards a QSO pair separated by 8\\arcmin, each being at the same velocity. Recent results from narrow-band imaging of the Francis \\& Hewett field shows that spectroscopically confirmed \\lya\\ emitters are clustered at the redshift of the strongest \\HI\\ cloud at $z=2.9$ ($\\log N_{\\HI}=20.9$) towards Q2138-4427 \\citep{fynbo03}.  \\citet{RocheN_00a} identified eight Lyman-alpha emitting galaxies near the DLA at $z=2.3$ towards PHL 957 in addition to the previously discovered Coup Fourr\\'e galaxy \\citet{lowenthal91}, implying the presence of a group, filament, or proto-cluster associated with the DLA. Other evidence of clustering include the work of \\citet{ellison01} and \\citet{dodorico02}. \\citet{ellison01} found that the DLA at $z_{\\rm abs}=3.37$ towards Q0201+1120 is part of a concentration of matter that includes at least four galaxies (including the DLA) over transverse scales greater than $5 h^{-1}$~Mpc. \\citet{dodorico02} showed  that out of ten DLAs in QSO pairs, five are matching systems within 1000\\kms. They concluded that this result indicates a highly significant over-density of strong  absorption systems over separation lengths from $\\sim 1$ to $8 h^{-1}$ Mpc.   Despite numerous attempts to reproduce the properties of DLAs   \\citep[e.g][]{katz96,mcdonald99,mo99,gardner01,nagamine03,okoshi03} in numerical simulations or semi-analytical models,  there are no predictions of the clustering of galaxies around DLAs. In \\citet[][in preparation]{bouche03b}, we discuss the cross-correlation of LBGs around DLAs in Smoothed Particle Hydrodynamical (SPH) numerical simulations. There, we find that the DLA-LBG cross-correlation is weaker than the LBG auto-correlation by approximately 30\\%, indicating that DLAs reside in less massive dark matter halos than those of LBGs.  In \\citet{bouche03}, we presented a preliminary measurement of clustering of LBGs around a DLA at $z\\sim3$ towards APM 08279+5255, based on part of this data set. We detected an over-density of $\\Sigma/\\Sigma_g\\simeq 3$ at the 95\\% level on scales $2.5<r_\\theta<5$~Mpc, implying that at least some DLAs reside in dense environments. In this paper, we expand on that study and present results on the DLA-LBG cross-correlation in our full survey of three fields.  In section~\\ref{section:data},  we present the imaging data used here, our completeness limits and  pre-selection of high-redshift galaxy candidates. We used photometric redshifts to select our $z=3$ LBG candidates as discussed in section~\\ref{section:sample}. Our clustering  analysis is shown in section~\\ref{section:cluster}. The results on the DLA-LBG cross-correlation  are shown in  section~\\ref{section:results}, and discussed in section~\\ref{section:discussion}. Section~\\ref{section:conclusions} contains our conclusions.  Throughout this paper, we adopt $\\Omega_M=0.3$, $\\Omega_\\Lambda=0.7$ and $H_o=100 h$~\\kms~Mpc$^{-1}$; thus, at $z\\sim3$, $1\\arcsec$ corresponds to  $\\sim 21.5h^{-1}$~kpc and $1\\arcmin$ to $\\sim 1.29h^{-1}$~Mpc, both  {\\em co-moving}. At that redshift, $H(z)\\sim 4.46 H_o$, so $\\delta z=0.1$ corresponds to $67 h^{-1}$~Mpc in co-moving coordinates~\\footnote{Using the latest cosmological parameters from WMAP ($\\Omega_M=0.268$, $\\Omega_\\Lambda=0.728$, $h=0.71$) changes these numbers by $\\sim$ 3\\%.}. ",
        "conclusions": "\\label{section:conclusions}  Based  on deep ($\\mu_{I,\\rm AB}(sky) \\simeq 27.6$~mag~arcsec$^{-2}$) wide-field images (0.31$\\deg^2$ or $\\sim65\\times 65 h_{71}^{-1}$~Mpc co-moving at redshift $z=3$) around three DLAs, we  identify   LBG candidates  brighter than $I_{AB}=\\complAB$mag using photometric redshift techniques that included the   $I$ magnitude as a prior estimate in addition to the colors.  From the redshift likelihood distributions, we selected LBG galaxies within a redshift slice of width $W_z=0.15(\\simeq\\sigma_z)$ centered on the redshift of the  DLAs $z_{\\rm abs}$. Within that slice, we cross-correlated the LBGs with the position of the DLAs and found that \\begin{itemize} \\item the amplitude of the DLA-LBG cross-correlation $w_{dg}$ relative to the auto-correlation $w_{gg}$ was $w_{dg}/w_{gg}\\equiv a=\\crossPdla \\pm \\crossPdlaerr$, corresponding to  $r_o=5\\pm 4.5h^{-1}$~Mpc (co-moving), \\item the amplitude of the DLA-LBG cross-correlation is  $a>0$, which is  significant at the $>95$\\% confidence level based on Monte Carlo simulations, \\item the clustering signal was not present in three redshift slices that did not contain the DLAs. \\end{itemize}  Given the uncertainty of our results, we cannot put constraints on the halo masses of DLAs and discriminate between the large disk hypothesis \\citep[e.g][]{wolfe86}, and small sub-L$^*$ hypothesis \\citep[e.g][]{maller00,haehnelt00,moller02}. Our observation of the clustering on large scales ($>4h^{-1}$~Mpc) is not inconsistent with previous clustering studies \\citep{gawiser01,adelberger03} since these were limited to small scales. In order to be able to directly compare these studies with our present results on scales $<4h^{-1}$~Mpc, a larger sample of DLAs  and multi-object spectroscopy of our LBG candidates are needed. This will enable to test whether the cross-correlation is stronger or weaker than the auto-correlation."
    },
    "0403/gr-qc0403029_arXiv.txt": {
        "abstract": "We  present a  new three-dimensional  fully  general-relativistic hydrodynamics   code    using   high-resolution   shock-capturing techniques and a  conformal traceless formulation of the Einstein equations.  Besides presenting a  thorough set of tests which the code passes  with very high accuracy, we  discuss its application to the study of  the gravitational collapse of uniformly rotating neutron stars  to Kerr black  holes.  The initial  stellar models are  modelled   as  relativistic  polytropes   which  are  either secularly  or dynamically  unstable and  with  angular velocities which range  from slow rotation  to the mass-shedding  limit.  We investigate the gravitational  collapse by carefully studying not only the  dynamics of  the matter, but  also that of  the trapped surfaces,  {\\it  i.e.}  of  both  the  apparent  {\\it and}  event horizons formed during the  collapse.  The use of these surfaces, together  with  the dynamical  horizon  framework,  allows for  a precise measurement of the black-hole mass and spin.  The ability to successfully  perform these simulations  for sufficiently long times  relies on excising  a region  of the  computational domain which  includes  the  singularity  and  is  within  the  apparent horizon.   The  dynamics of  the  collapsing  matter is  strongly influenced  by the  initial  amount of  angular  momentum in  the progenitor star  and, for  initial models with  sufficiently high angular velocities, the collapse can  lead to the formation of an unstable disc  in differential rotation.  All  of the simulations performed with uniformly rotating initial data and a polytropic or ideal-fluid equation of state show no evidence of shocks or of the presence of matter on stable orbits outside the black hole. ",
        "introduction": "\\label{sec:introduction}   The numerical investigation of gravitational collapse of rotating stellar configurations leading to black-hole formation is a long standing problem in numerical relativity. However, it is through numerical simulations in general relativity that one can hope to improve our knowledge of fundamental aspects of Einstein's theory such as the cosmic censorship hypothesis and black-hole no-hair theorems, along with that of current open issues in relativistic astrophysics research, such as the mechanism responsible for gamma-ray bursts. Furthermore, numerical simulations of stellar gravitational collapse to black holes provide a unique mean of computing the gravitational waveforms emitted in such events, believed to be among the most important sources of detectable gravitational radiation.  However, the modelling of black-hole spacetimes with collapsing matter-sources in multidimensions is one of the most formidable efforts of numerical relativity. This is due, on one hand, to the inherent difficulties and complexities of the system of equations which is to be integrated, the Einstein field equations coupled to the general-relativistic hydrodynamics equations, and, on the other hand, to the immense computational resources needed to integrate the equations in the case of three-dimensional (3D) evolutions. In addition to the practical difficulties encountered in the accurate treatment of the hydrodynamics involved in the gravitational collapse of a rotating neutron star to a black hole, the precise numerical computation of the gravitational radiation emitted in the process is particularly challenging as the energy released in gravitational waves is much smaller than the total rest-mass energy of the system.   The ability to perform long-term numerical simulations of self-gravitating systems in general relativity strongly depends on the formulation adopted for the Einstein equations. The covariant nature of these equations (the ``many-fingered time'' of relativity) leads to difficulties in constructing an appropriate coordinate representation which would allow for stable and accurate simulations.  Over the years, the standard approach has been mainly based upon the unconstrained solution of the $3+1$ ADM formulation of the field equations, which, despite large-scale and dedicated collaborations~\\cite{Cook97a, Abrahams97a, Gomez98a} has gradually been shown to lack the stability properties necessary for long-term numerical simulations. In recent years, however, a considerable effort has been invested in extending the set of ADM equations solved by including at some level the solution of the constraint equations on each spatial hypersurface~\\cite{AM:03,Schnetter03c,Bonazzola03a}, or by reformulating the ADM approach in order to achieve long-term stability (see, {\\it e.g.,}~\\cite{Lehner01a} and references therein). Building on the experience developed with lower-dimensional formulations, Nakamura, Oohara and Kojima~\\cite{Nakamura87} presented in 1987 a conformal traceless reformulation of the ADM system which subsequent authors (see, {\\it e.g.,}~\\cite{Shibata95, Baumgarte99, Shibata99e, Shibata03,Alcubierre99d, Font02c, Yoneda02a, Kawamura03}) gradually showed to be robust enough to accomplish such a goal for different classes of spacetimes, including black holes and neutron stars (both isolated and in coalescing binary systems). The most widespread version developed from this formalism, which we refer to here as the NOK formulation, was given by~\\cite{Shibata95,Baumgarte99} and is commonly referred to as the BSSN formulation.    In addition to the improvements achieved in the formulation of the field equations, successful long-term 3D evolutions of black holes in vacuum have been obtained in the last few years using excision techniques (see, {\\it e.g.,}~\\cite{Seidel92a, Brandt00, Alcubierre00a, Kidder00a, Kidder01a,Alcubierre01a, Yo02a, Laguna02, Choptuik:2003ac, Sperhake:2003fc}), although the original idea is much older~\\cite{MayWhite67}. In this approach, the spacetime region within the black-hole horizon, and so causally disconnected, can be safely ignored without affecting the evolution outside the horizon as long as suitable boundary conditions are specified at the excision surface. The simulations presented here show the applicability of excision techniques {\\it also} in non-vacuum spacetimes, namely during the collapse of rotating neutron stars to Kerr black holes. The excision technique, which is applied once the black-hole apparent horizon is found, permits to extend considerably the lifetime of the simulations, at least with the resolutions used here. This, in turn, allows for an accurate investigation of the dynamics of the trapped surfaces formed during the collapse, from which important information on the mass and spin of the black hole, as well as on the amount of energy which is lost to gravitational radiation, can be extracted. While our study was nearing completion, we have learnt that a similar approach has also been implemented with success~\\cite{Duez04}.    The presence of rotation in the collapsing stellar models requires multidimensional investigations, either in axisymmetry or in full 3D. The numerical investigations of black-hole formation (beyond spherical symmetry) started in the early 1980's with the pioneering work of Nakamura~\\cite{Nakamura81b}. He adopted the (2+1)+1 formulation of the Einstein equations in cylindrical coordinates and introduced regularity conditions to avoid divergences at coordinate singularities. Nakamura used a ``hypergeometric\" slicing condition which prevents the grid points from reaching the singularity when a black hole forms. The simulations could track the evolution of the collapse of a $ 10 M_\\odot $ ``core'' of a massive star with different amounts of rotational energy, up to the formation of a rotating black hole. However, the numerical scheme employed was not accurate enough to compute the emitted gravitational radiation. In subsequent works, Nakamura \\cite{Nakamura83} (see also \\cite{Nakamura87}) considered a configuration consisting of a neutron star with an accreting envelope, which was thought to mimic mass fall-back in a supernova explosion. Rotation and infall velocity were added to such a configuration, simulating an evolution dependent on the prescribed rotation rates and rotation laws.  Later on, in a series of papers~\\cite{Bardeen83,Stark85,stark86,stark87}, Bardeen, Stark and Piran studied the collapse of rotating relativistic polytropes to black holes, succeeding in computing the associated gravitational radiation. The gravitational field and hydrodynamics equations were formulated using the $3+1$ formalism in two spatial dimensions, using the radial gauge and a mixture of singularity-avoiding polar and maximal slicings. The initial model was a spherically symmetric relativistic polytrope of mass $M$ in equilibrium. The gravitational collapse was induced by lowering the pressure in the initial model by a prescribed (and often very large) fraction and by simultaneously adding an angular momentum distribution approximating rigid-body rotation. Their simulations showed that Kerr black-hole formation occurs only for angular momenta less than a critical value.  Furthermore, the energy $\\Delta E$ carried away through gravitational waves from the collapse to a Kerr black hole was found to be $\\Delta E/Mc^2<7\\times 10^{-4}$, the shape of the waveforms being nearly independent of the details of the collapse.  The axisymmetric codes employed in the aforementioned works adopted curvilinear coordinate systems which may lead to long-term numerical instabilities at coordinate singularities. These coordinate problems have deterred researchers from building 3D codes in spherical coordinates. Recently, a general-purpose method (called ``cartoon''), has been proposed to enforce axisymmetry in numerical codes based on Cartesian coordinates and which does not suffer from stability problems at coordinate singularities~\\cite{alcubierre00b}. It should be noted, however, that the stability properties of the cartoon method are not fully understood yet, as discussed by~\\cite{Frauendiener02}. Using this method, Shibata~\\cite{Shibata00a} investigated the effects of rotation on the criterion for prompt adiabatic collapse of rigidly and differentially rotating polytropes to a black hole, finding that the criterion for black-hole formation depends strongly on the amount of angular momentum, but only weakly on its (initial) distribution. The effects of shock heating when using a non-isentropic equation of state (EOS hereafter) are important in preventing prompt collapse to black holes in the case of large rotation rates.  More recently, Shibata~\\cite{Shibata03,Shibata03b} has performed axisymmetric simulations of the collapse of rotating supramassive neutron stars to black holes for a wide range of polytropic EOSs and with an improved implementation of the hydrodynamics solver (based on approximate Riemann solvers) with respect to the original implementation used in~\\cite{Shibata00a}. Parameterizing the ``stiffness\" of the EOS through the polytropic index $N$, the final state of the collapse is a Kerr black hole without any noticeable disc formation, when the polytropic index $N$ is in the range $2/3 \\le N \\le 2$. Based on the specific angular momentum distribution in the initial star, Shibata has estimated an upper limit to the mass of a possible disc as being less than $10^{-3}$ of the initial stellar mass~\\cite{Shibata03b}.  Unfortunately, such small masses cannot currently be confirmed with the presently-available resolutions in 3D simulations on uniform grids.   Three-dimensional, fully relativistic simulations of the collapse of supramassive uniformly rotating neutron stars to rotating black holes were presented in~\\cite{Shibata99e}. The simulations focused on $N=1$ polytropes and showed no evidence of massive disc formation or outflows. These results are in agreement with those obtained in axisymmetry~\\cite{Shibata03,Shibata03b} and with the new simulations reported by \\cite{Duez04} (both in axisymmetry and in 3D) which show that for a rapidly rotating polytrope with $J/M^2<0.9$ ($J$ being the angular momentum) all the mass falls promptly into the black hole, with no disc being formed. Hence, all existing simulations agree that massive disc formation from the collapse of neutron stars requires differential rotation, at least for a polytropic EOS with $1 \\le N \\le 2$.   Here, we present the results of new, fully 3D simulations of gravitational collapse of uniformly rotating neutron stars, both secularly and dynamically unstable, which we model as relativistic polytropes.  The angular velocities of our sample of initial models range from slow rotation to the mass-shedding limit. For the first time in such 3D simulations, we have detected the {\\it event} horizon of the forming black hole and showed that it can be used to achieve a more accurate determination of the black-hole mass and spin than it would be otherwise possible using the area of the {\\it apparent} horizon. We have also considered several other approaches to measure the properties of the newly formed Kerr black hole, including the recently proposed {\\it isolated} and {\\it dynamical-horizon} frameworks. A comparison among the different methods has indicated that the dynamical-horizon approach is simple to implement and yields estimates which are accurate and more robust than those of the equivalent methods.  The simulations are performed with a new general relativistic hydrodynamics code, the {\\tt Whisky} code, in which the Einstein and hydrodynamics equations are finite-differenced on a Cartesian grid and solved using state-of-the-art numerical schemes (a first description of the code was given in \\cite{Baiotti03a}). The code incorporates the expertise developed over the past few years in the numerical solution of the Einstein equations and of the hydrodynamics equations in a curved spacetime (see~\\cite{Alcubierre99d,Font02c}, but also~\\cite{Font03} and references therein) and is the result of a collaboration among several European Institutes~\\cite{eunetworkweb}.  As mentioned before, we have implemented in the {\\tt Whisky} code a robust excision algorithm which warrants the extension of the lifetime of the simulations far beyond the evolution times when the black holes first form. Our calculations are starting from initially axisymmetric stellar models but are performed in full 3D to allow for departures from the initial axial symmetry. Overall, our results show that the dynamics of the collapsing matter is strongly influenced by the initial amount of angular momentum in the progenitor neutron star, which, when sufficiently high, leads to the formation of an unstable flattened object. All of the simulations performed for realistic initial data and a polytropic equation of state show no evidence of shock formation preventing a prompt collapse to a black hole, nor the presence of matter on stable orbits outside the black hole. It should be remarked, however, that both of these conclusions may change if the initial stellar models are rotating differentially.  The use of numerical grids with uniform spacing and the present computational resources have placed the outer boundary of our computational box in regions of the spacetime where the gravitational waves have not yet reached their asymptotic form. As a result, the information on the gravitational waveforms that we extract through perturbative techniques~\\cite{Rupright98,Rezzolla99a} does not provide interesting information besides the obvious change in the stellar quadrupole moment. Work is now in progress to use mesh refinement techniques~\\cite{Schnetter-etal-03b} to move the outer boundary sufficiently far from the source so that important information can be extracted on the gravitational wave emission produced during the collapse. The results of these investigations will be presented in a companion paper~\\cite{Baiotti04}.   The paper is organized as follows: Section~\\ref{equations} describes the formulation we adopt for the Einstein and hydrodynamics equations, the way they are implemented in the code and a brief discussion of how the excision techniques can be employed within a hydrodynamical treatment making use of high-resolution shock-capturing (HRSC) schemes. To avoid detracting the reader's attention from the physical problem considered here, we have confined most of the technical details concerning the numerical implementation of the hydrodynamical equations to Appendix~\\ref{sec:numerical-methods}. Section~\\ref{sec:model} is therefore devoted to describing the various properties of the initial stellar models. The following two Sections, ~\\ref{sec:dynamics-fluid} and~\\ref{sec:dynamics-horizons}, present our results regarding the dynamics of the collapsing stars and of the trapped surfaces, respectively. In both cases we will be considering and comparing the dynamics of slowly and of rapidly rotating stellar models. The paper ends with Section~\\ref{sec:conclusion}, which contains a summary of the results obtained and the perspectives for further investigations. Finally, Appendix~\\ref{sec:numerical-tests} is devoted to presenting some of the tests the code passes with very high accuracy.  We here use a spacelike signature $(-,+,+,+)$ and a system of units in which $c=G=M_\\odot=1$ (unless explicitly shown otherwise for convenience). Greek indices are taken to run from 0 to 3, Latin indices from 1 to 3 and we adopt the standard convention for the summation over repeated indices. ",
        "conclusions": "\\label{sec:conclusion}  Although 3D numerical relativity has been a very active research area for several years now, there are still a number of technical issues to be addressed and physical problems to be investigated in detail. Separate progress has been made so far in obtaining long-term stable evolutions of vacuum spacetimes and of spacetimes with matter. Both of them have posed significant numerical problems because of the existence of horizons containing physical singularities, in one case, and the development of non-linear hydrodynamical phenomena such as shocks, in the other. In black-hole vacuum spacetimes, these problems have successfully been dealt with by using better suited formulations of the Einstein equations and by employing excision techniques for the regions of the spacetime containing the singularity. In spacetimes containing matter, on the other hand, sophisticated numerical techniques (such as the HRSC methods) have been employed to accurately track the dynamics of the shocks.  Here, we have combined these two different approaches by implementing excision techniques within a {\\it forming} horizon, thus following the dynamics of the matter as it accretes onto the developing black hole. We have shown that doing so allows the numerical evolution to proceed uninhibited from fully regular initial conditions of matter in equilibrium and devoid of trapped surfaces, up to a vacuum spacetime featuring an event horizon enclosing an excised physical singularity. This new important ability in numerical relativity evolutions will help in a more detailed investigation of complex astrophysical systems, such as the coalescence of neutron star binaries, considered as a prime candidate for the detection of gravitational waves, and of the collapse of stellar cores, considered as the progenitors of gamma-ray bursts.  Our hydrodynamical excision technique is implemented within a new 3D general-relativistic numerical evolution code that combines state-of-the-art numerical methods for the spacetime evolution ({\\it i.e.} the NOK formulation of the Einstein equations with Gamma-driver shift conditions) with an accurate hydrodynamical evolution employing several high-order HRSC methods. The evolution of the spacetime and of the hydrodynamics is coupled transparently through the method of lines, which allows for the straightforward implementation of various different time-integrators.  As a first astrophysical problem for this novel setup, we have here focused on the collapse of rapidly rotating relativistic stars to Kerr black holes. The stars are assumed to be in uniform rotation and dynamically unstable to axisymmetric perturbations. While the collapse of slowly rotating initial models proceeds with the matter remaining nearly uniformly rotating, the dynamics is shown to be very different in the case of initial models rotating near the mass-shedding limit, for which strong differential rotation develops. Although the stars become highly flattened during collapse, attaining a disc-like shape, the collapse cannot be halted because the specific angular momentum is not sufficient for a stable disc to form. Instead, the matter in the disc spirals towards the black hole and angular momentum is transferred inward to produce a spinning black hole.  Several different approaches have been employed to compute the mass and angular momentum of the newly formed Kerr black hole. Besides more traditional methods involving the measure of the geometrical properties of the apparent {\\it and} event horizons, we have fitted the oscillations of the perturbed Kerr black hole to specific quasi-normal modes obtained by linear perturbation theory. In addition, we have also considered the recently proposed {\\it isolated} and {\\it dynamical} horizon frameworks, finding it to be simple to implement and yielding estimates which are accurate and more robust than those of other methods. This variety of approaches has allowed for the determination of both the mass and angular momentum of the black hole with an accuracy unprecedented for a 3D simulation. These measures, in turn, have allowed us to set upper limits on the energy and angular momentum that could be lost during the collapse in the form of gravitational radiation.  Work using mesh-refinement techniques is already in progress to extract more precise information and waveforms for the gravitational radiation recorded at large distances from the collapsing stars. This aspect of the gravitational collapse has not yet been considered in full 3D simulations and will be reported in a forthcoming paper~\\cite{Baiotti04}. Finally, all of the techniques discussed here will also be applied to the study of the collapse of differentially rotating stars, governed by more realistic and non-isentropic EOSs. The expectation is that initial data with $J/M^2 \\gtrsim 1 $ can be constructed in this case, whose collapse could lead to the formation of a massive disc orbiting around the newly formed Kerr black hole~\\cite{Shibata03,Shibata03b,Duez04}.   \\appendix"
    },
    "0403/astro-ph0403291_arXiv.txt": {
        "abstract": "We study models in which the universe exits reheating at temperatures in the MeV regime. By combining light element abundance measurements with cosmic microwave background and large scale structure data we find a fairly robust lower limit on the reheating temperature of $T_{\\rm RH} \\gtrsim 4$ MeV at 95\\% C.L. However, if the heavy particle whose decay reheats the universe has a direct decay mode to neutrinos, there are some small islands left in parameter space where a reheating temperature as low as 1 MeV is allowed. The derived lower bound on the reheating temperature also leads to very stringent bounds on models with $n$ large extra dimensions. For $n=2$ the bound on the compactification scale is $M \\gtrsim 2000$ TeV, and for $n=3$ it is 100 TeV. These are currently the strongest available bounds on such models. ",
        "introduction": "The standard big bang model has been tested thoroughly up to temperatures around 1 MeV where big bang nucleosynthesis occurred. At much higher temperatures the universe is assumed to have undergone inflation, during which the primordial density perturbations are produced.  Towards the end of inflation the inflaton potential steepens so that slow roll is violated, and the universe enters the reheating phase. During this phase all particles which are kinematically allowed are produced, either by direct decay or from the thermal bath produced by the inflaton decay.  Finally the universe enters the radiation dominated phase at a temperature $T_{\\rm RH}$, which is a function of the inflaton decay rate. The only certain bound on this reheating temperature comes from big bang nucleosynthesis, and has in several previous studies been found to be around 1 MeV \\cite{kolb,kks}.  It should be noted that even if the reheating temperature after inflation is much higher there can still be subsequent \"reheating\" phases, in the sense that reheating is defined to be a period where the energy density is dominated by an unstable non-relativistic particle species. In standard reheating this is the inflaton, but in supersymmetric models it could for instance be the gravitino.  In the present paper we update previous calculations of this reheating phenomenon, using data from cosmic microwave background and large scale structure observations. Furthermore we extend the analysis to include the possibility of having a direct decay mode of the heavy particle into light neutrinos. If the heavy particle is a scalar this decay is normally suppressed by a factor $(m_\\nu/m_\\phi)^2$ because of the necessary helicity flip. However, the heavy particle could either be a non-scalar particle, or it could be a pseudo-scalar like the majoron which couples only to neutrinos. Even though such models are slightly contrived it is of interest to study whether the temperature bound on reheating is significantly affected by the possibility of direct decay into neutrinos.  In section II we discuss the set of Boltzmann equations necessary to follow the evolution of all particle species. In section III present results of the numerical solution of these equation, and in section IV we compare model predictions with observational data. Finally, section V is a review of other astrophysical constraints on heavy, decaying particles, and section VI contains a discussion. ",
        "conclusions": "We have carefully calculated constraints on models with extremely low reheating temperature, where a massive particle decays around $T \\sim 1$ MeV. By combining constraints on light element abundances with constraints on $\\eta$ and $N_\\nu$ from CMB and large scale structure we derived a fairly robust limit of \\begin{equation} T_{\\rm RH} \\gtrsim 4 \\,\\,\\, {\\rm MeV}. \\end{equation} This bound is a significant improvement over the previous bound of $T_{\\rm RH} \\gtrsim 0.7$ MeV, calculated from BBN alone. It is interesting that the lower bound is significantly higher than the $n \\leftrightarrow p$ conversion freeze-out temperature, $T \\sim 0.8$ MeV, end even higher than the neutrino decoupling temperature $T_{\\rm D} \\sim 2$ MeV. This shows that even small residual effects can be measured with present observational data.   Models with reheating temperature in the MeV regime are in general difficult to reconcile with such features as baryogenesis. However, in models with large extra dimensions a low reheating temperature is essential in order to avoid overproduction of massive Kaluza-Klein graviton states. This means that we can use our present bound to derive limits on the compactification scale in such models. For the case of two extra dimensions the bound is $M \\gtrsim 2000$ TeV and for $n=3$ it is $M \\gtrsim 100$ TeV. This bound is somewhat stronger than the bound coming from considerations of neutron star cooling and gamma ray emission."
    },
    "0403/astro-ph0403258_arXiv.txt": {
        "abstract": "{ A recent variant of the inflationary paradigm is that the ``primordial'' curvature perturbations come from quantum fluctuations of a scalar field, subdominant and effectively massless during inflation, called the ``curvaton'', instead of the fluctuations of the inflaton field. We consider the situation  where the primordial curvature perturbations  generated by the quantum fluctuations of  an inflaton and of a curvaton field are of the same order of magnitude. We compute the total curvature perturbation and its spectrum in this case and we discuss the observational consequences. } ",
        "introduction": "The first observations of the WMAP satellite \\cite{WMAP1} have confirmed the basic predictions of the inflation scenarios \\cite{guth} (see also \\cite{liddle_lyth}), namely the existence of super-Hubble primordial fluctuations, mostly adiabatic and characterized by a quasi-scale-invariant spectrum. Because direct information on the primordial universe is scarce, it is essential, when trying to interpret the cosmological data, to determine to which extent a precise measurement of the primordial spectrum gives us information on the inflationary scenario itself.  In this perspective, whereas in the standard inflationary paradigm, the primordial perturbations are generated by quantum fluctuations of the {\\it inflaton}, it has been realized recently that the observed primordial fluctuations could be accounted for by quantum fluctuations of a scalar field other than the inflaton, dubbed the {\\it curvaton} field \\cite{lw02,es02,mt01} (and see also \\cite{luw} and \\cite{mt02}). The curvaton is a scalar field $\\sigma$ whose effective mass must be much smaller than the Hubble parameter during inflation so that it acquires quantum fluctuations with a quasi scale-invariant spectrum. After inflation, the primordial curvature  perturbations are generated by the conversion of the isocurvature perturbations due to the fluctuations of the subdominant curvaton field, into adiabatic perturbations.  In the curvaton scenario, inflation is still necessary since this is the source of the curvaton quantum fluctuations but the separation of the r\\^oles played by the inflaton that drives the evolution of the universe and by the curvaton which generates the perturbations, liberates the inflaton from some observational constraints. This alternative possibility is mixed news for the inflationary scenarios. It is good news in the sense that it allows to construct more easily explicit inflationary models with a firmer footing from a particle physicist's point of view \\cite{Lythmodels}. But it is also bad news in the sense that it shows a new limitation to extract information on the nature of the inflaton field from the measurement of the primordial fluctuations.   In the curvaton scenario, the curvature  fluctuations generated by the inflaton are usually assumed to be negligible, and the primordial fluctuations are thus essentially due to the curvaton. However, as was pointed out in \\cite{dlnr03}, one can envisage  cases where   the fluctuations generated by  both the inflaton   and a curvaton-like field are relevant. This is the situation that we explore in the present work.  It can occur, for instance,   when  $\\sigma_*$ is  of the order of the Planck mass $\\mP$ during inflation, still assuming that the curvaton energy density is much smaller than that of the inflaton (otherwise we would be in the context of double inflation \\cite{ps92}), and  implies that the curvaton starts  to oscillate when its  energy density  already contributes significantly to the total energy density. This is in contrast with the usual curvaton scenario where the curvaton starts oscillating long before it dominates and  can thus be treated as a pressureless fluid during the first radiation dominated phase.  Curvaton and inflaton generated perturbations can also be of the same order of magnitude for  small $\\sigma_*$, i.e.\\ $\\sigma_*\\ll \\mP$, but with a slow-roll inflation parameter $\\epsilon$ sufficiently small during inflation to produce big inflaton perturbations.      This paper is organized as follows. In the next section we write the background evolution equations and we derive their initial conditions explaining the context of our present analysis. In Sec.~III we discuss the perturbation equations and their initial conditions. We can analytically derive the curvature perturbation after curvaton decay, i.e.\\ at the onset of the standard big bang cosmology, in two limits. In the pure curvaton limit, for $\\sigma_* \\ll \\mP$, we recover the curvature perturbation already derived in previous works. In the limit where the curvaton acts as a secondary inflaton, for $\\sigma_* \\gsim \\mP$, we recover the perturbation of the double inflation model. In the intermediate case we resort to a  numerical analysis to derive the mixed inflaton and curvaton perturbation, which is the main result of our work. In Sec.~IV we derive the spectrum of the curvature perturbation, its spectral index, and discuss the observational consequences of the model. Finally we conclude in the last section. ",
        "conclusions": "In the present work, we have explored one particular facet of the curvaton mechanism: we have considered the case where primordial curvature perturbations generated by the curvaton are of the same order of magnitude as the traditional inflaton-generated curvature perturbations. This can happen when the vacuum expectation value for the curvaton is of the order of the Planck mass during inflation. This situation can  occur in some scenarios, such as in the string axion scenario discussed in \\cite{dlnr03}.  Independently of the specific model of inflation, we have obtained the expression for the ``primordial'' curvature perturbation, given as a linear combination of the post-reheating curvature perturbation (i.e.\\ generated during inflation) and of the curvaton perturbation. In terms of the curvature perturbation on the uniform energy density hypersurface this is given as \\beq \\zeta= \\left[ { 1 \\over \\sqrt{ \\epsilon}} e_\\phi + {\\tilde f}(\\sigma)   e_\\sigma \\right] \\frac{H}{2 \\sqrt{2} \\pi \\mP}, \\quad \\langle e_\\phi e_\\sigma \\rangle = 0, \\eeq where $H$, $\\sigma$, and $\\epsilon$ are the Hubble parameter, the vacuum expectation value of the curvaton, and the first slow-roll parameter, all evaluated at Hubble crossing during inflation, respectively [see Eq.~(\\ref{curvature})]. This expression is only valid for large $\\f$, otherwise corrections which are second order in the slow-roll parameters must be introduced. The function $\\tilde f(\\sigma)$ depends on the initial background value for the curvaton and interpolates between two regimes, which have been investigated previously in the literature: the standard curvaton regime, where the curvaton starts oscillating long before it dominates, and for which $\\tilde f = (2\\sqrt{2}/3)(\\mP/\\sigma)$, and the secondary  inflaton regime, where the ``curvaton'' is still frozen when it dominates and thus produces an additional period of inflation, and $\\tilde f = (1/\\sqrt{2})(\\sigma/\\mP)$.     In the case of slow-roll inflation, we have computed the spectral amplitude and index associated with  curvature perturbations of mixed origins and shown how the resulting expressions nicely interpolate between the pure inflaton and pure curvaton cases. As an illustration, we have considered the case of a quartic potential for the inflaton and shown how allowing for a mixing with a curvaton changes the usual predictions.  The results derived here  are quite generic and can be applied to any inflationary model. Although  we have considered, for simplicity,  only a quadratic potential for the curvaton, which is a good approximation near a local minimum of most potentials, it is in principle straightforward to extend our results to other potentials. As we have shown, a generic consequence of mixed perturbations is a modification of the usual consistency relation, which gives some hope to be able to distinguish  observationally between the standard inflationary scenario and its mixed version.    \\vskip 1cm {\\bf Acknowledgements} We wish to thank David Lyth, J\\'er\\^ome Martin and David Wands for stimulating discussions and instructive comments.  F.V. acknowledges financial support from the Swiss National Science Foundation."
    },
    "0403/astro-ph0403634_arXiv.txt": {
        "abstract": "We present spectroscopy of the microquasar SS~433 obtained near primary eclipse and disk precessional phase $\\Psi = 0.0$, when the accretion disk is expected to be most ``face-on''.  The likelihood of observing the spectrum of the mass donor is maximized at this combination of orbital and precessional phases since the donor is in the foreground and above the extended disk believed to be present in the system.  The spectra were obtained over four different runs centered on these special phases. The blue spectra show clear evidence of absorption features consistent with a classification of A3-7~I. The behavior of the observed lines indicates an origin in the mass donor.  The observed radial velocity variations are in anti-phase to the disk, the absorption lines strengthen at mid-eclipse when the donor star is expected to contribute its maximum percentage of the total flux, and the line widths are consistent with lines created in an A supergiant photosphere. We discuss and cast doubt on the possibility that these lines represent a shell spectrum rather than the mass donor itself. We re-evaluate the mass ratio of the system and derive masses of $10.9\\pm 3.1~M_\\odot$ and $2.9\\pm0.7~M_\\odot$ for the mass donor and compact object plus disk, respectively.  We suggest that the compact object is a low mass black hole.  In addition, we review the behavior of the observed emission lines from both the disk/wind and high velocity jets. ",
        "introduction": "The well-studied binary SS~433, despite the wealth of observational and theoretical studies, is still one of the most enigmatic systems. SS~433 is an X-ray binary consisting of a compact star and a mass donor star, called such because it donates mass to a precessing extended disk surrounding the compact companion, typically considered to be a black hole \\citep[e.g.][]{lei84} or neutron star \\citep[e.g.][]{dod91}.  SS~433 is also well known as a source of collimated relativistic jets with speeds $v_{jet}\\approx 0.26c$ \\citep[e.g.][]{mar82}.  The number of components in this system combined with their complexity greatly obscures interpretation of observational data and theoretical modeling, and thus many physical parameters of the system are as yet undetermined.  The 13 day orbital period of the binary as well as the 162 day period of the disk precession are well established \\citep{gor98,eik01,gie02}. The radial velocity behavior of the compact companion and inner disk, while complicated by emission from other components, has also been convincingly established with a semiamplitude of $\\approx 170$ km s$^{-1}$\\citep{fab90,gie02}. \\citet*{ant87} suggested that observation of the mass donor spectrum could be possible and would lead to direct kinematical masses for the two components.  An accurate detection of the mass donor spectrum is crucial in determining if the compact companion is a black hole or neutron star.  A first attempt at this was made by \\citet[][= GHM03]{gie03} who found similarities between absorption features in the mid-eclipse spectrum of SS~433 and an A-type evolved star, specifically the A7 Ib star HD~148743.  Their detection provided a tentative determination of the radial velocity curve of the mass donor resulting in a mass determination for both components.  We set out to confirm this tentative detection of the mass donor with additional observations.  Absorption features from the mass donor star may best be observed during primary eclipse (donor star inferior conjunction) when much of the continuum light from the disk is blocked and the donor star is in the foreground.  The precession of the extended disk complicates detection of the mass donor since we expect the disk to have both a high opacity and large vertical extent. This means that light from the mass donor will be obscured by the disk except near precessional phase zero, the phase of maximum disk opening. The combination of precessional phase $\\Psi = 0.0$ and orbital phase $\\phi = 0.0$ limits the possible observing times to roughly two five-night windows per year.  Here we present new spectroscopy of SS~433 at three epochs corresponding to $\\Psi = 0.0$ and $\\phi = 0.0$ that we combine with the original data from GHM03.  We show that the absorption spectrum is well matched by the spectrum of HD~9233, an A4~Iab star.  We also utilize the emission features in the spectrum to confirm both the orbital and precessional phases of observation.  The observed jet features and the ``stationary'' emission lines are also discussed. ",
        "conclusions": "We have shown that the mass donor in SS~433 is likely to be an A3-7~I star.  A similar kind of absorption spectrum may be produced by circumstellar gas, but we argue that the strength of the spectrum and its orbital variation in intensity and Doppler shift all indicate a photospheric origin. Further observations with higher S/N ratio spectra would allow us to classify more accurately the spectrum and calculate temperature, abundances, and surface gravity for the mass donor. With additional radial velocity measurements to help constrain the fit, more accurate masses would also be obtained which would lead to a more secure identification of the compact star as either a black hole or neutron star.  Our estimates of the mass ratio and masses fall within the values predicted in the literature and begin to give us a more definite understanding of the physical parameters of this peculiar system.  The evolutionary scenario which best fits our determined physical parameters is that of \\citet{kin00}, in which the mass donor star is in a phase of rapid evolution as it crosses the Hertzsprung gap and the compact companion is a low mass black hole.  We have also shown that the orbital ephemeris of \\citet{gor98} and disk precession/nodding model parameters of \\cite{gie02} are still valid."
    },
    "0403/astro-ph0403345_arXiv.txt": {
        "abstract": "We use a numerical simulation of cosmological reionization to estimate the likelihood of detecting Lyman-$\\alpha$ emitting galaxies during the pre-reionization era. We show that it is possible to find galaxies even at $z\\sim9$ that are barely affected by the dumping wing of the Gunn-Peterson absorption from the neutral IGM outside of their $\\HII$ regions. The damping wing becomes rapidly more significant at $z>9$, but even at $z>10$ is it not inconceivable (although quite hard) to see a Lyman-$\\alpha$ emission line from a star-forming galaxy. ",
        "introduction": "A recent landslide in the number of detections of high redshift Lyman-$\\alpha$ emitters (Dey et al.\\ 1998; Hu, Cowie, \\& McMahon 1998; Weymann et al.\\ 1998; Hu, McMahon, \\& Cowie 1999; Stern \\& Spinrad, 1999; Ellis et al.\\ 2001; Ajiki et al.\\ 2002; Hu et al.\\ 2002; Kodaira et al.\\ 2003; Taniguchi et al.\\ 2003; Rhoads et al.\\ 2003; Maier et al.\\ 2003; Cuby et al.\\ 2003; Bunker et al. 2003; Hu et al.\\ 2004; Pello et al.\\ 2004; Kneib et al.\\ 2004) offers a unique probe of the ionization history of the universe during cosmological reionization. As observations of Lyman-$\\alpha$ absorption in the spectra of high redshift quasars (Djorgovski et al.\\ 2001; Becker et al.\\ 2001; Fan et a.\\ 2002, 2003; White et al.\\ 2003; Songaila 2004) unambiguously indicate that the universe was reionized shortly before $z=6$ (Gnedin 2002), the observations of Lyman-$\\alpha$ at $z>6$ directly probe the pre-reionization era.  Specifically, since the damping wing of the Gunn-Peterson absorption (Gunn \\& Peterson 1965) of the neutral IGM will affect the Lyman-$\\alpha$ emission line if the $\\HII$ region around the emitting galaxy is not large enough (Miralda-Escud\\'e 1998; Haiman 2002), the observations of high redshift Lyman-$\\alpha$ emitters tell us something about the distribution and sizes of cosmological $\\HII$ regions - although, to understand what specific constraints the observations place is a horrendously difficult task\\footnote{Since most Lyman-$\\alpha$ emitters are discovered because they are highly lensed by foreground clusters of galaxies, determining the observational selection function is extremely difficult.}.  However, the very fact of discovering a high-redshift Lyman-$\\alpha$ emitter may potentially be used to constraint the ionization history of the universe. For example, Pello et al.\\ (2004) report a tentative discovery of a $z=10$ galaxy. Is it possible to see a Lyman-$\\alpha$ emitting galaxy at such a high redshift? Loeb et al.\\ (2004) used this fact to put some preliminary constraints on the ionization state of the IGM around that galaxy. While a real situation is likely to be more complicated because of the variety of reasons ($\\HII$ regions are typically highly asymmetric because I-fronts propagate much faster across underdense voids then across dense filaments; bright galaxies are biased, so it is likely that more than one galaxy is located inside a single $\\HII$ region; Lyman-$\\alpha$ emitters can also be located inside $\\HII$ regions of luminous quasars, which are often many times larger than $\\HII$ regions of galaxies, etc), it is interesting to investigate how likely it is to find a high redshift galaxy unaffected by the damping wing of the Gunn-Peterson absorption from the surrounding IGM. ",
        "conclusions": "Somewhat unexpectedly, even at as high a redshift as $z\\sim9$, one percent of all Lyman-$\\alpha$ emitters should have their emission lines barely affected by the damping wing. At higher redshifts the damping wing quickly becomes more significant, but finding Lyman-$\\alpha$ emitters at $z>10$ should still be possible, although they must be quite rare.  The reason for this rapid transition is quite simple: as the universe expands, $\\HII$ regions increase in size, so at each moment there exists a characteristic scale for the size distribution of $\\HII$ regions (say, a typical size of an $\\HII$ region around an $L_*$ galaxy). At the same time, the damping wing of the Gunn-Peterson absorption introduces another spatial scale (about $\\Delta z=0.01$ or proper distance of $200\\dim{kpc}$ at $z\\sim10$; Miralda-Escid\\'e 1998). Thus, when typical $\\HII$ regions around bright galaxies exceed this scale, a substantial fraction of all Lyman-$\\alpha$ emitters becomes observable.  It is important to emphasize here that our estimate is likely only a lower limit: because of the finite size of the computational box, we miss the brightest galaxies which create the largest $\\HII$ regions (although the size of the simulation box was selected so as to make this effect not significant at $z>6$). In addition, our simulation does not include quasars, which can create $\\HII$ regions several times that of galaxies. High redshift quasars are rare (there should be only about 2 SLOAN quasars in the whole observable universe between redshifts 9 and 10, Fan et al.\\ 2001), so that effect is most likely not large too, but without a detailed simulation we cannot estimate it, so we treat our results as a lower limit."
    },
    "0403/astro-ph0403173_arXiv.txt": {
        "abstract": "The hydromagnetic structure of a neutron star accreting symmetrically at both magnetic poles is calculated as a function of accreted mass, $M_{\\rm a}$, and polar cap radius, starting from a centered magnetic dipole and evolving through a quasistatic sequence of two-dimensional, Grad-Shafranov equilibria. The calculation is the first to track fully the growth of high-order magnetic multipoles, due to equatorward hydromagnetic spreading, while simultaneously preserving flux freezing and a self-consistent mass-flux distribution. Equilibria are constructed numerically by an iterative scheme and analytically by Green functions. Two key results are obtained, with implications for recycled pulsars. (i) The mass required to significantly reduce the magnetic dipole moment, $10^{-5}\\Msun$, greatly exceeds previous estimates ($\\sim 10^{-10}\\Msun$), which ignored the confining stress exerted by the compressed equatorial magnetic field. (ii) Magnetic bubbles, disconnected from the stellar surface, form in the later stages of accretion ($M_{\\rm a} \\gtrsim 10^{-4}\\Msun$). ",
        "introduction": "} Observations of low-field neutron stars in binary systems containing white-dwarf and supergiant companions, with a history of disc-fed and wind-fed accretion respectively, suggest that the magnetic dipole moment $|\\vv{m}|$ of a neutron star decreases monotonically with accreted mass, $M_{\\rm a}$ (\\citealt{taa86,van95}; see \\citealt{wij97} for a dissenting view). Several mechanisms have been proposed to explain why $|\\vv{m}|$ is reduced: (i) accelerated Ohmic decay, where the electrical conductivity of the crust is lowered by accretion-induced heating \\citep{urp95,urp97}; (ii) interactions between superfluid neutron vortices and superconducting magnetic fluxoids in the stellar interior \\citep{mus85,sri90}; and (iii) magnetic screening or burial, where the currents generating the natal magnetic field are partially neutralized by accretion-induced currents in the crust \\citep{blo86,aro80}. For a critical review of these mechanisms, see \\citet{mel01}.  In this paper, we study the mechanism of magnetic burial in detail. In the early stages of accretion ($M_{\\rm a} \\lesssim 10^{-10}\\Msun$), accreted matter accumulates in a column at the polar cap, minimally distorting the magnetic field. The mass-flux distribution in this regime has been calculated by Grad-Shafranov methods, with the prediction that $|\\vv{m}|$ is reduced by $\\sim 1$ per cent for $M_{\\rm a} \\approx 10^{-10}\\Msun$ \\citep{ham83,bro98,lit01}. We show that these calculations overestimate the amount of screening; in fact, $M_{\\rm a} \\gtrsim 10^{-5}$ is required to reduce $|\\vv{m}|$ by 10 per cent when the confining stress of the compressed, equatorial magnetic field is modelled faithfully. In this regime, inaccessible to previous analyses due to numerical breakdown \\citep{ham83,bro98,lit01}, the latitudinal pressure gradient at the base of the polar column forces the polar magnetic field to buckle, and the accreted material spreads equatorward together with frozen-in magnetic flux \\citep{mel01}. We compute the structure of the highly distorted magnetic field, and hence $|\\vv{m}|$, as a function of $M_{\\rm a}$.  A key advance in the present work is that the mass-flux distribution in each equilibrium state is self-consistent; our equilibria are generated by a continuous deformation of the flux surfaces of the initial magnetic field (say, a dipole), in a manner which preserves flux-freezing. This is not true of previous calculations, where the mass-flux distribution is unconstrained relative to the initial state \\citep{bro98,ham83,lit01,mel01}. However, several other important effects are not included to keep the problem manageable. (i) Ohmic dissipation is neglected, even though the diffusion and accretion time-scales are comparable for the smallest magnetic structures predicted by the theory \\citep{cum01,bro98}. (ii) We do not investigate the stability of the hydromagnetic equilibria we compute; sharp magnetic-field gradients are potentially disrupted by Rayleigh-Taylor and interchange instabilities \\citep{cum01,mel01,bha99}. (iii) We treat the neutron star as a hard surface; subsidence of accreted material, and incorporation into the crust, are neglected \\citep{bha99}.  The paper is structured as follows. In Section 2, we introduce the theoretical framework for calculating the self-consistent hydromagnetic equilibrium state of an accreting neutron star. Analytic and numerical methods of solution are given in Section 3. The properties of the equilibria are investigated in Section 4, $|\\vv{m}|$ is computed as a function of $M_{\\rm a}$ and the radius of the polar cap, and the novel feature of magnetic bubbles is explored. The limitations of our results, with respect to time-dependent processes like hydromagnetic instabilities and ohmic dissipation, are assessed in Section 5. ",
        "conclusions": "Observations of low-field binary neutron stars and recycled pulsars imply that the magnetic dipole moment of a neutron star is reduced by accretion. In this paper, we undertake a self-consistent analysis (numerical and analytic) of one mechanism that may account for the reduction observed: polar magnetic burial in the ideal-magnetohydrodynamic regime. Our analysis has several new features. (i) Flux freezing is strictly enforced when connecting the final and initial magnetic configurations by solving self-consistently for the mass-flux distribution rather than specifying it ad hoc (see Figure \\ref{fig:compareF}). (ii) The Lorentz force due to equatorial magnetic field lines compressed by equatorward hydromagnetic spreading is included when calculating the confinement of the polar accreted column. (iii) Numerical methods are developed for treating accreted masses up to $\\approx 10^{-4}\\Msun$, (cf. $10^{-10}\\Msun$ in previous work), where the field is dramatically distorted and high-order multipoles dominate.  We report two key results. (i) $M_{\\rm a} \\gtrsim 10^{-5}\\Msun$ must be accreted in order to reduce significantly the magnetic dipole moment $|\\vv{m}|$ of the star, contrary to previous estimates ($M_{\\rm a} \\approx 10^{-10}\\Msun$) which neglected equatorial magnetic stresses. For small $M_{\\rm a}$, we find $|\\vv{m}| = |\\vv{m}_{\\rm i}|(1 - M_{\\rm a}/M_{\\rm c})$, with $M_{\\rm c} = 1.2\\times 10^{-6}\\Msun$, (cf.\\ \\citealt{shi89}). For $M_{\\rm c} \\lesssim M_{\\rm a} \\lesssim 10^{-4}\\Msun$, we find $|\\vv{m}| = |\\vv{m}_{\\rm i}|(M_{\\rm a}/4.6\\times 10^{-5}\\Msun)^{2.25\\pm 0.22}$. (ii) When enough mass is accreted, such that $M_{\\rm a} \\geq 10^{-4}b^{-1}\\Msun$, the hydrostatic pressure gradient generates flux surfaces with $\\psi < 0$ or $\\psi > \\psi_{*}$, creating closed magnetic bubbles that are disconnected topologically from the stellar surface. The bubbles are valid solutions of the Grad-Shafranov boundary problem, as confirmed by analytic, Green function calculations; they are not numerical artefacts or fingerprints of hydromagnetic instabilities (e.g.\\ Parker).  Several of our assumptions need to be relaxed in future work, including (i) perfect conductivity, (ii) an impenetrable stellar surface, and (iii) axisymmetry (which tends to suppress hydromagnetic instabilities, for example). Finally, the uniqueness of the hydromagnetic equilibria we compute numerically is yet to be established.  Our equilibria serve as useful starting points for exploring the stability of the magnetic configuration during and after accretion. Our theoretical results will be tested against observational data from binary neutron stars and recycled pulsars in a companion paper."
    },
    "0403/astro-ph0403490_arXiv.txt": {
        "abstract": "We present R-band images covering more than 11 square degrees of sky that were obtained in preparation for the Spitzer Space Telescope First Look Survey (FLS). The FLS was designed to characterize the mid-infrared sky at depths 2 orders of magnitude deeper than previous surveys.  The extragalactic component is the first cosmological survey done with \\sst.  Source catalogs extracted from the R-band images are also presented.  The R-band images were obtained using the MOSAIC-1 camera on the Mayall 4m telescope of the Kitt Peak National Observatory.  Two relatively large regions of the sky were observed to modest depth: the main FLS extra galactic field ($17^h18^m00^s$ $+59^o30'00''.0$ J2000; l$=88.3$, b$=+34.9$) and ELAIS-N1 field ($16^h10^m01^s$ $+54^o30'36''.0$ J2000; l$=84.2$, b$=+44.9$).  While both of these fields were in early plans for the FLS, only a single deep pointing test observation was made at the ELAIS-N1 location. The larger Legacy program SWIRE (Lonsdale et al., 2003) will include this region among its surveyed areas.  The data products of our KPNO imaging (images and object catalogs) are made available to the community through the World Wide Web (via the Spitzer Science Center and NOAO Science Archives).  The overall quality of the images is high. The measured positions of sources detected in the images have RMS uncertainties in their absolute positions of order 0.35 arc-seconds with possible systematic offsets of order 0.1 arc-seconds, depending on the reference frame of comparison.  The relative astrometric accuracy is much better than 0.1 of an arc-second. Typical delivered image quality in the images is 1.1 arc-seconds full width at half maximum.  Images are relatively deep since they reach a median 5$\\sigma$ depth limiting magnitude of R=25.5 (Vega), as measured within a 1.35~$\\times$~FWHM aperture for which the S/N ratio is maximal.  Catalogs have been extracted using SExtractor using thresholds in area and flux for which the number of false detections is below 1\\% at R=25. Only sources with S/N greater than 3 have been retained in the final catalogs.  Comparing the galaxy number counts from our images with those of deeper R-band surveys, we estimate that our observations are 50\\% complete at R=24.5. These limits in depth are sufficient to identify a substantial fraction of the infrared sources which will be detected by {\\it Spitzer}. ",
        "introduction": "One of the main advantages of the Spitzer Space Telescope (formerly known as the Space Infrared Infrared Telescope Facility; SIRTF, Gallagher et al. 2003) is the possibility to make extragalactic surveys of large regions of the sky in a relatively short time covering wavelengths from the near-IR to the far-IR with the instruments IRAC (Fazio et al. 1998) and MIPS (Rieke et al. 1996). Compared to {\\it Spitzer}'s predecessors (e.g. IRAS, Soifer et~al. 1983, and ISO, Kessler et al. 1996), there are improvements in the detectors (number of pixels and better responsivity), the collecting area of the primary mirror (85cm diameter), and Sun-Earth-Moon avoidance constraints due {\\it Spitzer}'s heliocentric orbit. \\sst\\ can also make observations simultaneously in multiple bands (with IRAC--3.6 and 4.5 or 5.8 and 8 $\\mu$m; with MIPS--24, 70 and 160 $\\mu$m).  Many extragalactic surveys are already scheduled with {\\it Spitzer} as Legacy programs (SWIRE, Lonsdale et al. 2003; GOODS, Dickinson \\& Giavalisco 2001) or as observations by the Instrument Teams (Wide, Deep and Ultra-Deep {\\it Spitzer} surveys which will cover regions like the Bo\\\"otes field of the NOAO Deep Wide-Field Survey, the Groth strip, Lockman Hole, XMM-Deep and so on, see e.g.~Dole et al. 2001). The First Look Survey utilizes 112 hours of Director's Discretionary time on \\sst\\ and includes extragalactic, galactic, and ecliptic components\\footnote{see the FLS website at ssc.spitzer.caltech.edu/fls}. These data will be available to all observers when the Spitzer Science Archive opens in May, 2004. The purpose of the FLS is to characterize the mid-infrared sky at previously unexplored depths and make these data rapidly available to the astronomical community. The extragalactic component is comprised of a 4 square degree survey with IRAC and MIPS near the north ecliptic pole centered at J1718+5930. These observations were executed on 2003 December 1-11.  To fully exploit the \\sst\\ FLS data we have obtained ancillary surveys at optical (this paper) and radio wavelengths (Condon et~al. 2003). Given the modest spatial resolution of the {\\it Spitzer} imagers (the point spread function is large, especially in the mid and far-IR, e.g. 5.7 arcseconds full width at half maximum, FWHM, for the 24 $\\mu$m channel), the first problem to solve for the infrared sources detected by {\\it Spitzer} will be to associate these sources with an optical counterpart, when possible. This will then allow the higher spatial resolution of the available optical images to assist with the source classification (e.g. as stars, galaxies or QSO) and enable targeting of subsets of the sources for spectroscopy with optical or near-IR spectrographs. Since many of the infrared sources that will be detected by {\\it Spitzer} will be dust-obscured galaxies with faint optical counterparts, the complementary optical imaging must be relatively deep.  Although a deep multi-wavelength optical survey would be more useful, allowing one to compute photometric redshifts (e.g. the NOAO Deep Wide-Field Survey, Jannuzi and Dey 1999, Brown et al. 2003), the task of deeply covering a large region of sky in a homogeneous manner is quite time-consuming. Therefore, for the initial optical ancillary survey we chose to observe the entire field in the R-band.  NOAO provided 4 nights of Director's Discretionary time on the KPNO 4m in May, 2000, for this survey. We have limited multi-wavelength optical observations to the central portion of the FLS field. The Sloan Digital Sky Survey included the FLS field in their early release observations (Stoughton et al. 2002) and mosaics and catalogs for the region are also now available (Hogg et~al. 2004).  In this paper, we present the R-band optical observations made with the MOSAIC-1 camera on the Mayall 4m Telescope of Kitt Peak National Observatory. Centered on the main FLS field, a region 9.4 square degrees in area was imaged. In addition, 2.3 square degrees covering the ELAIS-N1 field was also observed.  Although originally the ELAIS-N1 field was planned to be part of the FLS program, the FLS observations of the ELAIS-N1 field have now been revised to a very deep 10'$\\times$10' pointing to evaluate the confusion limits of the MIPS instrument.  The remainder of the ELAIS-N1 field will now be imaged as part of a larger survey in this region, a portion of the SWIRE {\\it Spitzer} Legacy Survey (Lonsdale et al., 2003).  In section 2 we review the overall observing strategy and describe the MOSAIC-1 observations. In section 3 we discuss the techniques used in the data reduction including the astrometric and photometric calibration of the images. We describe in section 4 the data products made publicly available.  We detail in section 5 the criteria used to detect, classify, and photometrically measure objects in the images. Section 5 also includes a description of the information available in our catalogs.  In section 6 we examine the quality of the imaging data by comparing them with other available data sets. Finally, a brief summary is given in section 7. ",
        "conclusions": "A deep NOAO/KPNO survey in the R band has been carried out to observe a field of more than 9 square degrees centered at 17:18:00 +59:30:00 (J2000) aimed to find optical counterparts for the {\\sl First Look Survey} which surveys 7 different infrared wavelengths with the instruments IRAC and MIPS using the Spitzer Space Telescope. Another 2.3 square degrees have been surveyed in the ELAIS-N1 region which will be observed in the \\sst\\ SWIRE Legacy survey.  This paper describes the observation strategy, the data reduction and the products which are publicly available to the astronomical community on the World Wide Web at the Spitzer Science Center and the NOAO Science Archives.  The overall quality of the data is good and homogeneous: the average seeing is 1.1 and typically varies between 0.9 and 1.2. The limiting magnitude of the images, measured inside an aperture of 1.35 $\\times$ FWHM for which the S/N ratio is maximal, is around R=25.5 at 5$\\sigma$, deep enough to detect optical counterparts for a substantial fraction of the new {\\it Spitzer} selected objects.  An average number of  35000 extragalactic sources are detected in each subfield (40' $\\times$ 40', approximatively) with a 50\\% completeness limit of R=24.5  as deduced by comparing the counts with other deeper surveys.  Images and catalogs are available to the astronomical community along with the first release of the FLS infrared data to exploit in the best way the wealth of extragalactic data expected from the new infrared observatory {\\it Spitzer}."
    },
    "0403/astro-ph0403459_arXiv.txt": {
        "abstract": "We develop a new method of modeling microlensing events based on a Monte Carlo simulation that incorporates both a Galactic model and the constraints imposed by the observed characteristics of the event. The method provides an unbiased way to analyze the event especially when parameters are poorly constrained by the observed lightcurve. We apply this method to search for planetary companions of the lens in OGLE-2003-BLG-423, whose maximum magnification $A_{\\rm max}=256\\pm 43$ (or $A_{\\rm max}=400\\pm 115$ from the lightcurve data alone) is the highest among single-lens events ever recorded.  The method permits us, for the first time, to place constraints directly in the planet-mass/projected-physical-separation plane rather than in the mass-ratio/Einstein-radius plane as was done previously.  For example, Jupiter-mass companions of main-sequence stars at 2.5 AU are excluded with 80\\% efficiency. ",
        "introduction": "\\label{sec:intro} High-magnification microlensing events are exceptionally sensitive to the presence of planetary companions to the lens. As the projected separation of the source and the (parent-star) lens decreases, the size of the images increases, thus enhancing the probability that the planet will pass close enough to an image to generate a noticeable deviation in the lightcurve \\citep{gl}.  And if the source gets sufficiently close to the lens, the lightcurve can be perturbed by the central caustic associated with the parent star itself \\citep{gs}.  Groups that monitor microlensing events to search for planets are well aware of this enhanced sensitivity and so devote special attention to these events.  It is therefore somewhat surprising that the vast majority of events monitored by these groups in the past have not been particularly high-magnification, and the vast majority of the observations of the few that did reach high magnification were actually performed well away from the peak, when the event had far less sensitivity to planets.  In particular, of the 43 events monitored by the PLANET (Probing Lensing Anomalies NETwork) collaboration \\citep{five,gaudi02} over 5 years, most of the sensitivity to planets came from just 5 or 6 events, and most of that from the near-peak regions of these events. \\citet{tsa} and \\citet{sno} used the relatively sparse OGLE data to put limits on planetary systems, although \\citet{gh04} have argued that planets could not be reliably detected from such data alone.  The main reason for this apparent discrepancy was simply a shortage of microlensing alerts.  Hence, at any given time, there just were no high-magnification events in progress, or at least none near their peak.  The available telescope time then had to be applied to less favorable events.  In addition, when devising their observational strategy, PLANET considered that they would have to characterize the events they were monitoring entirely with their own data. Such characterization is absolutely essential to evaluating the sensitivity of each event to planets, and it requires a very large number of observations on the wings and at baseline when the event has very little sensitivity to planets.  With the commencement of the Optical Gravitational Lens Experiment's OGLE-III project \\citep{ogleIII}, the situation is radically changed.  Using its dedicated 1.3 m telescope, large-format ($35'\\times 35'$) camera, generally excellent seeing, and ambitious observing strategy, OGLE-III is alerting microlensing events toward the Galactic bulge at a rate of 500 per season.  Since microlensing events are uniformly distributed in impact parameter $u_0$, and since peak magnification scales $A_\\max\\sim u_0^{-1}$ (for $u_0\\ll 1$), this implies that there are dozens of events with $A_\\max\\ga 10$ each year, and a handful with $A_\\max\\ga 100$.  Moreover,  OGLE-III photometry is publicly available (literally hours after it is taken), so there is generally no need for microlensing followup groups to monitor the wings or baseline in order to characterize the event. That is, followup observing time can be concentrated on the highly sensitive peaks of the high-magnification events.  In the 2003 season, such peaks were occurring almost continuously.  OGLE-III is therefore generating substantial new opportunities for microlensing planet search groups such as PLANET \\citep{planet}, the Microlensing Planet Search (MPS, \\citealt{mps,rhie}), and the Microlensing Follow-Up Network ($\\mu$FUN, \\citealt{ob262}).  However, the OGLE-III approach also generates substantial challenges. In order to monitor a very large area during the 2002 and 2003 seasons, OGLE-III returned to each field only of order 50 times over the roughly 9 month season.  During the long nights around 21 June, when the bulge transits near midnight, the cadence was relatively high, once every two or three nights. But at the edges of bulge season, the rate of return dropped as low as once per week or less.   Hence, some extremely high magnification events appeared quite ordinary as of their last observation before peak, and then could be recognized for what they were only very close to (or past) their peak.  Indeed, it is quite possible that their true nature as high magnification events could not be recognized at all from the OGLE-III lightcurve alone.  Thus, without additional work, the riches generated by OGLE-III could easily pass by unnoticed. Beginning in the 2004 season, OGLE adjusted its strategy to concentrate on a reduced number of fields that have relatively higher expected event rates.  Hence, it is expected that the above-mentioned problems will be mitigated in future seasons.  Here we develop a new method of modeling microlensing events that incorporates both a Galactic model via a Monte Carlo simulation and the constraints imposed by the observed characteristics of the event. We apply this method to the extreme microlensing event (EME) OGLE-2003-BLG-423, which at $A_\\max\\sim 250$, proves to have the highest magnification ever recorded among single-lens events.  As such, the event also has the greatest potential sensitivity to planetary companions of the lens, with substantial probability of detecting even Neptune mass planets, whose event timescale would typically be only about 6 hours. This enhanced sensitivity poses special challenges to the analysis because both the form and amplitude of the impact of such small planets on the lightcurve will depend on the relative size of the source compared to the Einstein ring.  If this relative size were known, it would be straightforward to calculate its effect.  However, since the lightcurve is consistent with a point source, our information on the source size is limited.  Similarly, using the lightcurve data alone the impact parameter $u_0$ is measured only to about 30\\%. If $u_0$ were known much more precisely (as it often is for events with relatively bright sources), then we would be able to specify with equal precision where in relation to the Einstein ring a planet could be and still avoid detection.  With our less perfect knowledge of $u_0$, however, we must be satisfied with a more probabilistic statement about these locations.  Both of these challenges are likely to be generic to the analysis of EMEs.  Because such events occur with low probability, their sources are most likely to be the relatively common main-sequence stars that normally lie unnoticed in ground-based bulge images, but which can briefly leap to prominence in an EME. Since these main-sequence stars are faint and hence small, they will most often avoid finite-source effects even in EMEs.  Their faintness also induces large photometric errors in the wings of the lightcurve, the region that must be well-measured to accurately determine $u_0$.  For similar reasons, these two challenges are likely to be key issues in future, even more aggressive, microlensing experiments that aim to detect Earth-mass planets either by space-based \\citep{gest} or ground-based \\citep{ghg} observations.  In our analysis, we will take as our starting point the method pioneered by \\citet{gsa} and \\citet{ob14}, which was then applied to a much larger sample by \\citet{gaudi02}. However, we improve upon this method in several respects. First, we fix the impact parameter $u_0$ at a series of different values consistent with the event data and evaluate the sensitivity to companions at each $u_0$.  To find the net sensitivity, we must weight each of these outcomes by the relative probability that the actual event had that particular $u_0$.  Second, we determine these relative probabilities not just from the fit to the lightcurve data, but by incorporating the results of a Monte Carlo simulation of events toward the actual line of sight.  For each trial $u_0$, we weight the simulated events by how well they reproduce both the observed characteristics of the lightcurve and the probability that the source has the luminosity inferred from the lightcurve combined with the Monte Carlo event parameters, as determined from the Hipparcos luminosity distribution at the observed color of the source. This method not only allows us to more accurately estimate the planetary sensitivity, it also permits us to characterize this sensitivity as a function of planet mass and planet-star separation, since each simulated event has a definite lens mass (drawn from the adopted mass function) and definite lens and source distances (and so definite Einstein radius). In contrast, the original approach of \\citet{ob14} yielded sensitivities in terms of two lightcurve-fit parameters, the planet-star mass ratio and the separation in units of the Einstein radius.  This method would also permit a similarly rigorous statistical treatment of finite source effects, since each simulated event has a definite ratio of source size to Einstein radius. However, based on the Monte Carlo, we show that in the case of OGLE-2003-BLG-423, finite-source effects are negligible. ",
        "conclusions": ""
    },
    "0403/astro-ph0403203_arXiv.txt": {
        "abstract": "We study the $\\gamma$-ray emission properties of pulsars by using a new self-consistent outer gap model. The outer gap can exist in pulsars with age over million years old if the effect of magnetic inclination angle as well as the average properties of the outer gap are considered. The mature $\\gamma$-ray pulsars, whose ages are between 0.3 to 3 million years old,  are able to move up to high galactic latitude. Moreover, their $\\gamma$-ray luminosity are weaker and their spectra are softer than those of younger $\\gamma$-ray pulsars in the galactic plane significantly. We use a Monte Carlo method to simulate the statistical properties of $\\gamma$-ray pulsars in the Galaxy as well as in the Gould Belt. We find that $\\gamma$-ray pulsars located at $\\mid b \\mid < 5^{\\circ}$ and located at $\\mid b \\mid > 5^{\\circ}$ have very different properties. High galactic latitude $\\gamma$-ray pulsars are dominated by mature pulsars with longer periods, weaker fluxes and softer spectra. If the pulsar birth rate in the Galaxy and the Gould Belt are $\\sim 10^{-2}yr^{-1}$ and $\\sim 2\\times 10^{-5}yr^{-1}$ respectively, there are 42 and 35 radio-quiet $\\gamma$-ray pulsars for $\\mid b \\mid < 5^{\\circ}$ and $\\mid b \\mid > 5^{\\circ}$ respectively. Radio-quiet $\\gamma$-ray pulsars from the Gould Belt are 2 and 13 for $\\mid b \\mid < 5^{\\circ}$ and $\\mid b \\mid > 5^{\\circ}$ respectively. We suggest that a good fraction of unidentified EGRET $\\gamma$-ray sources may be these radio-quiet $\\gamma$-ray pulsars. Furthermore $\\gamma$-ray pulsars located at $\\mid b \\mid > 5^{\\circ}$ satisfies $L_{\\gamma} \\propto L^{\\beta}_{sd}$ whereas $L_{\\gamma} \\propto L^{\\delta}_{sd}$ for $\\gamma$-ray pulsars in the galactic plane, where $\\beta \\sim 0.6$ and $\\delta \\sim 0.3$ respectively. ",
        "introduction": "There are 170 unidentified $\\gamma$-ray sources in the third EGRET catalog (Hartman et al. 1999), where $\\sim$50 sources close to the Galactic plane with $|b|<5^{\\circ}$ and $\\sim$70 sources in the medium latitudes with $|b|$ between $5^{\\circ}$ and $30^{\\circ}$. For those unidentified $\\gamma$-ray sources in the Galactic plane, many of them are associated with Wolf-Rayet and Of stars, SNRs and OB stars(Montmerle 1979; Kaaret \\& Cottam 1996; Yadigaroglu \\& Romani 1997, Romero et al. 1999). Most of these objects are considered as pulsar tracers, therefore it is natural to suggest that these low latitude sources may be Geminga-like pulsars, which are radio-quiet pulsars (Yadigaroglu \\& Romani 1995; Cheng \\& Zhang 1998; Zhang, Zhang \\& Cheng 2000). Since $\\gamma$-ray pulsars are known to be steady $\\gamma$-ray emitters, it has been suggested that the variability should be a good indicator to identify the real pulsar candidates from the unidentified EGRET sources (Romero, Combi \\& Colomb 1994; Mclaughlin et al. 1996; Zhang et al. 2000; Torres et al. 2001; Nolan et al. 2003). Gehrels et al. (2000) define a class of sources  as steady: a source is steady if the most significant detection of a  source in 3EG catalogue is for a timescale of years and if that particular flux is within $3\\sigma$ of the flux calculated for the full data set. From this classification, 48 unidentified sources at $|b|< 5^{\\circ}$  and 72 unidentified sources at $|b|> 5^{\\circ}$ are steady sources respectively. On the other hand, Grenier (2001) suggests another definition called persistent sources, which are those detected sources with a significance $\\sqrt{TS}>4$ for every observation at $|b|>2.5^{\\circ}$. If we assume that the sources with a significance $\\sqrt{TS}>5$ at $|b|\\le 2.5^{\\circ}$ are persistent sources, there are 40 (45) persistent unidentified EGRET sources at $|b|< 5^{\\circ}$ ($|b|> 5^{\\circ}$).  However, the spectral properties of  medium latitude sources are significantly softer, fainter and have a steeper logN-logS function than those at low latitudes (Gehrels et al. 2000). It has been suggested that they are associated with recent supernovae in the nearby Gould Belt (Grenier 1997; Gehrels et al. 2000; Grenier 2000). Their natures remain as mystery. Harding and Zhang (2001) used the polar cap models (Daugherty \\& Harding 1996; Harding \\& Muslimov 1998) to investigate whether $\\gamma$-ray pulsars viewed at a large angle to the neutron star magnetic pole could contribute to unidentified EGRET sources in the medium latitudes associated with the Gould Belt. They suggest that the off-beam $\\gamma$-rays come from high-altitude curvature emission of primary particles can radiate over a large solid angle and have a much softer spectrum than that of the main beams, and at least some of radio-quiet Gould Belt sources detected by EGRET could be such off-beam $\\gamma$-ray pulsars.  Recently, the brightest of unidentified EGRET sources in the medium latitude, 3EG J1835+5918, is strongly suggested as the second Geminga-like pulsar ( radio-quiet pulsar) by multi-wavelength observations including Chandra, HST and Jodrell Bank (Mirabal et al. 2000; Mirabal \\& Halpern 2001; Halpern et al. 2002). The X-ray spectrum can be described by two components: a soft thermal X-ray spectrum with a characteristic temperature ($T_{\\inf} \\approx 3\\times 10^5$K and a power law hard tail with a photon index $\\gamma \\approx $2, which closely resembles to the X-ray spectrum of Geminga. The repeated radio observations at Jodrell Bank did not show any periodicity. The X-ray data suggests that the distance is between 250-800pc, which is consistent with the distance to the Gould Belt.  In this paper, we use the revised outer gap models (Zhang et al. 2004) to investigate emission properties of $\\gamma$-ray pulsars in the Galaxy as well as in the Gould Belt. This revised outer gap model is based on the original outer gap models (Cheng, Ho \\& Ruderman 1986a, {\\bf{hereafter CHR I}}, 1986b; Zhang \\& Cheng 1997) but takes into account the effect of the inclination angle ($\\alpha$), which is the angle between the magnetic axis and the rotation axis, in determining the size of the outer gap,which is defined as the ratio between the dimension of the gap perpendicular to the magnetic field and the light cylinder radius. The revised model also takes into account the fact that the typical radiation region of the outer gap is not necessary at half of the light cylinder instead it should be better represented by an appropriate average over the entire outer gap. This effect is particular important for old pulsars. If the outer gap is only represented at half of the light cylinder, then the outer gap is assumed to be turn off when the gap size at this region is larger than unity.  This new model has taken the entire active region of the outer gap into account. As long as the gap size is less than unity in some parts of magnetosphere, the outer gap still exist. This effect allows some pulsars with appropriate combination of $\\alpha$, $P$ and $B$, can maintain their outer gaps until a few million years old. These pulsars are able to move up to high galactic latitude and their ages make them weak $\\gamma$-ray sources. Furthermore, we will show that these pulsars will emit softer spectra and their relation between $\\gamma$-ray luminosity and spin-down power differs from that of younger pulsars located in the galactic plane. We organize the paper as follows. In section 2, we review the new outer gap models. In section 3, we study the $\\gamma$-ray emission properties of pulsars by using the new self-consistent outer gap model. In section 4, we describe a Monte Carlo simulation method to determine the statistical distributions of radio-quiet and radio-loud $\\gamma$-ray pulsars. In section 5, we summarize the simulation results and discuss their implications. Finally, a brief conclusion is given in section 6. ",
        "conclusions": "}  We have studied the $\\gamma$-ray emission properties of pulsars in the Galaxy as well as in the Gould Belt by using a new self-consistent outer gap model, which includes the effects of inclination angle and average properties of outer gap. We have found that this new model can produce more mature $\\gamma$-ray pulsars, which have longer period, smaller inclination angle and typical age of one million years old, than the old model. In fact the mature $\\gamma$-ray pulsars dominate in the high galactic latitude. The spectra of these mature pulsars are significantly softer and weaker than the young pulsars in the galactic plane. It is because the spectral hardness is roughly determined by the typical photon energy given in equation (1), which characterizes the position of spectral break. The explicit form of how the typical photon energy depends on the inclination angle is given in equations (28) and (38) of Zhang et al. (2004). Although the form is quite complicated, we can roughly understand why smaller inclination angle gives smaller typical photon energy and hence {\\bf {softer}} spectrum by equation (1). It is because the typical position of the outer gap is the null surface where becomes larger for smaller inclination angle. From the simulation results, it is not clear why $\\gamma$-ray pulsars in high latitude should have smaller inclination angle. It is because all of these high latitude pulsars are older pulsars, which have longer period. Without the compensation of the effect of inclination angle, the fractional size of the outer gap $f$ is larger than unity ( cf. equation 3). In order to maintain the outer gap size less than unity, $\\eta$ in equation (4), which contains the effect of inclination angle must be small. In fact $\\eta$ becomes small for small inclination angle (cf. Fig. 1 of Zhang et al. 2004). Therefore high latitude gamma-ray pulsars tend to have smaller inclination angle and harder spectrum. On the other, this seems to imply that the confirmed radio-loud $\\gamma$-ray pulsars in the galactic plane should have larger inclination. This statistical correlation has not yet been established. But we would like to point out that the values of inclination angle of pulsars are very difficult to be measured accurately and there are only 7 confirmed $\\gamma$-ray pulsars. So it is very difficult to find such correlation in radio-loud $\\gamma$-ray pulsars. However, we do believe that this correlation should be there when more radio-loud $\\gamma$-ray pulsars are discovered in the galactic plane.   We have also used a Monte Carlo method to simulate the properties of the $\\gamma$-ray pulsar population in the Galaxy as well as in the nearby Gould Belt in terms of the revised outer gap models. The initial magnetic field, spatial and velocity distributions of the neutron star at birth which are obtained by the radio pulsar statistical studies have been used in our simulations. We have obtained the spatial, distance, period, age, magnetic field, inclination angle, photon flux distributions of the radio-loud and radio-quiet $\\gamma$-ray pulsars. We find that the properties of $\\gamma$-ray pulsars between $\\mid b \\mid < 5^{\\circ}$ and $\\mid b \\mid > 5^{\\circ}$ are very much different. Galactic plane $\\gamma$-ray pulsars are younger, shorter in period and have larger inclination angles. They satisfy different relation between $L_{\\gamma}$ and $L_{sd}$. The high latitude $\\gamma$-ray pulsars satisfy $L_{\\gamma}\\propto L^{0.6}_{sd}$ but the galactic plane pulsars satisfy $L_{\\gamma}\\propto L_{sd}^{0.3}$.   The present model predicts very similar numbers of $\\gamma$-ray pulsars as old models  but many more high latitude $\\gamma$-ray pulsars than old models (Cheng \\& Zhang 1998; Zhang, Zhang \\& Cheng 2000). It is because the new model allows the outer gap can exist about a few million years old for appropriate combination of B,P and $\\alpha$, so old $\\gamma$-ray pulsars can move up to high galactic latitude. Furthermore, the nearby Gould Belt also contribute significant number of $\\gamma$-ray pulsars in high latitude. Torres et al. (2003) have predicted that AGILE can detect more $\\gamma$-ray pulsars (for a general review of AGILE cf. Tavani et al. 2001). Perhaps AGILE can provide some clues to differentiate polar gap model predictions (Gonthier et al. 2002) and outer gap model predictions, and GLAST makes the final verdict. Finally we would like to remark that the better predictions on how many $\\gamma$-ray pulsars should be detected by AGILE and GLAST should include a better expression of $\\gamma$-ray solid angle, which depends on the inclination angle as well as the outer gap size. However, in order to have a reliable expression of $\\gamma$-ray solid angle, we need to carry out three-dimensional model calculation and come up with an approximate expression from Monte Carlo simulation."
    },
    "0403/astro-ph0403035_arXiv.txt": {
        "abstract": "{ As a part of an on-going program to explore the signature of $p$-modes in solar-like stars by means of high-resolution absorption line spectroscopy, we have studied four stars ($\\alpha$\\,CMi, $\\eta$\\,Cas A, $\\zeta$\\,Her A and $\\beta$\\,Vir). We present here new results from two-site observations of Procyon A acquired over twelve nights in 1999. Oscillation frequencies for $l$=1 and 0 (or 2) $p$-modes are detected in the power spectra of these Doppler shift measurements. A frequency analysis points out the difficulties of the  classical asymptotic theory in representing the $p$-mode spectrum of Procyon A. ",
        "introduction": "Precise measurements of the frequencies and amplitudes of the solar five-minute oscillations have led to detailed inferences about the solar internal structure \\citep[cf. review paper by][]{Christ02}. Because of the wealth of knowledge obtained from helioseismology there is a strong impetus to use the pulsation analysis techniques to probe the interiors of solar-like stars. From the ground, the application of seismic techniques to these stars is difficult because of extremely small variations in intensity and velocity associated with $p$-modes and the  need to have an adequate temporal coverage to resolve the modes. To detect such small amplitudes, photometric methods \\citep{Gilliland93} suffer from scintillation noise unless conducted from space, e.g. MOST \\citep{Walker03,Matthews02} or the future COROT \\citep{Baglin03} and Eddington \\citep{RoxburghFavata03} missions. Recent developments in high-precision spectrometric techniques and, in particular, Doppler spectroscopy have led to the detection of $p$-mode frequencies in individual solar-like stars \\citep[e.g.,][]{Martic99,Bedding01,Martic01b,BouchyCarrier02,Frandsen02,Kjeldsen03,CarrierBourban03}. A recent review of these and other measurements has been given by \\citet{BeddingKjeldsen03}.   One of the most interesting target for solar-like seismology is Procyon A ($\\alpha$\\,CMi, HR 2943, HD61421), a F5\\,IV star with $m_{\\rm v}$=0.363 at a distance of only 3.53\\,pc. Its fundamental stellar parameters are now well determined. Procyon A is in a 40-year period visual binary system, sharing the systemic velocity with a white dwarf. Procyon's astrometric orbit was recently updated by \\citet{Girard00}. With a parallax, $\\Pi = 283.2 \\pm 1.5$ mas, they estimated a mass of $1.470\\pm0.045 M_{\\odot}$, which agrees well with previous excellent predictions from stellar interior models \\citep{GuentherDemarque93}. In addition, \\citet{Mozurkewich91}, using optical interferometry, measured a stellar angular diameter $\\theta$ = 5.51$\\pm 0.5$\\,mas. Adopting the very precisely measured parallax by \\textit{Hipparcos}, \\mbox{$\\Pi = 285.93 \\pm\\,0.88$\\,mas}, \\citet{Prieto02} derived a slightly lower mass of 1.42\\,$\\pm\\,0.06 \\rm\\thinspace M_{\\odot}$, a radius  R/$\\rm\\thinspace R_{\\odot} = 2.071\\,\\pm\\,0.02$ and a gravity log\\,$g=3.96\\,\\pm0.02$\\,(in cgs units).  An exploratory single-site Doppler observation of Procyon A by \\citet{Martic99}, hereafter Paper I, revealed the presence of several frequencies that fell in a comb-like pattern with roughly equal spacing of $55\\,\\mu$Hz in the region of excess power around 1\\,mHz in the power spectrum. Because of the large daytime gaps in these single-site measurements, it was not possible to identify these frequencies unambiguously. Note however that the set of  frequencies reported in Paper I were determined very close ($\\leq 2\\,\\mu$Hz) to the model frequencies published independently by \\citet{Chaboyer99}.  Following the findings of these exploratory studies we have carried out a two-site campaign of Procyon A to identify as many $p$-modes as possible without alias ambiguities. In section~\\ref{section_obs}, we describe the context of the observing runs, in section~\\ref{section_period} we present the periodograms of the time sequences from coordinated two-site measurements. In  sections~\\ref{section_echdiag} and \\ref{section_modes}, we outline the different techniques we used to search for the $p$-modes in the power spectra. In the last sections, we present the principal mode characteristics (large and small separations, amplitudes) and compare the results with the existing models. ",
        "conclusions": ""
    },
    "0403/astro-ph0403279_arXiv.txt": {
        "abstract": "{ In an \\xmm\\ raster observation of the bright Local Group spiral galaxy \\m33\\ we study the population of X-ray sources (X-ray binaries, supernova remnants) down to a 0.2--4.5 keV luminosity of \\oergs{35} -- more than a factor of 10 deeper than earlier ROSAT observations. EPIC hardness ratios and optical and radio information are used to distinguish between different source classes. The survey detects 408 sources in an area of 0.80 square degree. We correlate these newly detected sources with earlier \\m33\\ X-ray catalogues and information from optical, infra-red and radio wavelengths. As \\m33\\ sources we detect 21 supernova remnants (SNR) and 23 SNR candidates, 5 super-soft sources, and 2 X-ray binaries (XRBs). There are 267 sources classified as hard, which may either be XRBs or Crab-like SNRs in \\m33\\ or background AGN. The 44 confirmed and candidate SNRs more than double the number of X-ray detected SNRs in \\m33. 16 of these are proposed as SNR candidates from the X-ray data for the first time. On the other hand, there are several sources not connected to \\m33: five foreground stars, 30 foreground star candidates, 12 active galactic nucleus candidates,  one background galaxy and one background galaxy candidate. Extrapolating from deep field observations we would expect 175 to 210 background sources in this field. This indicates that about half of the sources detected are sources within \\m33. ",
        "introduction": "\\begin{table*} \\begin{center} \\caption[]{\\xmm\\ \\m33\\ observation log (proposal numbers 010264 and 014198).} \\begin{tabular}{rllrrlrlrlr} \\hline\\noalign{\\smallskip} \\hline\\noalign{\\smallskip} \\multicolumn{1}{c}{Field} & \\multicolumn{1}{c}{Obs. id.} &\\multicolumn{1}{c}{Obs. dates} & \\multicolumn{2}{c}{Pointing direction} & \\multicolumn{2}{c}{EPIC PN} & \\multicolumn{2}{c}{EPIC MOS1} & \\multicolumn{2}{c}{EPIC MOS2}  \\\\ \\noalign{\\smallskip} M33 & & & \\multicolumn{2}{c}{RA/DEC (J2000)} & \\multicolumn{1}{c}{Mode$^{+}$/Filter}  & \\multicolumn{1}{c}{T$_{exp}^{\\dagger}$} & \\multicolumn{1}{c}{Mode$^{+}$/Filter}  & \\multicolumn{1}{c}{T$_{exp}^{\\dagger}$} & \\multicolumn{1}{c}{Mode$^{+}$/Filter}  & \\multicolumn{1}{c}{T$_{exp}^{\\dagger}$}\\\\ \\noalign{\\smallskip} \\multicolumn{1}{c}{(1)} & \\multicolumn{1}{c}{(2)} & \\multicolumn{1}{c}{(3)} & \\multicolumn{1}{c}{(4)} & \\multicolumn{1}{c}{(5)} & \\multicolumn{1}{c}{(6)} & \\multicolumn{1}{c}{(7)} & \\multicolumn{1}{c}{(8)} & \\multicolumn{1}{c}{(9)} & \\multicolumn{1}{c}{(10)} & \\multicolumn{1}{c}{(11)} \\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} 1  & 0102640101 & 2000-08-04 & 1:33:51.0 & 30:39:37 & FF/medium &  6.29 & PW2/medium &  6.23 & PW2/medium &  6.22 \\\\ 1  & 0141980501 & 2003-01-22 & 1:33:51.0 & 30:39:37 & FF/medium &  1.23 & FF/medium  &  1.51 & FF/medium  &  1.51 \\\\ 1  & 0141980801 & 2003-02-12 & 1:33:51.0 & 30:39:37 & FF/medium &  7.25 & FF/medium  &  8.86 & FF/medium  &  8.87 \\\\ 2  & 0102640201 & 2000-08-04 & 1:34:40.0 & 30:57:48 & EFF/medium & 11.40 & FF/medium  & 15.24 & FF/medium  & 15.24 \\\\ 3  & 0102640301 & 2000-08-07 & 1:33:32.0 & 30:52:13 & EFF/medium &  3.80 & FF/medium  &  4.58 & FF/thin    &  4.58 \\\\ 3  & 0141980401 & 2003-01-24 & 1:33:32.0 & 30:52:13 & FF/medium &  0.00 & FF/medium  &  0.00 & FF/medium  &  0.00 \\\\ 4  & 0102640401 & 2000-08-02 & 1:32:51.0 & 30:36:49 & EFF/thick &  8.66 & FF/thick   & 12.99 & FF/thick   & 12.99 \\\\ 4  & 0141980601 & 2003-01-23 & 1:32:51.0 & 30:36:49 & FF/medium & 11.18 & FF/medium  & 12.56 & FF/medium  & 12.56 \\\\ 5  & 0102640501 & 2001-07-05 & 1:33:02.0 & 30:21:24 & FF/medium &  9.26 & FF/medium  & 11.76 & FF/medium  & 11.76 \\\\ 6  & 0102640601 & 2001-07-05 & 1:34:08.0 & 30:46:06 & FF/medium &  2.86 & FF/medium  &  3.42 & FF/medium  &  3.42 \\\\ 6  & 0141980301 & 2003-07-25 & 1:34:08.0 & 30:46:06 & FF/medium &  4.12 & FF/medium  &  4.82 & FF/medium  &  4.82 \\\\ 7  & 0102640701 & 2001-07-05 & 1:34:10.0 & 30:27:00 & FF/medium &  6.97 & FF/medium  & 10.96 & FF/medium  & 10.95 \\\\ 8  & 0102640801 & 2001-07-07 & 1:34:51.0 & 30:42:22 & FF/medium &     - & FF/medium  &  1.19 & FF/medium  &  1.19 \\\\ 8  & 0102642001 & 2001-08-15 & 1:34:51.0 & 30:42:22 & FF/medium &  8.65 & FF/medium  & 11.15 & FF/medium  & 11.15 \\\\ 9  & 0102640901 & 2001-07-07 & 1:34:04.0 & 30:57:25 & FF/medium &  2.42 & FF/medium  &  3.15 & FF/medium  &  3.15 \\\\ 9  & 0141980201 & 2003-07-11 & 1:34:04.0 & 30:57:25 & FF/medium &  3.74 & FF/medium  &  4.02 & FF/medium  &  4.02 \\\\ 10 & 0102641001 & 2001-07-08 & 1:33:07.0 & 30:45:02 & FF/medium &  1.25 & FF/medium  &  8.33 & FF/medium  &  8.37 \\\\ 10 & 0141980101 & 2003-07-11 & 1:33:07.0 & 30:45:02 & FF/medium &  5.45 & FF/medium  &  5.75 & FF/medium  &  5.77 \\\\ 11 & 0102641101 & 2001-07-08 & 1:32:46.0 & 30:28:19 & FF/medium &  8.10 & FF/medium  & 10.36 & FF/medium  & 10.36 \\\\ 12 & 0102641201 & 2000-08-02 & 1:33:38.0 & 30:21:49 & EFF/thick & 11.34 & FF/thick   &  3.67 & FF/thick   &  3.66 \\\\ 12 & 0141980701 & 2003-01-24 & 1:33:38.0 & 30:21:49 & FF/medium &  4.40 & FF/medium  &  5.72 & FF/medium  &  5.73 \\\\ 13 & 0102642101 & 2002-01-25 & 1:34:34.0 & 30:34:11 & FF/medium & 10.00 & FF/medium  & 12.27 & FF/medium  & 12.27 \\\\ 14 & 0102642201 & 2002-01-25 & 1:34:56.0 & 30:50:52 & FF/medium & 11.60 & FF/medium  & 13.87 & FF/medium  & 13.87 \\\\ 15 & 0102642301 & 2002-01-27 & 1:33:33.0 & 30:33:07 & FF/medium & 10.00 & FF/medium  & 12.27 & FF/medium  & 12.27 \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\label{observations} \\end{center} Notes:\\\\ $^{ +~}$: FF: full frame, EFF: extended full frame, PW2: partial window 2 \\\\ $^{ {\\dagger}~}$: Exposure time in units of ks after screening for high background (see text)\\\\ \\end{table*} The Local Group Sc galaxy \\m33\\ is located at a distance of 795 kpc \\citep[][ i.e. 1\\arcsec\\ corresponds to 3.9 pc and the flux to luminosity conversion factor is 7.6\\expo{49} cm$^2$]{1991PASP..103..609V}, is seen under a relatively low inclination of 56$^{\\circ}$ \\citep{1989AJ.....97...97Z} and the optical extent can be approximated by an inclination-corrected D$_{25}$ ellipse with large diameter of 64\\farcm4 and axes ratio of 1.66 \\citep{1991trcb.book.....D,1988ngc..book.....T}. With its moderate Galactic foreground absorption \\citep[\\nh $= 6\\times10^{20}$\\,cm$^{-2}$, ][]{1992ApJS...79...77S} \\m33\\ is ideally suited to study the X-ray source population and diffuse emission in a nearby spiral galaxy. The \\ein\\ X-ray Observatory detected diffuse emission from hot gas in \\m33\\ and 17 unresolved sources \\citep{1981ApJ...246L..61L,1983ApJ...275..571M,1988ApJ...329.1037T}. First ROSAT HRI and PSPC observations revealed 57 sources and confirmed the detection of diffuse X-ray emission which may trace the spiral arms within a 10' radius around the nucleus \\citep{1995ApJ...441..568S,1996ApJ...466..750L}. Combining all archival ROSAT observations of the field, \\defcitealias{2001A&A...373..438H}{HP01} \\citet[][ hereafter \\citetalias{2001A&A...373..438H}]{2001A&A...373..438H} found 184 X-ray sources within a 50' radius around the nucleus, identified some of the sources by correlations with previous X-ray, optical and radio catalogues, and in addition classified sources according to their X-ray properties. They found candidates for super-soft X-ray sources (SSS), X-ray binaries (XRBs), supernova remnants (SNRs), foreground stars and active galactic nuclei (AGN).  To follow up on these findings, we proposed as part of the telescope scientist guaranteed time (GT) program a deep homogeneous \\xmm\\ survey of \\m33\\ with a sensitivity of $10^{35}$\\,erg~s$^{-1}$ in the 0.5--10 keV band, a factor of ten deeper than previous surveys. The survey consisted of a raster of 15 pointings of about 10 ks each. The directions were selected in a way that each area within the \\m33\\ optical D$_{25}$ extent was covered by the \\xmm\\ EPIC detectors at least three times with the medium filter. The survey was complemented by an additional AO2 program, initiated to fill observation gaps due to high detector background, or, because the thick filter was used during the GT observations. First results of this survey were presented in \\citet{2003AN....324...85P}.  Two \\m33\\ sources are known for their outstanding X-ray properties \\citep{1989ApJ...336..140P}. The brightest source \\citep[X--8 in the nomenclature of][ luminosity of $10^{39 - 40}$\\,erg~s$^{-1}$]{1981ApJ...246L..61L} is the most luminous X-ray source in the Local Group of galaxies and coincides with the optical center of \\m33. Its time variability \\citep{1997ApJ...490L..47D}, its point-like nature seen by ROSAT HRI \\citep{2000immm.proc..149P} and \\chandra\\ \\citep{2002MNRAS.336..901D} and its X-ray spectrum best described by an absorbed power law plus disk blackbody model \\cite[e.g. ][]{2001ncxa.conf..300E,2003ApJ...583..758L,fosch2003}, point towards a black hole XRB. A possible periodicity of 106 days was not confirmed in later observations \\citep[see ][]{2001A&A...368..420P}. The second source (\\x7) is an eclipsing high mass XRB  (HMXB) with a binary period of 3.45 d and possible 0.31 s pulsations, discovered in ROSAT PSPC and HRI observations \\citep{1993ApJ...418L..67S,1997ApJ...490L..47D, 1999MNRAS.302..731D,1997AJ....113..618L}. \\citet{px7} analyzed \\xmm\\ and \\chandra\\  data and determined improved binary parameters, however, they could not confirm the proposed pulsations. In a special analysis of DIRECT data, they identified an O7I star of 18.89 mag in V, as the optical counterpart, which shows the ellipsoidal heating light curve of a HMXB with the X--7 binary ephemeris. X--7 was the most distant eclipsing XRB until the detection of another such source in \\xmm\\ and \\chandra\\ observations of the starburst galaxy \\n253, which is located at more than three times  the distance of \\m33\\ \\citep{2003A&A...402..457P}.  Here we present merged medium and thin filter images for the three EPIC instruments, in five energy bands, using only times of low background from our \\xmm\\ survey of \\m33, as well as a source catalogue and source population study based on these images. An analysis of the individual pointings with less stringent background rejection and using all filters is in progress, and will result in a source catalogue with improved positions (taking care of individual pointing offsets), that will also provide information on source variability and spectral characteristics. ",
        "conclusions": "\\begin{figure} \\caption[]{\\xmm\\ EPIC 0.2-2 keV band contours of (6,8,14,32,86)$\\times 10^{-5}$ ct s$^{-1}$ overlayed on an optical image of \\m33\\ extracted from the digitized sky survey (DSS2) red plates. The image scale is marked. \\label{opt}} \\end{figure} Besides the large number of point-like sources, there was a strong component of unresolved X-ray emission in the energy band below 2 keV. The overlay of X-ray on the DSS2 red image (Fig.~\\ref{opt}) clearly showed the strong correlation with the inner disk and the southern and NW spiral arm. While some of this emission certainly is due to unresolved point sources, there is a significant contribution from hot gas in the disk which is more concentrated in the spiral arms and \\HII\\ regions like NGC~604 12\\arcmin\\ NE and BCLMP 290 15\\arcmin\\ NW of the nucleus. A detailed analysis of this emission will be presented in a separate paper.  In the \\xmm\\ source catalogue we presented 408 sources in the \\m33\\ field, 35 of which could be identified or classified as foreground stars and 14 as background AGN or galaxies. As sources in \\m33\\ we identified or classified 5 SSS, 44 SNRs, and 2 XRBs. 267 sources were classified as hard and may be either XRBs or Crab-like SNRs in \\m33\\ or background AGN. There remained 43 sources unclassified.  We can estimate the contribution of background sources to the source catalogue extrapolating from X-ray deep field number counts in the 0.5--2.0 keV and 2.0--8.0 keV bands \\citep[see e.g.][]{2001AJ....122.2810B,2002ApJ...566..667R}. We used rough sensitivity maps for our survey created by the SAS task {\\tt esensmap} for the 0.5--2.0 keV and 2.0--4.5 keV bands and an estimate for additional absorption in the line-of-sight. One expects about 175 to 210 background sources detected in our field-of-view, if we assume an overlap of sources detected in the hard and soft band of 75\\% or only 50\\%. Fourteen of the background sources have already been identified (see Sect. 6.2). Taking these numbers into account, up to 50\\% of the hard or unclassified sources should be sources within \\m33\\ (mostly XRBs or Crab-like SNRs).  We compared the \\xmm\\ source catalogue with the ROSAT catalogue of HP01. While most of the bright ROSAT sources were also detected by \\xmm\\, many of the sources, only detected with the ROSAT HRI with a detection likelihood below 10, were not detected by \\xmm\\ and seem to be spurious detections. From the seven ROSAT sources classified as SSS, we did not detect four. As expected for this source class, these might be transients, that were not active during the \\xmm\\ observations (HP01 sources 43, 69, 76, 82). The remaining three seem to be false identifications: HP01 source 61 is now classified by \\xmm\\ as a SNR, 84 as a foreground star and 124 as a hard source. Two ROSAT correlations with radio sources are no longer valid with the improved \\xmm\\ positions (HP01 sources 63 and 147).  HP01 identified and classified 17 SNRs. The \\xmm\\ source catalogue marks 21 identifications and 23 classifications. While the HR criteria seem to be rather robust for thermal SNRs, the identification of Crab-like remnants is doubtful and can only be secured by measurement of source extent. This may be possible with \\xmm\\ EPIC for very extended SNRs (optical extent of SNR in \\m33\\ vary from 9\\arcsec\\ to 126\\arcsec\\ according to \\citet{1998ApJS..117...89G}). Most of the sources, however, will need dedicated \\chandra\\ on-axis observations for a final decision on the extent. The \\chandra\\ count rates will not be sufficient to get good spectra and a combined \\xmm\\ and \\chandra\\ analysis will be important for this class of sources.  The importance of combined  \\xmm\\ and \\chandra\\ data for \\m33\\ XRB research was demonstrated in \\citet{px7}. This type of analysis extended to all brighter sources will be  the subject of a future paper."
    },
    "0403/astro-ph0403179.txt": {
        "abstract": "In this paper we comment on the possibility for intermittent jet activity in quasar 3C 273 on different time-scales. We propose, that striking morphology of the large-scale radio jet in this source, as well as the apparent lack of its counterpart on the opposite side of the active center, may be explained in a framework of a restarting jet model. In particular, we propose that 3C 273 radio source is intrinsically two-sided, and represents an analogue of double-double radio galaxies, but only inclined at a small angle to the line of sight. In this case, the apparent one-sideness of the kiloparsec-scale radio structure may be due to combined Doppler and time-travel effects alone, if the 3C 273 large-scale jet itself is relativistic and matter-dominated. We also propose, that knotty morphology of the discussed jet, which is observed now additionally at optical and X-ray frequencies, indicates modulation in the jet kinetic power. This, together with the variability of the jet at small (parsec) scales, indicates that the jet activity in 3C 273, and possibly in other similar sources, is variable/modulated/intermittent over many different time-scales. ",
        "introduction": "3C 273 is the well-known quasar at the redshift $z = 0.158$, discovered together with its famous large-scale jet at the very beginning of quasar optical researches \\citep[see a review by][]{cou98}. At radio frequencies it appears as the one-sided, core-dominated source of non-thermal radiation, with the radio jet extended up to a few tens of kpc from the core, similarly to its optical counterpart. The jet is also prominent at X-ray frequencies, although the origin of this radiation is still being debated (in contrast with the radio-to-optical emission, whose synchrotron nature is surely established). Until now, neither the jet composition nor its velocity on the kpc-scales is known. In addition, the unusual morphology of the radio structure -- lack of an extended radio cocoon in particular -- is not easy to interpret, suggesting, in the opinion of some authors like \\citet{cla86}, Poynting flux-dominated flow \\citep[see in this context also][]{lin89,kom99}. Finally, the existence of the counterjet in 3C 273 is also a controversial issue, because the jet-counterjet brightness asymmetry at the jet terminal point is exceptionally high, with a lower limit $10000:1$ \\citep{dav85,con94}. At the working surface of the jet only subrelativistic bulk velocities, and therefore minor Doppler-hiding effects, are expected \\citep[e.g.,][see also section 2.3 below]{beg84}. Hence, some authors \\citep{con81,dav85} have concluded that 3C 273 is an intrinsically one-sided jet source, that seems to be consistent with the lack of \\emph{any} radio feature (like some extended lobe, for example) on the 3C 273 counter-side.  In this paper we comment on the possibility for intermittent jet activity in quasar 3C 273 on different time-scales. We propose that the striking morphology of the large-scale jet in this source, as well as the apparent lack of its counterpart on the opposite side of the active center, may be explained in a framework of restarting and modulated jet model. In particular we propose that 3C 273 radio source is intrinsically two-sided, and represents an analogue of double-double radio galaxies \\citep{sch00a}, but only inclined at small angle to the line of sight, and in addition that the 3C 273 large-scale jet itself represents a continuous flow which is relativistic, matter-dominated and highly modulated in kinetic power. We discuss how present observations support this idea, and how the future ones can eventually verify it. Section 2 contains a brief presentation of the 3C 273 jet. In Section 3 we consider the proposed model of intermittent jet activity in this source. Discussion and final conclusions are presented in the last Section 4. ",
        "conclusions": "In this paper we propose that the large-scale jet in 3C 273 is an analogue of the inner lobes observed in double-double radio galaxies. This model can explain the narrowness of the radio cocoon in the considered source, as well as the lack of the counterjet, counter-lobe and the counter-head on the opposite side of the active center, if the jet inclination on the large scales is small, $\\theta \\sim 10^{\\circ} - 20^{\\circ}$, the jet is matter-dominated and relativistic up to its head, $\\Gamma_{\\rm jet} = 5 - 10$, and if the original intragroup medium surrounding 3C 273 is rarefied by previous jet activity by a factor of $100 - 1000$. If this is the case, then 3C 273 constitutes a quasar-counterpart to a few known double-double sources, which are inclined at much larger angles to the line of sight, $\\sim 90^{\\circ}$ \\citep[see][]{sch00a,sch00b,sar02,sar03}. Let us mention that the intermittent case between 3C 273 and other `classical' double-doubles can be radio galaxy 3C 219. The jet in this source shares many morphological similarities to the 3C 273 jet, and is considered by some authors as the example of the restarting (`partial') beam \\citep{bri86,cla92}. The overall radio morphology of 3C 219 was recently interpreted as a double-double one, but only observed at intermittent jet viewing angles \\citep{sch00a}. Interestingly, on the opposite side of the 3C 219 active center a single radio knot was detected and subsequently identified with the counter-head of the restarting counter-jet \\citep{cla92}. The ratio of the projected distances of the head and the counter-head from the 3C 219 core is $\\sim 3.6$, which can be compared with the appropriate value for double-double radio galaxies, $\\sim 1$, and for 3C 273, as estimated in this paper, $> 11$.  Besides the DDRGs, decaying activity of the extragalactic radio sources manifests in formation of `radio relics' found in rich clusters of galaxies. However, we believe that the analogy between 3C 273 and the DDRGs is more appropriate than the analogy between this object and cluster radio sources, for the following reasons. First, we note unusual narrowness of the large-scale radio jet/cocoon in 3C 273, which resembles morphology of the inner lobes of DDRGs. Second, 3C 273 is located in poor group of galaxies (similarly to the DDRGs) and not in rich cluster like radio relics considered by, e.g., \\citet{dey03}. As the evolution of radio lobes depends crucially on the properties of the surrounding medium, we believe that the model by \\citet{kai00} regarding extended outer lobes of DDRGs is more appropriate to adopt in this paper than the models regarding radio relics observed in clusters \\citep[see in this context][]{kom94,kai02}. Also, the DDRGs -- exhibiting in addition `pure' FR II morphology -- constitutes the \\emph{direct} proof for the intermittent/recurrent jet activity in powerful radio sources. Ghost radio cavities found in galaxy clusters are widely interpreted as the remnants of ceased radio activity, but there is no observed case for the new-born FR II jet propagating within the cluster radio relic \\citep[but see][]{jon01,hei02,ven04}.  We estimated the age of the large-scale jet in 3C 273 as $\\sim 10^5$ yrs. This would be then the time-scale of the jet active period. Unfortunately, the time-scale of the quiescence epoch before the present activity cannot be evaluated at the moment, as we did not detect its outer lobes until now (but see below). In analogy to the double-double radio galaxies one can only suggest that it is $\\geq 10^7$ yrs. We also propose, that knotty morphology of the discussed jet indicate modulation in the jet kinetic power on the time scale $\\sim 10^4$ yrs \\citep[see][]{sta04}. This, together with the VLBI jet variability time scale of the order of years, indicates that the jet activity in 3C 273, and possibly in other similar sources, is variable/modulated/intermittent over many decades of the source lifetime. This may have several potential consequences for the jet formation models \\citep[see in this context also][]{rey97,sie97,mar03}. We note, that recent coronographic observations of quasar 3C 273 with the Advanced Camera for Survays \\citep{ma03} revealed complicated morphology of the host galaxy, with the complex structure of filaments and clumpy regions, spiral-shaped plume, dust lanes, etc. It was suggested, that the possible explanation for such a complexity could be a relatively recent merger of 3C 273 with a massive companion galaxy. It is then possible, that such an eventual merger event started the epoch of the \\emph{recurrent} jet activity in the considered source, resulting in formation of the extended outer lobes some $10^7 - 10^8$ yrs ago, and next of the present $10^5$-year-old jet \\citep[see in this context][]{mer02,gop03}.  During the review process for this paper the author became aware of the new, unpublished $1.4$ GHz map of 3C 273 at $4''$ FWHM resolution with the dynamic range $30,000$ \\citep{per04}. It shows for the very first time some faint extended radio emission surrounding the 3C 273 quasar, enveloping its large-scale jet and also present at the counter-jet side. Interestingly, this extended radio structure seems to be elongated in the direction of the large-scale jet, as expected in the case of double-double radio morphology. Hence, in the near future, with the new data published, the proposed model can be verified and enriched by a more detailed analysis of the old lobe parameters."
    },
    "0403/astro-ph0403086_arXiv.txt": {
        "abstract": "We present optical candidates for 75 X-ray sources in a $\\sim 1$ deg$^2$ overlapping region with the medium deep ROSAT survey (Molthagen et al. 1997). These candidates are selected using the multi-color CCD imaging observations made for the T329 field of the Beijing-Arizona-Taipei-Connecticut (BATC) Sky Survey, which utilizes the NAOC 0.6/0.9m Schmidt telescope with 15 intermediate-band filters covering the wavelength range 3360-9745 \\AA. These X-ray sources are relatively faint (CR $<< 0.2 s^{-1}$) and thus mostly are not included in the RBS catalog, they also remain as X-ray sources without optical candidates in a previous identification program carried out by the Hamburg Quasar Survey. Within their position-error circles, almost all the X-ray sources are observed to have one or more spatially associated optical candidates within them down to the magnitude $m_V \\sim 23.1$.  We have classified 149 of 156 detected optical candidates with 73 of the 75 X-ray sources with a new method which predicts a redshift for non-stellar objects, which we have termed the SED-based Object Classification Approach (SOCA).  These optical candidates include: 31 QSOs, 39 stars, 37 starburst galaxies, 42 galaxies, and 7 ``just'' visible objects.  Twenty-eight X-ray error circles have only one visible object in them: 9 QSOs, 3 normal galaxies, 8 starburst galaxies, 6 stars, and two of the ``just'' visible objects.  We have also cross-correlated the positions of these optical objects with NED, the FIRST radio source catalog and the 2MASS catalog.  Separately, we have also SED-classified the remaining 6011 objects in our field of view. Optical objects are found at the $6.5\\sigma$ level above what one would expect from a random distribution, only QSOs are over-represented in these error circles at greater than 4$\\sigma$ frequency.  We estimate redshifts for all extragalactic objects, and find a good correspondence of our predicted redshift with the measured redshift (a mean error of 0.04 in $\\Delta z$.  There appears to be a supercluster at z $\\sim$ 0.3-0.35 in this direction, including many of the galaxies in the X-ray error circles are found in this redshift range. ",
        "introduction": "Combined optical and X-ray data allow one to obtain information about the luminosity functions of various types of X-ray sources as well as their evolution with redshift. In turn, this information can be used to further constrain models for the production of the X-ray background at different flux levels.  Much effort has so far been made on the optical identification of the X-ray sources in the ROSAT/Bright Source (RBS) catalog (e.g., Voges et al. 1999; Rutledge et al. 2000) as well as of X-ray sources in some individual ROSAT deep survey observations (e.g., Lehmann et al. 2001).  Yet, it is often unknown how much of the detections that occur in X-ray error circles are real associations of optical counterparts to these X-ray sources, and how much of these associations are due to random chance.  To do this, one needs to have detected and identified all optical objects in a given image, and then see what percentages of these objects (QSOs, galaxies, stars) are then found near or within the areas covered by the X-ray error circles. This is precisely the kind of data we have for 75 X-ray sources detected with the ROSAT PSPC to a flux limit $S_x({\\rm 0.1-2.4 keV}) \\geq 5.3\\times10^{-14}$\\ \\fxu, in a 1 deg$^2$ field of view, as this field of view was also observed in our multicolor images for the Beijing-Arizona-Taipei-Connecticut Sky Survey (BATC survey).  The relevant X-ray and optical data are presented in \\S~2.  Details of the object classification procedure as well as selection of the X-ray candidates are given in \\S~3. Associated information that can be gleaned from these data are given in \\S~4. We summarize our results in \\S~5. ",
        "conclusions": "Based on the 15 color photometric observations, and SED-based object classification approach (SOCA), as well as the multiwavelength cross-correlations, we find 156 optical candidates, to a magnitude limit close to $V \\approx 23$, within or near the error circles for all 75 X-ray sources in our field of view. Among them are 31 QSOs (nine of which have spectroscopic confirmation), 37 starburst galaxies (2 of which have spectroscopic confirmation), 42 normal galaxies (including 3 possible galaxy groups), 39 stars (one of which is a BD object), and 7 just ``visible'' objects with no classifications (two of which are coincident with known radio sources). Two of the X-ray error circles have only just ``visible'' objects in them. We find 8 radio sources (out of 100 in our image) that are coincident with an optical object within the X-ray error circles, making it likely that many of these are the optical counterparts to these radio sources.  Separately, we have also SED-classified 6011 additional objects in our 3364 arcmin$^2$ field of view to the same apparent magnitude limit (i.e., not including those found near or in the X-ray error circles). Of these 6011 objects, 341 are QSOs, 1912 are stars, 1508 are starburst galaxies, and 2076 are galaxies and 174 are unclassified for a number of objective reasons.  The area of our X-ray circles subtends 31.52 arcmin$^2$, or 0.00937 of our full 3364 arcmin$^2$ field of view.  Considering the objects detected to objects that could be randomly found in these error circles, we have: $6.5\\sigma$ detections for all 149 classified objects, $4.7\\sigma$ for the QSOs, and from $2.8-3.2\\sigma$ for stars, starburst galaxies and galaxies. Twenty-eight error circles have only one object in or near them, including: 9 QSOs, 3 normal galaxies, 8 starburst galaxies, 6 stars and 2 ``just'' visible objects.  In sum, in this paper we perform an exercise that few have been able to do with their X-ray data. By being able to SED-classify all objects in our 3364 arcmin$^2$ field of view down to V $\\sim$ 23, we can ask ourselves what are the kinds of optical candidates we find within X-ray error circles to those randomly found in the field of view. question is that while all classified objects: QSOs, starburst galaxies, normal galaxies and stars are overrepresented in the X-ray error circles compared to a random distribution, it is only the QSOs that are highly statistically found in these X-ray circles.  Yet, at the same time, there are 6.5$\\sigma$ more objects within these X-ray circles than if randomly distributed in this image. So, while we know that one of the objects in our X-ray error circles is likely the X-ray source, in absence of independent knowledge, choosing which object it is is still more arbitrary than scientific."
    },
    "0403/astro-ph0403615_arXiv.txt": {
        "abstract": "We discuss spectral principal component analysis (SPCA) and show examples of its application in analyzing AGN spectra in both small and large samples.  It can be used to identify peculiar spectra and classify AGN spectra.  Its application to correlation studies of AGN spectral properties and spectral measurements for large samples is  promising. ",
        "introduction": "To study AGN spectra, one has to deal with many emission features as well as continuum.  Simple statistics often cannot handle the large number of measured parameters efficiently, therefore, a multivariate analysis is needed.  Principal component analysis (PCA) is one such powerful tool (see also Boroson 2004).  Suppose we have $n$ AGNs in a sample, each has $p$ measured parameters $X_1, X_2, \\cdots, X_p$. We can write {\\small \\[ \\begin{array}{ccccc} & X_1   & X_2   & \\cdots        & X_p   \\\\ QSO1    & x_{11}& x_{12}& \\cdots        & x_{1p}\\\\ QSO2    & x_{21}& x_{22}& \\cdots        & x_{2p}\\\\ \\vdots  &      \\vdots & \\vdots &        & \\vdots \\\\ QSOn    & x_{n1}& x_{n2}& \\cdots        & x_{np}\\\\ \\end{array}. \\] } \\noindent $X_1, X_2, \\cdots, X_p$ are unit vectors and $x_{ij}$ are data.  PCA defines a set of new orthogonal variables, principal components $W_j$ ($j=1, \\cdots, p$), which are linear combinations of the original variables\\\\ \\makebox[5in][c]{ $W_j = e_{j1}X_1 + e_{j2}X_2 + \\cdots + e_{jp}X_p ,$ } \\hfill (1) \\\\ An easy way to understand PCA is to consider it from the geometrical point of view (Francis \\& Wills 1999). PCA aims to define orthogonal principal axes in a multi-dimensional space, as shown in Fig.~1, where a correlation exists between measured parameters $X_1$ and $X_2$, and two principal components (PCs) are defined.  PC1 (or $W_1$) accounts for most variance (the correlation), PC2 (or $W_2$) is small and can be ignored. \\begin{figure} \\plotfiddle{ShangZfig1a.eps}{.8in}{0}{20}{20}{-142}{-64} \\plotfiddle{ShangZfig1b.eps}{.0in}{0}{20}{20}{0}{-40} \\caption {Geometrical point of view of PCA.  (a) Measured values.  (b) Mean subtracted values.  Two principal axes PC1 and PC2 are defined.} \\end{figure}    PCA can be applied to AGN spectra directly (Francis et al. 1992; Shang et al. 2003).  In this spectral principal component analysis (SPCA), each spectrum is divided into small wavelength bins, and the flux in each bin (e.g., $F_{\\lambda_i}$) is an input variable.  Eq.~(1) can be written as\\\\ \\makebox[5in][c]{ $W_j = e_{j1}F_{\\lambda_1} + e_{j2}F_{\\lambda_2} + \\cdots + e_{jp}F_{\\lambda_p} ,$ } \\hfill (2) \\\\ The result principal components can then be also represented as spectra, $SPC_j = [e_{j1}\\  e_{j2}\\   \\cdots \\  e_{jp} ]$, namely, the spectral principal components (SPCs). An original spectrum can be reconstructed by adding the weighted principal component spectra to the mean spectrum,\\\\ \\makebox[5in][c]{ $ Spectrum_i = Mean + w_{i1} {SPC_1} + w_{i2} {SPC_2} + \\cdots + w_{ip} {SPC_p}, $ } \\\\ where $w_{ij}$ are calculated from Eq.~2 for each object $i$, and are referred as the weights (or scores) of $SPC_j$. In other words, all original spectra are made from the principal component spectra. This implies that a limited number of SPCs can be used to measure spectra in large samples (see also Yip et al. 2004).  If there are strong linear relationships among the original variables (or wavelength bins), each of these relationships will be represented by a principal component.  Fewer principal components may be required to describe the total variation of a sample, thus providing a simpler description of the dataset (Francis et al. 1992).  Any principal component accounting for a significant fraction of the total sample variance might be related to one or more underlying fundamental physical parameters, giving some physical insight into the cause of the variations (Boroson \\& Green 1992, Boroson 2002). However, PCA is only a statistical tool and it is still up to the investigators to judge whether the resulting PCs have any physical meanings.  PCA is a linear analysis.  If there are strong non-linear relationships involved, PCA is not able to identify them directly and there will be crosstalk among the resulting principal components.  When a non-linear part of a relationship is not strong, PCA is still able to follow the ``linear'' trend of the relationship.  In SPCA, good redshifts are needed because line shifts can cause strong non-linear effects. Line width change also introduces non-linear relationships among binned fluxes, so caution is needed in interpreting SPCA results. ",
        "conclusions": ""
    },
    "0403/astro-ph0403109_arXiv.txt": {
        "abstract": "We compare the results of two--dimensional simulations to experimental data obtained at Los Alamos National Laboratory in order to validate the FLASH code.  FLASH is a multi--physics, block--structured adaptive mesh refinement code for studying compressible, reactive flows in various astrophysical environments.  The experiment involves the lateral interaction between a planar Ma=1.2 shock wave with a cylinder of gaseous sulfur hexafluoride (SF$_6$) in air.  The development of primary and secondary flow instabilities after the passage of the shock, as observed in the experiments and numerical simulations, are reviewed and compared.  The sensitivity of numerical results to several simulation parameters are examined.  Computed and experimentally measured velocity fields are compared.  Motivation for experimental data in planes parallel to the cylinder axis is provided by a speculative three--dimensional simulation. ",
        "introduction": "The impulsive acceleration of a material interface can lead to complex fluid motions due to the Richtmyer--Meshkov (RM) instability.  Here, the misalignment of pressure and density gradients deposits vorticity along the interface, which drives the flow and distorts the interface. At later times the flow may be receptive to secondary instabilities, most prominently the Kelvin--Helmholtz instability, which further increase the flow complexity and may trigger transition to turbulence.  Experimental investigations of impulsively accelerated interfaces have focused mainly on interfaces with single--mode perturbations, and on the case we consider here, shock--accelerated cylindrical gas columns~\\cite{jacobs93,tompre03,zoldi2002}.  The experiments are relatively inexpensive and turnaround times are short.  The challenges are diagnostics and repeatability: after the shock passes, the flowfield evolution is driven by flow instabilities and vortex dynamics, which are sensitive to the initial conditions and noise in the system.  Specifically, for the present case of a single shock--accelerated gas cylinder, the flow is dominated by a counter--rotating vortex pair.  Verification and validation are critical in the development of any simulation code, without which one can have little confidence that the code's results are meaningful.  The sensitivity and the complex evolution of the vortex pair are desirable properties for our primary purpose, which is to use the experiments to validate our simulation code.  A well--designed, well--characterized, and accurately diagnosed experiment is essential for validation.  FLASH is a multi--species, multi--dimensional, parallel, adaptive--mesh--refinement, fluid dynamics code for applications in astrophysics~\\cite{fryols00}.  Calder et al. discuss initial validation tests of the FLASH code~\\cite{calfry02}.  Herein we continue our validation effort by comparing FLASH simulations to an experiment performed at the Los Alamos National Laboratories~\\cite{tompre03,zoldi2002}. ",
        "conclusions": "To date we have made a large number of two--dimensional simulations to validate the FLASH code for problems dominated by vortex dynamics. While this is work in progress, we can make the following remarks: \\begin{itemize} \\item The overall morphology of the flowfield is captured by the simulations, but differences exist in the location and extent of small--scale structure. \\item Simulation velocity magnitudes lie within 10\\% of experimental values. \\item For vortex--dominated subsonic flows, FLASH users should be cautious regarding the choice of CFL number, mesh refinement criteria, and if PPM is used, contact steepening. \\end{itemize} Our three--dimensional simulation, despite issues with the initial conditions, suggests that three--dimensional effects might be important for this experiment, and measurements should be made parallel to the cylinder axis to examine the issue.  \\vspace*{-0.2in}"
    },
    "0403/astro-ph0403423_arXiv.txt": {
        "abstract": "When applied to deep photometric catalogs, the two-point angular correlation function, \\wt, is a sensitive probe of the evolution of galaxy clustering properties.  Here we present measurements of \\wt \\ as a function of $I_{\\rm AB}$ magnitude and $(R-I)$ color to a depth of $I_{\\rm AB}=24$ on scales $\\theta=7\\arcsec -3\\arcmin$, using a sample of $\\sim350,000$ galaxies covering 5 degrees$^2$ in total over 5 separate fields.  Using redshifts of 2954 galaxies in early DEEP2 Galaxy Redshift Survey data, we construct robust galaxy redshift distributions as a function of $I_{\\rm AB}$ and $R_{\\rm AB}$ magnitude and $(R-I)$ color for galaxies between $0<z<2$.  We constrain models of the redshift evolution of galaxy clustering and find that significant growth of clustering has occurred from $z\\geq1$ to $z=0$.  A model in which the comoving scale-length, $x_0$, evolves linearly with redshift, $x_0(z)=x_0(0)(1-Bz)$, fits the data better than the $\\epsilon$ model proposed by \\cite{Groth77}. The clustering properties depend strongly on observed $(R-I)$ color, with both the reddest and bluest galaxies exhibiting large clustering amplitudes and steeper slopes. Different observed $(R-I)$ color ranges are sensitive to very disparate redshift regimes.  Red galaxies with $(R-I)\\sim1.5$ lie in a narrow redshift range centered at $z\\sim0.85$ and have a comoving scale length of clustering at $z=0.85$ of $x_0=5.0 \\pm0.3$ \\mpch.  These galaxies have early-type spectra and are likely progenitors of massive local ellipticals.  The bluest galaxies with $(R-I)\\sim0$ appear to be a mix of star-forming galaxies, both relatively local ($z\\sim0.3-0.6$) dwarfs and bright $z>1.4$ galaxies, and broad-line AGN. We find that local blue dwarfs are relatively unclustered, with $x_0=1.6 \\pm0.2$ \\mpch. The $z>1.4$ blue galaxies have a larger clustering scale-length, $x_0\\gtrsim$5 \\mpch. ",
        "introduction": "Studies of galaxy clustering can constrain cosmological parameters as well as galaxy formation models, since the observed galaxy clustering strength depends on both the underlying dark matter distribution and the efficiency of galaxy formation inside dark matter halos.  While theory and simulations both provide detailed predictions for the dark matter distribution in various cosmological models, the processes of galaxy formation and evolution are less well-understood.  The redshift dependence of the galaxy clustering strength, in particular, can be used to constrain galaxy evolution theories if the cosmology is known.  Galaxy clustering is ideally studied in three dimensions, where the redshift of each galaxy is known.  However, cosmic variance dictates the need for very large data samples, probing a large volume of space and sampling multiple independent fields, in order to obtain statistically robust results.  As photometric data exist over larger areas and to fainter magnitude limits than are currently available for spectroscopic samples, galaxy clustering has long been measured from two-dimensional photometry alone.  The angular two-point correlation function, \\wt, measures the excess probability above random of finding a galaxy at a specific angle $\\theta$ from another galaxy.  This quantity is derived from counts of galaxy pairs as a function of their separation, and while relatively straightforward to measure, its interpretation is not trivial.  From the projected angular two-point correlation function, \\wt, one can infer the three-dimensional spatial two-point correlation function, $\\xi(r)$, if the redshift distribution of the sources in each sample is well-known and the redshift dependence of the clustering strength is known or assumed.  The two-point correlation function, $\\xi(r)$, is usually fit as a power law, $\\xi(r)=(r/r_0)^{-\\gamma}$, where $r_0$ is the scale-length of the galaxy clustering.  However, the assumed galaxy redshift distribution ($dN/dz$) has varied widely in different studies, such that similar observed angular clustering results have led to widely different conclusions.  A further complication is that each sample usually spans a large range of redshifts and is magnitude-limited, such that the mean intrinsic luminosity of the galaxies is changing with redshift within a sample.  As the clustering properties of galaxies are expected to depend on luminosity and have been shown to do so locally \\citep{Norberg01}, this can hinder interpretation of the evolution of clustering measured in \\wt \\ studies.  Many of the first useful measurements of large-scale structure were studies of angular clustering. One of the earliest measurements of the \\wt \\ statistic was the pioneering work of \\cite{Peebles75} using photographic plates from the Lick survey.  They found \\wt \\ to be well fit by a power law, \\wt$=A_{\\omega}\\theta^{\\delta}$, with a slope of $\\delta=-0.77 \\pm0.06$. Later studies using CCDs were able to reach deeper magnitude limits and found that fainter galaxies had a lower clustering amplitude. However, most of these surveys covered only one or two fields, typically with a field of view of $\\sim50$ arcmin$^2$ due to the small sizes of CCDs, and were strongly affected by cosmic variance \\citep[e.g.,][]{Efstathiou91, Roche93, Brainerd95, Hudon96, Woods97, McCracken00}.  More recently, wide-field cameras on intermediate-class telescopes have led to deep wide-field photometric surveys \\citep[e.g.,][]{Neuschaefer95,Roche99,Cabanac00,McCracken01,Wilson03}. A ground-breaking study was conducted by \\cite{Postman98}, who used a 4-m telescope with a 16 arcmin field of view to cover a much wider area by tiling together several pointings, surveying a contiguous 4 x 4 degree$^2$ field to a depth of $I_{\\rm AB}=24$, reaching to $z\\sim1$. \\cite{McCracken01} present a survey of 1 degree$^2$ to a depth of $I_{\\rm AB}=25.5$ over four independent fields using the UH8K camera on CFHT, and \\cite{Wilson03} use the same camera to image 1.5 degrees$^2$ over 3 separate fields to a depth of $I=24.0$.  Other recent studies employ accurate photometric redshifts to measure \\wt \\ for galaxy samples of a specific type over well-defined redshift ranges.  \\cite{Brown03} investigate the clustering of red galaxies in three redshift bins between $0.3<z<0.9$ in a 1.2 degree$^2$ field from the NOAO Deep Wide-Field Survey \\citep{Jannuzi99}, while \\cite{Budavari03} use data from the Sloan Digital Sky Survey (SDSS, \\cite{York00}) to construct a volume-limited sample with more than 2 million galaxies between $0.1<z<0.3$ with which to study galaxy clustering as a function of color and spectral type.  In this paper we present a survey undertaken with the CFH12K camera covering a total of 5 degrees$^2$ over 5 separate fields to a depth of $I_{\\rm AB}=24.0$. There is a wide diversity of adopted redshift distributions for magnitude-limited samples in the literature, which can lead to different interpretations of similar results. Here we use much improved $dN/dz$ distributions, relying on recent data from the DEEP2 Galaxy Redshift Survey.  Unlike previous studies, we require no spectral modelling in fitting for these $dN/dz$ distributions and no assumptions are made about the evolution of the galaxy population.  An outline of the paper is as follows. In Sections 2 and 3 we describe the observations and data reduction process.  As a large subset of this data was used for the DEEP2 Galaxy Redshift Survey photometry, we provide a detailed account of the data reduction pipeline. In section 4 we present galaxy counts in three photometric bands, while in Section 5 we measure \\wt \\ as a function of $I_{\\rm AB}$ magnitude, describe the $dN/dz$ distributions we use and discuss constraints on the evolution of galaxy clustering.  Section 6 parallels Section 5 but focuses on clustering as a function of $(R-I)$ color.  In this section we discuss in particular the strong clustering seen for the reddest and bluest galaxies and discuss the most likely interpretations of this effect.  We conclude in Section 7. ",
        "conclusions": ""
    },
    "0403/astro-ph0403437_arXiv.txt": {
        "abstract": "The first accurate measurements of Vela's proper motion were produced using the Parkes-Tidbinbilla Interferometer (PTI) (Bailes et al 1989). This 275 km baseline at 1.6~GHz gave a resolution of 140 mas and detected the proper motion and thus, by extrapolation, the birth site. We have extended this work to allow the third dimension to be derived via the parallax. We find the offset between the spin axis and the proper motion vector to be $9^o$, which limits acceleration time from the SN blast. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403601_arXiv.txt": {
        "abstract": "We report the results of \\xmm\\ observations of the X-ray bright, radio-loud blazar \\gb\\ at a redshift of $z=4.72$. We confirm the presence of soft X-ray spectral flattening at energies $\\lesssim0.7\\kev$ as reported in previous \\rosat\\ and \\bepposax\\ observations. At hard X-ray energies the spectrum is consistent with a power-law although we find the spectral slope varied between both \\xmm\\ observations and is also significantly different from that reported previously. Whilst we cannot rule-out intrinsic cold absorption to explain the spectral depression, we favour a dust-free warm absorber. Cold absorption requires a column density $\\sim1.4-1.6\\times10^{22}\\pcmsq$ whilst a warm absorber could have up to $\\sim10^{23}\\pcmsq$ and an ionization parameter $\\sim10^{2}$. The spectrum of \\gb\\ shows remarkable parallels with that of the $z=4.4$ blazar \\pmn, in which the available evidence is also most consistent with a warm absorber model. ",
        "introduction": "High-redshift quasars are important objects in the study of the early evolution of massive black holes and the interaction with their host galaxies. A number of radio-loud quasars have been found to show the spectral characteristics typical of blazars. \\gb\\, at $z=4.72$ \\citep{hook98,fabian97,fabian98,fabian99} displays spectral flattening at soft X-ray energies in \\rosat\\ \\citep{boller00} and \\bepposax\\ \\citep{fabian01_gb} data. Here we confirm the flattening with \\xmm\\ observations. Soft X-ray flattening has also been reported from \\rx\\ \\citep{yuan00} and \\pmn\\ \\citep{fabian01_pmn} at redshifts of $z=4.28$ and $z=4.4$ respectively. \\xmm\\ data shows only marginal evidence for flattening in the case of \\rx\\ \\citep{grupe04}. Conversely, the effect has been confirmed in \\pmn\\ \\citep{worsley04}. Spectral flattening is consistent with a trend seen in a number of studies of lower-redshift radio-loud quasars \\citep[e.g.][]{cappi97,fiore98,reeves00}.  The most consistent explanation for the flattening is intrinsic absorption with column densities of the order $10^{22-23}\\pcmsq$. Intergalactic absorption systems remain an unlikely possibility given their low metallicities. Furthermore, in order to account for the flattening the line of sight to the blazar would need to intercept two or more very high column density systems of which the probability is very low. Certainly this explanation cannot account for the spectral depressions seen in an increasing number of objects. An alternative explanation for flattening is a break in the underlying blazar continuum. Hard X-ray production in blazars is likely to be due to the inverse Compton scattering of seed photons by relativistic electrons in the jet. A spectral break can arise through a low energy cut-off in the electron population or by a peaked seed photon spectrum, however the break appears much too sharp for this to be plausible.  Given an intrinsic absorber hypothesis there remains much discussion about the nature of the absorber. X-ray data alone have been unsuccessful in breaking the degeneracy between a cold (neutral) absorber or one that is warm (ionized). Optical observations however, can play a crucial role in this regard. A number of the blazars showing soft X-ray flattening have optical spectra (corresponding to rest frame UV emission) which indicate very little dust along the line of sight or show any evidence of a Lyman-limit system (for the optical data for \\gb\\ and \\pmn\\ see \\citealp{hook98} and \\citealp{peroux01}). These point to a dust-free yet metal-enriched absorber where hydrogen has been ionized, consistent with the warm absorber such as that seen in Seyfert galaxies \\citep[e.g.][]{reynolds97,crenshaw99}. Strong intrinsic C\\thinspace\\textsc{iv} absorption is common in these galaxies and it is also a feature of \\pmn. ",
        "conclusions": "We have studied the X-ray spectrum of the high-redshift, radio-loud blazar \\gb, providing definitive confirmation of the soft X-ray spectral flattening as previously observed with \\rosat\\ and \\bepposax. The flattening is not consistent with Galactic or intergalactic absorption and we propose that the effect is due to an absorber that is intrinsic to the source, or to an underlying break in the blazar continuum.  A continuum break is highly speculative given the present lack of knowledge concerning emission processes in the nuclear regions. Continuum breaks are not seen in nearby radio-loud objects where they should be readily detectable and there is no satisfactory way to account for the sharpness of the break seen. An intrinsic absorber is therefore our favoured explanation. Cold absorption requires a column density $N_{\\textrm{H}}\\sim1.4-1.6\\times10^{22}\\pcmsq$, whilst a warm absorber could have a column density of up to $\\sim10^{23}\\pcmsq$ and an ionization parameter $\\xi\\sim10^{2}$. X-ray data alone are unable to distinguish between a warm or cold absorber model. Good constraints are expected from future optical observations which could place strong limits on the neutral hydrogen column density.  It is interesting to note the parallels with the blazar \\pmn\\ \\citep{worsley04} i.e. a relatively blue UV continuum with strong Ly-$\\alpha$ and C\\thinspace\\textsc{iv} emission lines as well as the flat hard X-ray spectral slope and the remarkably similar form of the soft X-ray depression. The \\pmn\\ X-ray data are consistent with both warm/cold absorption but optical observations rule out a cold absorber scenario. The lack of significant reddening in the UV spectra of both objects is consistent with a dust-free warm absorber.  Comparisons can also be drawn to the $z=3.27$ radio-loud quasar \\pks\\ which also shows definite soft X-ray absorption \\citep{ferrero03,fiore03}. Both warm and cold absorbers are consistent although ionized absorption is again most able to explain the lack of significant UV reddening. The implied iron abundance of a warm absorber however is at odds with the super-solar abundances implied by broad emission lines.  The spectrum of \\gb\\ are suggestive of an additional component contributing to an excess of emission in the $\\sim5-10\\kev$ rest-frame energy range. There are a number of possibilities for the emission including a pile-up in the electron energy distribution or a feature caused by bulk Comptonization, however the detection is marginal and further work is required."
    },
    "0403/astro-ph0403092_arXiv.txt": {
        "abstract": "We examine the fine structure lines $^{3}P_{1}\\rightarrow $$^{3}P_{0}$ (492~GHz)  and $  ^{3}P_2\\rightarrow $$^{3}P_1$  (809 GHz)  of neutral atomic carbon  as bulk molecular gas  mass tracers and  find that they can be good and on  many occasions better than $ ^{12}$CO transitions, especially at  high redshifts. The notion  of CI emission  as an H$_2$ gas mass tracer challenges the long-held view of its distribution over only a  relatively narrow  layer in the  CII/CI/CO transition  zone in FUV-illuminated  molecular  clouds.   Past  observations  have  indeed consistently pointed  towards a more  extended CI distribution  but it was  only recently,  with the  advent of  large scale  imaging  of its $^{3}P_{1}\\rightarrow  $$^{3}P_{0}$  transition,  that its  surprising ubiquity in molecular clouds has  been fully revealed.  In the present work we show that under  {\\it typical} ISM conditions such an ubiquity is  inevitable  because  of  well known  dynamic  and  non-equilibrium chemistry  processes   maintaining  a  significant   [C]/[$  ^{12}$CO] abundance  throughout Giant  Molecular Clouds  during  their lifetime. These processes are more  intense in star-forming environments where a larger ambient  cosmic ray  flux will also  play an important  role in boosting [C]/[$ ^{12}$CO].   The resulting CI lines can  be bright and effective   H$_2$   mass  tracers   especially   for  diffuse   ($\\sim 10^2-10^3\\rm  \\ cm^{-3}$)  gas while  in UV-intense  and/or metal-poor environments  their  H$_2$-tracing  capability diminishes  because  of large scale  CII production but nevertheless remains  superior to that of $ ^{12}$CO.  The best place to take full advantage of CI's capacity to trace H$_2$ is not in the low-$z$ Universe, where large atmospheric absorption at  492 and 809~GHz precludes routine  observations, but at high redshifts ($\\rm z\\ga 1$). ",
        "introduction": "The  use of  $  ^{12}$CO  and its  isotopologues  rotational lines  to measure H$_2$  gas mass and physical  properties in galaxies  is now a well-established  technique  (e.g.  Dickman  et  al.  1986;  Elmegreen 1989; Bryant  \\& Scoville  1996) employed successfully  throughout the local Universe  (z$\\la 0.3$) (e.g.  Sanders, Scoville  \\& Soifer 1991; Young \\& Scoville 1991; Solomon  et al.  1992; Downes \\& Solomon 1998) and at increasingly  high redshifts (e.g.  Brown \\&  Vanden Bout 1991; Downes et  al.  1992;  Omont et  al.  1996; Frayer  et al.   1998). In high-z observations  usually only high-J CO  lines $\\rm J+1\\rightarrow J$,  ($\\rm J+1\\geq  4$) are  accessible  for evaluating  the mass  and physical conditions of H$_2$  gas, thus omitting cooler ($\\rm T_{kin}< 30$  K)  and/or sub-thermally  excited  gas  ($\\rm n<10^4\\  cm^{-3}$). Recently  observations   of  the   lowest  $  ^{12}$CO   J=1--0,  2--1 transitions in  several high redshift objects have  been made possible with the Very Large Array  (e.g. Carilli, Menten, \\& Yun 1999; Carilli et al.  2002a, 2002b) and  in some cases thermalized transitions up to fairly  high  rotational levels  (e.g.   J=5--4)  have been  observed. However, differential  lensing can complicate  standard interpretation of  line  ratios by  preferentially  amplifying  the  emission of  the usually more spatially confined warm  and dense gas (e.g.  Blain 1999) yielding  the  false impression  that  such  gas  properties (and  the thermalization of high-J CO  transitions) characterize the bulk of the molecular  gas in such  systems.  In  one case  of an  unlensed high-z starburst large scale sub-thermal  excitation of $ ^{12}$CO J=5--4 has been inferred (Papadopoulos \\& Ivison 2002).      Another limitation  of the standard  method is encountered  in tracing metal-poor  H$_2$  gas, which  becomes  more  severe  in enhanced  FUV radiation  fields: the  reduced dust-shielding  and  CO self-shielding allows  UV   photons  to  dissociate  CO  while   leaving  the  mostly self-shielding  H$_2$  intact (Maloney  \\&  Black  1988; Israel  1988, 1997).   This is now  corroborated by  observations of  the metal-poor outer parts of typical spirals (Nakai \\& Kuno 1995; Arimoto, Sofue, \\& Tsujimoto  1996),  globally  metal-poor  objects like  the  Magellanic Clouds and Magellanic Irregulars (e.g.  Madden et al.  1997), and blue compact  dwarf galaxies  (e.g.  Barone  et al.   2000).  Unfortunately both types  of the aforementioned  bias can hinder efforts  to measure H$_2$ gas mass  at high redshifts since high-$J$  $ ^{12}$CO lines are those  used most  often,  and significant  amounts  of metal-poor  gas reside in UV-intense environments (e.g. in Ly-break galaxies).    The CI J=1--0 line was one of the first Photodissociation Region (PDR) lines to be detected from molecular clouds (Phillips \\& Huggins 1981), yet its H$_2$  tracing potential has gone unnoticed.   There are three main reasons for  that namely, a) the low  atmospheric transmission at its rest-frame frequency of  $\\sim 492$ GHz precluded, until recently, routine  observations  of  molecular  clouds in  this  transition,  b) instrumentation constraints  in terms of receiver  sensitivity and the ability to  map large areas fast (no  multi-beam receivers available), and c) an early  theoretical prejudice of one-dimensional steady-state PDR models  that pictured CI  distributed only in a  relatively narrow CII/CI/CO transition zone on  the surface of FUV-illuminated molecular clouds (e.g.  Tielens \\& Hollenbach  1985a,b) and thus unable to trace the bulk  of their  mass.  Interestingly this  prejudice was  at times reinforced by  the instrumentation and  atmospheric restrictions which allowed only limited mapping of  CI emission, which in turn was mainly conducted in regions where theory predicted CI should~be.  These early views  on the  CI  distribution  in PDRs  could  explain some  general features of its  emission such as the surprising  robustness of the CI J=1--0  line brightness  over  a  wide range  of  FUV intensities  and metallicities  (e.g. Kaufman  1999).   This simple picture is now challenged by two large scale CI surveys of the  Orion A  and B  molecular  clouds (Ikeda  et al.   2002) and  the Galactic Center  (Ojha et al.  2001)  that find CI, $  ^{12}$CO, and $ ^{13}$CO J=1--0  emission to be fully  co-extensive, their intensities tightly  correlated, and  with a  surprisingly constant  $\\rm N(CI)/N( ^{12}CO)\\sim 0.1-0.2$ ratio over dramatic changes in ambient FUV field and physical conditions.  These observations along with a considerable body  of similar  past observational  evidence (see  Keene 1997  for a review) point to a CI  distribution fully concomitant with that of CO. The  present  work uses  time-varying  chemical  models under  typical conditions found in Giant Molecular Clouds (GMCs) to further elucidate this issue, and then examines the  potential of CI line emission as an H$_2$ mass tracer.  We focus mainly on the following issues,   \\noindent 1) The  current state  of the  art  of CI  observations revealing  its distribution  throughout molecular clouds  (thus allowing  it to  be a tracer of their mass), and  the role of non-equilibrium and/or dynamic processes as well as cosmic rays under typical ISM conditions.   \\noindent 2) The  expected CI  line  flux  densities over  a  range of  physical conditions and  their relative  observational advantage (if  any) with respect to the $ ^{12}$CO transitions.   \\noindent 3) The uncertainties  of H$_2$  mass estimates from CI  emission with respect  to those associated  with the  standard method  employing CO lines, as well  as the observational prospects of  using the CI lines to trace molecular gas at high redshifts.  \\noindent Throughout  this  work, we  adopt a $\\rm  q_{\\circ  }=0.5$ cosmology with $\\rm \\Omega_{\\Lambda}=0$. ",
        "conclusions": "In  this work  we examined  the prospects  of using  the two  lines of atomic carbon at 492~GHz  (J=1--0) and 809 GHz (J=2--1) as alternative H$_2$ mass tracers in galaxies.   Our conclusions can be summarized as follows:   1. Non-equilibrium  chemical states,  turbulent diffusive  mixing, and cosmic  rays can  easily make  CI ubiquitous  throughout  a typical Giant Molecular Cloud.  This is now firmly supported by large scale CI, $ ^{12}$CO, $ ^{13}$CO surveys of such clouds in the Milky Way, which also find a constant average [C]/[$ ^{12}$CO] abundance ratio over the bulk of their mass.   2. In star  forming and/or dynamically  fast-evolving environments the processes responsible for  the ubiquity of CI in  GMCs will further enhance its abundance.  In  UV-intense and metal-poor objects large scale CII production diminishes the H$_2$-tracing capability of CI, which  nevertheless remains better  than that  of the  now severely dissociated $ ^{12}$CO.   3. In  terms of  emerging flux  density  as a  function of  excitation conditions CI lines have a potentially strong advantage with respect to $ ^{12}$CO  lines in tracing H$_2$ gas  mass under diffuse and/or cooler conditions as well as  in the metal poor environments of mild starbursts.   However  the  large  atmospheric absorption  at  their rest-frame frequencies  allows such  an advantage to  materialize at high redshifts ($\\rm z\\ga 1$).   4. The simpler 3-level structure of  the CI lines, their small optical depths  and  modest excitation  requirements  with  respect to  $ ^{12}$CO lines of similar frequency make the associated H$_2$ mass estimates  more  straightforward.   However  the  usually  unknown (particularly   for  high   redshift   sources)  mean   excitation conditions introduce large  uncertainties in such estimates, which are reduced considerably if an  estimate of the gas temperature is available.   5. Concentrated  observational effort  in the  local  Universe remains critical.  Large  scale CI, $  ^{12}$CO, and $ ^{13}$CO  surveys of GMCs  in a  variety of  environments  will help  address issues  of optical depths, $\\rm [CI]/[H_2] $ abundances and its variations, and ``cross-calibrate'' the two H$_2$ gas mass tracing techniques."
    },
    "0403/astro-ph0403571_arXiv.txt": {
        "abstract": "We study a cosmological model in which the boson dark matter gradually condensates into dark energy. Negative pressure associated with the condensate yields the accelerated expansion of the Universe and the rapid collapse of the smallest scale fluctuations into many black holes, which become the seeds of the first galaxies. The cycle of gradual sedimentation and rapid collapse of condensate repeats many times and self-regularizes the ratio of dark energy and dark matter to be order one. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/hep-th0403236_arXiv.txt": {
        "abstract": "Due to the tremendous red-shift that occurs during the inflationary epoch in the early universe, it has been realized that trans-Planckian physics may manifest itself at energies much lower than the Planck energy.  The presence of a fundamental scale suggests that local Lorentz invariance may be violated at sufficiently high energies. Motivated by this possibility, recently, different models that violate Lorentz invariance locally have been used to evaluate the trans-Planckian corrections to the inflationary density perturbation spectrum.  However, certain astrophysical observations seem to indicate that local Lorentz invariance may be preserved to extremely high energies.  In such a situation, to study the trans-Planckian effects, it becomes imperative to consider models that {\\it preserve}\\/ local Lorentz invariance {\\it even as they contain a fundamental scale}.\\/ In this work, we construct one such model and evaluate the resulting spectrum of density perturbations in the power-law inflationary scenario.  While our model reproduces the standard spectrum on small scales, it {\\it naturally}\\/ predicts a suppression of power on large scales.  In fact, the spectrum we obtain has some features which are similar to the one that has recently been obtained from non-commutative inflation.  However, we find that the amount of suppression predicted by our model is {\\it far less}\\/ than that is required to fit the observations.  We comment on the fact that, with a suitable choice of initial conditions, our approach can lead to corrections at the infra-red {\\it as well as}\\/ at the ultra-violet ends of the spectrum. ",
        "introduction": "Inflation, a period of accelerated expansion in the high-energy phase of the universe, is currently considered to be the best paradigm for describing the early stages of the universe~\\cite{txts,infltnmdls}. The success of the inflationary paradigm rests on its ability to explain not only the homogeneity of the background, but also the characteristics of the inhomogeneities superimposed upon it.  The inflationary epoch magnifies the tiny fluctuations in the quantum fields present at the beginning of the epoch into classical perturbations that leave an imprint as anisotropies in the cosmic microwave background (CMB). These anisotropies in turn act as seeds for the formation of the large-scale structure that we observe at the present time as galaxies and clusters of galaxies.  With anisotropies in the CMB being measured with higher and higher precision, we are currently able to test the predictions of inflation better and better.  During the last couple of years or so, considerable amount of attention has been devoted to examining the possibility that Planck scale physics may leave a measurable imprint on the CMB~\\cite{mb01}--\\cite{mb03}.  There are three reasons that have led to such an enormous interest in the literature.  Firstly, there is no unique model of inflation~\\cite{infltnmdls}.  In many versions of inflation, most notably in chaotic inflation, the period of inflation lasts sufficiently long so that the comoving length scales that are of cosmological interest today would have emerged from sub-Planckian length scales at the beginning of inflation.  Hence, in principle, quantum gravitational effects should have left their signatures on the primordial spectrum of perturbations.  Secondly, in the simplest models of inflation, the scale of the vacuum energy during the period of exponential expansion is assumed to be $\\sim 10^{16}\\, {\\rm GeV}$, while the rate of the exponential expansion, viz. $H$, is considered to be $\\sim 10^{14}\\, {\\rm GeV}$.  These enormous energies suggest that during the inflationary epoch, various (as yet, unknown?) high-energy processes could have been activated which would have left their signatures on the primordial perturbation spectrum.  These signatures on the perturbation spectrum in turn will leave their imprints on the CMB.  Recent measurements of the CMB anisotropies by, say, WMAP~\\cite{MAP}, strongly indicate a primordial spectrum that is nearly scale-invariant, just as the inflationary scenario predicts. Therefore, further precise and accurate measurements of the CMB anisotropies by experiments such as PLANCK~\\cite{PLANCK} can, in principle, provide us with the form of the corrections to the scale-invariant primordial perturbation spectrum.  Lastly and, more importantly, the first results of the temperature-temperature correlation spectrum of the WMAP data show that the power in the quadrupole and (to a lesser extent) the octopole moment of the CMB are {\\it lower than}\\/ as expected in the best fitting cold dark matter models~\\cite{MAP}.  The deficit of power in the quadrapole moment of the CMB power spectrum can not be explained within the context of the standard inflationary models (unless these models are fine-tuned \\cite{contaldi}) and suggests a possible signature of the Planck scale physics~\\cite{nci}.  Several groups have pointed out to these possibilities and, broadly, there have been two approaches in the literature in order to study these effects. In the first approach, which is now commonly referred to as the minimal trans-Planckian approach, the specific nature of trans-Planckian physics is not assumed, but is rather described by the boundary conditions imposed on the mode at the cut-off scale~\\cite{minimaltp,act03}. In the second approach, one incorporates quantum gravitational effects by introducing the fundamental length scale into the standard field theory in a particular fashion and, it should be emphasized here that, the resulting modified theory may not even preserve local Lorentz invariance.  Indeed, most of the attempts in this direction have involved models which {\\it break}\\/ local Lorentz invariance.  Such models include those which introduce non-linear dispersion relations~\\cite{mb01,niemeyer,earliertp,br,shanki03} as well as the approaches which utilize the generalized uncertainty principle~\\cite{guptp} and non-commutative geometry~\\cite{nctp,nci}. (For a recent review on these various approaches, see Ref.~\\cite{mb03}.)  However, theoretically, there exists no apriori reason to believe that Lorentz invariance may be broken at the scales of inflation. More importantly, recent observations of photons inferred to arise from synchrotron emission off electrons in the Crab nebula have placed bounds on dispersive corrections involving additional powers of the particle momentum to be at or above the Planck scale~\\cite{liv}. In such a situation, in order to study the trans-Planckian effects on the primordial perturbation spectrum, it becomes important that we also consider models which preserve Lorentz invariance {\\it even as they contain a fundamental scale}.\\/ In this work, we consider one such model and evaluate the resulting spectrum of density perturbations during inflation in this model.  The model we shall consider is as follows: We shall {\\it assume}\\/ that, due to the Planck-scale effects, the standard $k-$space propagator for a massless scalar field in the Minkowski vacuum is modified in a particular manner.  We shall introduce the high-energy scale $k_{\\rm c}$ into the $k-$space propagator in a Lorentz invariant manner and, as we shall point out later, the modification we consider can be said to be minimal in nature.  We find that the resulting modified Wightman function can be expressed as the difference of the Wightman functions of the massless field and a massive field of mass $k_{\\rm c}$.  We shall further assume that the form of the modified Wightman function in the $x-$space remains the same in an inflationary background as well, i.e. it is the difference of the massless and the massive Wightman functions.  We then use this modified Wightman function to evaluate the corrections to the spectrum of density perturbations for the power-law inflationary scenario.  We show that our model naturally predicts a suppression of power at the large scales.  We find that the modified spectrum we obtain has some resemblance to the spectrum that has been obtained in the non-commutative inflationary scenario~\\cite{nci}.  However, the amount of suppression predicted by our model turns out to be {\\it far less}\\/ than that seems to be required to fit the WMAP data.  The rest of the paper is organized as follows.  In Section~(II), we outline the model we shall consider, present the motivations behind the model and also point out its attractive features.  In Section~(III), we shall briefly review the spectrum of fluctuations in the standard power-law inflationary scenario.  In Section~(IV), we shall evaluate the corrections to the spectrum of fluctuations using our model for the power-law inflation.  Finally, in Section~(V), we shall summarize the results of our analysis and also discuss its implications.  Before we proceed, a few comments on the notations we shall use are in order.  The metric signature we shall adopt is $(+, -,-, -)$, we shall set $\\hbar=c=1$ and we shall denote the set of four coordinates $x^{\\mu}$ simply as ${\\tilde x}$.  Also, the quantum field $\\Psi$ we shall consider will be a minimally coupled scalar field. ",
        "conclusions": "In this work, we have studied the trans-Planckian effects on the spectrum of the primordial density perturbations in the power-law inflationary scenario using an approach that preserves local Lorentz invariance.  Motivated by the fact that quantum gravitational effects can be expected to improve the divergence structure of standard quantum field theory, we {\\it assumed}\\/ that the trans-Planckian physics modifies the standard massless propagator such that the modified propagator can be expressed as the {\\it difference}\\/ of the original massless propagator and a massive propagator with a mass of the order of the Planck mass.  In the standard inflationary scenario, the primordial perturbation spectrum is determined by the Fourier transform of the massless scalar field propagator.  Therefore, modifying the scalar field propagator leads to modifications in the perturbation spectrum.  We find that, in our model, the resulting modified spectrum remains scale invariant at the ultra-violet end, but, interestingly, it exhibits a suppression of power at the infra-red end---a feature that seems to be necessary to explain the low quadrupole and octopole moments measured in the CMB~\\cite{MAP,contaldi,blwe03,tarun}.  However, at Hubble exit, the amount of suppression predicted by our model in power-law inflation turns out to be {\\it far less}\\/ than as expected from WMAP data. Nevertheless, the loss of power at small $k$ suggests that the power spectrum we have obtained may fit the WMAP data better than the standard $\\Lambda$CDM model.  It will be interesting to analyze the implications of the WMAP data for our model in the context of slow-roll inflation~\\cite{ss04}.  Naively, one would expect that very high energy effects will leave their imprints only at the ultra-violet end of the primordial spectrum.  However, we find that the high energy effects lead to a modification of the spectrum at the infra-red end.  As we had pointed out in the last section, this can be attributed to the fact that the longer wavelength modes leave the Hubble radius at earlier epochs thereby carrying the signatures of the high energy effects.  Finally, we would like to emphasize a very attractive feature of our approach which can result in corrections at the infra-red {\\it and}\\/ the ultra-violet ends of the perturbation spectrum.  Recall that in obtaining the modified spectrum (\\ref{eq:psmfr}) we had assumed that both the massless and the massive fields are in the Bunch-Davies vacuum.  This need not be the case.  While it is natural to assume that the massive field is in the Bunch-Davies vacuum (in particular, a very heavy field with a mass of the order of the Planck mass), one can assume the massless field to be in a state such as the minimal uncertainty state~\\cite{minimaltp}.  In such a case, the resulting modified spectrum will have corrections at both the ends.  While at the infra-red end, we will simply reproduce the corrections we have obtained, in the ultra-violet limit, the corrections will be of the form that has been recently obtained in the literature~\\cite{act03}."
    },
    "0403/hep-ph0403272_arXiv.txt": {
        "abstract": "We find the condition that the primordial baryon asymmetry is not washed out in the $SU(2)_{L}$ triplet Higgs model by solving the Boltzmann equation. We further require that the model is compatible with the recent results of the neutrino oscillation experiments and the Wilkinson Microwave Anisotropy Probe, and the constraints on the $\\rho$ parameter imposed by the CERN LEP. We finally obtain the allowed region of the parameters in the model. ",
        "introduction": "} The executed experiments of the neutrino oscillation have made it certain that the neutrinos have the very small, but finite masses \\cite{S-K, S-K2, Valle-Gon, SNO, SNO2, Chooz}. The recent KamLAND experiment gave us the more precise information about the neutrino mass matrix \\cite{kam, kam2,Valle-kam}. Although the neutrino oscillation experiments are not able to determine the magnitudes themselves of the neutrino masses, the recent results of Wilkinson Microwave Anisotropy Probe (WMAP) imposed the upper bound of the sum of the neutrino masses \\cite{Wmap1,Wmap2,Han,Elg,Wmap3}. Since the neutrinos have no masses in the standard model, it is an urgent issue that we extend the standard model and build the model where the neutrinos naturally have the very small masses in comparison with the other leptons and quarks. It is expected that the smallness of the neutrino masses is explained by the Majorana type mass, which only neutrinos among quarks and leptons are able to have. The Majorana mass terms have the noticeable feature that they violate the lepton number,  which is conserved at the tree level in the standard model. A few types of models of the Majorana neutrinos are proposed. The seesaw model \\cite{Yan,Yan2,Yan3}, the $SU(2)_{L}$ triplet Higgs model \\cite{Sch,Laz1}, and the Zee model \\cite{aZee} are the three typical models where the standard model is minimally extended to induce the Majorana masses of the neutrinos within the framework of the $SU(2)_{L} \\otimes U(1)_{Y}$ gauge theory.   In this paper we focus on the $SU(2)_{L}$ triplet Higgs model. Especially, we introduce the model which is built in the framework of the large extra dimension and can naturally explain the smallness of the neutrino masses by the \"shining\" mechanism \\cite{Ema1, Ema2}. We live in one brane and the lepton number is violated in the distant other brane. A messenger field informs us of the lepton number violation in the other brane. Since the messenger field spreads out in the extra dimensions, the density of the field is diluted, which is called the \"shining\" mechanism. Since we are informed of the only tiny lepton number violation, the very small neutrino masses are naturally induced in our brane. Furthermore, based on the 't Hooft's naturalness condition, the tiny lepton number violation, that is, the very small neutrino mass, is stabilized under the quantum corrections. In this paper we adopt the 4-dimensional effective theory in our brane of the $SU(2)_{L}$ triplet Higgs model in the large extra dimension. It is also noticed that the $SU(2)_{L}$ triplet Higgs model must not induce the dangerous pseudo Majoron which is inconsistent with the accelerator experiments as the CERN LEP, and the model is constrained from the observation of the $\\rho$ parameter in the LEP \\cite{gri, erl}. The lepton flavor violation in this model is analyzed in \\cite{Chu2, Kak}.    The baryon number in the present universe is well known to be \\cite{Wmap1} \\begin{eqnarray} B_{0}\\equiv \\frac{n_{B}}{s}=(8.4 \\sim 9.1) \\  \\times 10^{-11}, \\end{eqnarray} where $n_{B}$ is the baryon number density and $s$ is the entropy density. Some scenarios where the baryon number is dynamically generated from vanishing initial baryon number, $B_{ini}=0$, are proposed. The GUT scenario \\cite{Yos, Barg}, the electroweak scenario \\cite{Kuz, Coh}, the Fukugita-Yanagida scenario \\cite{Fuk}, and the Affleck-Dine scenario \\cite{Aff} are proposed as the typical scenarios of the baryogenesis. The GUT and electroweak scenarios have already been almost ruled out. It is often mentioned that these scenarios of the baryogenesis have a double-edged behaver, that is, the scenarios of the baryogenesis simultaneously have a potential to erase the existing baryon number. Especially, while the mass models of the Majorana neutrinos have a potential to generate the baryon number through the leptogenesis as in the Fukugita-Yanagida scenario, these models simultaneously have a potential to erase arbitrary initial lepton and therefore baryon numbers. There exist the sphaleron-anomaly processes which create and annihilate the baryon in an equal amount as that of the lepton in the standard model \\cite{Hoo, Man, Kli, Arn, Bod}. When the the sphaleron-anomaly processes and the lepton number violating processes in the mass models of the Majorana neutrinos are simultaneously in equilibrium, the arbitrary initial baryon and lepton numbers are actually washed out. Since the sphaleron-anomaly processes are in equilibrium at the large region of temperature between 100 GeV and $10^{12}$ GeV, whether the lepton number violating processes are in equilibrium or not in the region handles the fate of the cosmological baryon number.   In this paper, we assume that the primordial baryon number is generated by some unknown mechanism at the very early universe and require that the primordial baryon number is not washed out in the mass models of the Majorana neutrinos where the baryon number is never generated. The situation where the baryon number is never generated via leptogenesis in the Majorana mass models is typically realized when we assume that the lepton sector has no CP violation because of the Sakharov's three conditions \\cite{Sak}. Although we assume no CP violation in the following analyses, the results of this paper are valid as long as the mass models of the Majorana neutrinos are not able to generate the present baryon number. Actually, it is shown that even if the $SU(2)_{L}$ triplet Higgs model, which we use in this paper, has CP violation in the lepton sector, the baryon number can not be generated \\cite{Ern}. The leptogenesis in the other extended triplet Higgs models is analyzed in \\cite{Laz2,Laz3}. In  order to generate the baryon number via leptogenesis, two triplet Higgs fields must be introduced. The similar analyses are carried out in the GUT scenario \\cite{Fry1, Fry2} and in the Fukugita-Yanagida scenario \\cite{Fuk2, Hav, Nel, Buc3,  Fis, Oli1, Oli2, Dre}.   In order to escape from washing out of the baryon number in the mass models of the Majorana neutrinos, it is necessary that the lepton number violating processes are out of equilibrium. An out-of-equilibrium process means that the reaction rate of a process is smaller than the Hubble parameter at that time. In order to violate the lepton number explicitly in the $SU(2)_{L}$ triplet Higgs model, two kinds of interactions must be newly introduced. At first sight, it is not clear whether both of the two kinds of interactions must be out of equilibrium or either of the two kinds of interactions must be out of equilibrium in order to protect the primordial baryon number. The answer we obtain in this paper is that if either of the two kinds of interactions is out of equilibrium, the primordial baryon number is not washed out. If either of the two kinds of interactions is out of equilibrium, the lepton number is considered not to be explicitly broken and an approximately preserved global $U(1)$ charge which contains the lepton number exists. Thus the primordial baryon number is protected in proportion to the initial value of the approximate $U(1)$ charge. It is the first purpose in this paper that we can concretely estimate the condition to protect the primordial baryon number by solving the Boltzmann equation in the $SU(2)_{L}$ triplet Higgs model. We next analyze whether the $SU(2)_{L}$ triplet Higgs model can simultaneously satisfy the obtained condition to protect baryon number and the results of the neutrino oscillation experiments or not. We further require that the model satisfies the results of WMAP and the constraints on the $\\rho$ parameter. It is the second purpose in this paper that we require that the model satisfies all the  above conditions and obtain the allowed region of the parameters in the model. The similar analyses are carried out in the seesaw model \\cite{Has2} and in the Zee model \\cite{Haba, Has1}.   The outline of this paper is as follows: in Sec. \\ref{trimodel}, we review the $SU(2)_{L}$ triplet Higgs model where the smallness of the neutrino masses is naturally explained. We make sure that the dangerous pseudo Majoron which is inconsistent with the accelerator experiments as LEP does not appear in the model. In Sec. \\ref{Washing}, we obtain the condition to protect the primordial baryon number by solving the Boltzmann equation in the $SU(2)_{L}$ triplet Higgs model. In Sec. \\ref{triprotect}, we require that the $SU(2)_{L}$ triplet Higgs model satisfies the condition to protect the baryon number, the results of the neutrino oscillation experiments and WMAP, and the constraints on the $\\rho$ parameter. We finally obtain the allowed region of the parameters in the model. Sec. \\ref{trisum} is devoted to a summary. ",
        "conclusions": "} \\subsection{Conclusions} We introduce the $SU(2)_{L}$ triplet Higgs model where the smallness of the neutrino masses is naturally explained by the shining mechanism in the large extra dimension. In order to violate the lepton number explicitly, which is the source of the Majorana mass terms, we need to introduce the two kind of the interactions in addition to the ones in the standard model. It is a first fruitful consequence that we can obtain the condition (\\ref{tencon}) that the primordial baryon number is not washed out by solving the Boltzmann equations in the model. We find that in order to protect the primordial baryon number, either of the two kind of the interactions, both of which are needed to violate the lepton number explicitly, must be out of equilibrium under the in-equilibrium sphaleron-anomaly processes. The condition to protect the primordial baryon number is easily extended into the case that the differences of the three lepton flavors are taken into account. It is a second fruitful consequence that we can obtain the allowed region where the condition to protect the primordial baryon number, the results of the neutrino oscillation experiments and WMAP, and the constraints on the $\\rho$ parameter are all satisfied in the model. A typical allowed region is shown in Fig. \\ref{asfig}. We find that if both of the coupling constant $A$ and $f^{\\alpha \\beta}$ are large, the existing baryon number is washed out and when both of $A$ and $f^{\\alpha \\beta}$ are small, the mass differences which are induced from the neutrino mass matrix are too small  to explain the observed neutrino oscillations. In this model, the constraints on the $\\rho$ parameter give the upper bounds on the coupling constant $A$ for the fixed triplet Higgs masses. We find the two allowed regions where all the requirements are satisfied. One region where $A$ is very small and $f^{\\alpha \\beta}$ are as large as the Yukawa coupling constants of the charged leptons is expected to be a favorable region where the smallness of the neutrino masses is naturally explained in the $SU(2)_{L}$ triplet Higgs model in the large extra dimension.  \\subsection{Discussions} We first make comment on the pre-existing approximately conserved charge. We actually obtain the approximately conserved charges, $P=B-L+2\\Delta$ or $P=B-L$. The primordial baryon number is protected in proportion to the initial value of the conserved charge. Can we naturally expect that the initial amount of the conserved charge exists in the early universe ? We recall that in the $SU(5)$ GUT, the initial value of the conserved charge $(B-L)$ is naturally assumed to be zero, which regrettably leads to the washing out of the baryon number by the in-equilibrium sphaleron processes. The difference between the case we examine and the case of $SU(5)$ GUT is that in the former case the conserved charge is approximate one, while in the latter case the conserved charge is exact one. So there is no reason why an approximately conserved charge is not generated by the unknown mechanism at the very early universe.   We finally enumerate the remaining problems and the analyses which have not been carried out yet in this paper. The first problem is that we have omitted all four body processes such as $ll \\leftrightarrow ll$ or $ll \\leftrightarrow \\phi \\phi$, where $l$ is the lepton and $\\phi$ is the Higgs field. Although we expect that the rates of the  four body processes are much smaller than those of the three body processes, the  four body processes should be taken into account for the exactness. The second problem is that the analyses containing CP violation are not carried out. We have dealt in the restricted parameter region where there are no CP violation in the lepton sector of the models in order to focus on the feature that the baryon number is washed out in the Majorana neutrino mass models. If we try to investigate the feature more generally and more exactly, the analyses containing the CP violation should be carried out. However, even if the $SU(2)_{L}$ triplet Higgs model which has the only one triplet Higgs field contains CP violation in the lepton sector, the baryon number can not be generated. Hence the condition to protect the primordial baryon number which we obtain in this paper is valid in the case that the model contains the CP violation. On the other hand, since Majorana neutrino mass matrix generally has one Kobayashi-Maskawa phase and two Majorana phases, the analyses of the results of the neutrino oscillation experiments might be slightly changed."
    },
    "0403/astro-ph0403021_arXiv.txt": {
        "abstract": "{We study the conditions for convective instability in rotating, non-magnetic proto--neutron stars. The criteria that determine stability of nascent neutron stars are analogous to the Solberg--H{\\o}iland conditions but including the presence of lepton gradients. Our results show that, for standard angular velocity profiles, convectively unstable modes with wave-vectors parallel to the rotation axis are suppressed by a stable angular momentum profile, while unstable modes with wave-vectors perpendicular to the axis remain unaltered. Since the wave-vector is perpendicular to the velocity perturbation, the directional selection of the unstable modes may result in fluid motions along the direction of the rotation axis. This occurs in  rigidly rotating stars as well as in the inner core of differentially rotating stars. Our results provide a natural source of asymmetry for proto--neutron stars with the only requirement that angular velocities be of the order of the convective characteristic frequency. ",
        "introduction": "Hydrodynamical instabilities in newly born neutron stars play an important role in enhancing neutrino luminosities and increasing the energy deposition efficiency. It was first argued by Epstein (1979) that the negative lepton gradient that arises in the outer layers of the proto-neutron star (PNS) can be convectively unstable. However, convection in PNSs can be driven not only by the lepton gradient but by a negative entropy gradient as well (see, e.g., Arnett 1987, Burrows \\& Lattimer 1988). The development of negative entropy and lepton gradients is common in many simulations of supernova models (Bruenn \\& Mezzacappa 1994, Bruenn et al. 1995) and evolutionary models of PNSs (Burrows \\& Lattimer 1986, Keil \\& Janka 1995, Sumiyoshi et al. 1995, Pons et al. 1999). The nature of instabilities in PNSs has been considered by a number of authors (Grossman et al. 1993; Bruenn \\& Dineva 1996; Miralles et al.  \\cite{MPU00}, MPU in the following; Miralles et al.  \\cite{MPU02}). In non-rotating and non-magnetic PNSs, the criteria indicate the presence of two essentially different instabilities (MPU) with a convectively unstable region surrounded by a neutron-finger unstable region, the latter involving typically a larger portion of the stellar material. The unstable zones grow in size and move inward as the PNS deleptonizes (Keil et al. 1996), until neutrino diffusion progressively reduces the temperature and lepton gradients, and convective instabilities disappear, at most in 30-40 s. The growth time of convective and neutron finger instabilities are substantially different, as well as the efficiency of turbulent transport in the convective and neutron-finger unstable zones.  Less is known about convection in rotating PNSs. The failure of detailed simulations in spherical symmetry to get a successful Supernova explosion has led different authors to argue that multidimensional effects, such as convection or  rapid rotation can be the key to the problem (see e.g. Buras et al. 2003, and references therein). From theoretical modeling and simple analytic considerations it is commonly accepted that core collapse of a rotating progenitor leads to differential rotation of a newly born neutron star (Zwerger \\& M\\\"uller 1997, Rampp et al. 1998, Liu 2002, Dimmelmeier et al. 2002, M\\\"uller et al. 2003), mainly due to conservation of the angular momentum during collapse. A recent study on evolutionary sequences of rotating PNSs (Villain et al. \\cite{Loic03}) shows that the typical scale on which the angular velocity changes is in the range $\\approx 5-10$ km. Granted that a PNSs is born rapidly rotating, when the angular velocity is of the order or larger than the Brunt-V\\\"ais\\\"al\\\"a or Ledoux frequencies, even relatively fast convective instabilities can be substantially modified, and convection can be constrained to the polar regions (Fryer \\& Heger \\cite{FH00}). In this paper, we study this effect in the context of PNSs. ",
        "conclusions": "We have derived the criteria for convective instability in rotating, non--magnetic PNSs, obtaining two conditions analogous to the Solberg--H{\\o}iland criteria but including the effects of lepton gradients. In convectively unstable regions, the existence of a positive gradient of angular momentum can stabilize the unstable convective modes in some directions. The basic picture is that an unstable buoyant element of fluid is stabilized when it tries to move toward a direction in which the surrounding fluid elements have larger angular momentum, because it is subject to a smaller centrifugal force and therefore restored to its equilibrium position. However, this restoring force has no effect on buoyant waves that follow lines of constant angular momentum. For example, in a situation in which thermal, lepton, and pressure gradients are nearly aligned, the angular momentum depends weakly on the vertical coordinate, thus allowing convectively unstable modes to grow only in the direction approximately parallel to the rotation axis. In a general case, the geometry will be more complicated but the qualitative results will not change. In Fig. \\ref{fig4} we illustrate this effect in a realistic model of a PNS 0.5 s after birth, taken from Pons et al. (1999). The convectively unstable region extends from 17 to 28 km, and the arrows indicate the direction of the most rapidly growing mode, with the length being proportional to the inverse growth time. In the convectively unstable region near the equator, $\\Omega_e$ varies between 500 and 1000 rad/s, which is about one third the value of $\\omega_0$, and we have taken $R_0=10$ km (defined in Eq. (\\ref{slaw})). In a saturated state, hydrodynamical non-linear effects will probably smear out anisotropies on small scales but it is unlikely that this may occur on the largest convective scale.  Hydrodynamic simulations of the early evolution of rotating PNSs show that the scenario drawn in Fig. \\ref{fig4} is fulfilled during the very early phase of convection in PNSs (H-Th. Janka, private communication), which shows that our linear stability analysis reflects accurately the situation at the onset of convection. However, after several hundreds of milliseconds, the efficient neutrino and convective transport in the polar region will flatten the temperature, lepton and angular momentum profiles, resulting in a different situation: convection occurring only close to the equator, although convective shells are still elongated and parallel to the rotation axis (see e.g. Fig. 10 in Janka \\& Keil \\cite{JK97}, or Janka, Kifonidis \\& Rampp \\cite{JKR01} ).  This coupling between convective instability and stabilizing gradients of angular momentum can be another ingredient that helps to explain the large asymmetry that has been observed in recent simulations of rotational core collapse Supernovae (M\\\"uller et al., 2003), not observed in simulations of convective PNSs without rotation.  The mechanism we propose is a natural way to create anisotropic energy and momentum transport by convective motions, that only requires that the angular velocity be of the order of the Brunt--V\\\"ais\\\"al\\\"a or leptonic frequencies. Notice that our results are relevant for convective motions in the PNS, more than in the extended envelope as in Fryer \\& Heger (\\cite{FH00}). Our analysis is only valid in a steady--state situation, while the conditions in the accretion layer are extremely non--stationary, and other effects such as anisotropic accretion are certainly more relevant. If the collapsed core shows strong differential rotation, at scales larger than 50-100 km the angular momentum can be nearly constant and rotational effects are small. Nevertheless, the anisotropy originated in the inner core may have some effect at larger scales, and can help to explain the Supernova mechanism. Obviously, self--consistent 3D simulations involving convection, neutrino transport and angular momentum transport are needed prior to drawing robust conclusions.  In our study we only considered non--magnetic proto--neutron stars. The presence of a large magnetic field could change this picture.  It has recently been shown that turbulent mean--field dynamo can be effective for PNSs with periods shorter than 1 s, which would generate strong magnetic fields in the interior (Bonanno et al. \\cite{BRU}). However, magnetic field effects begin to be relevant only when the Alfven frequency is of the order of the other dynamical or dissipative frequencies, and this requires magnetic fields larger than 10$^{15}$ G (Miralles et al. \\cite{MPU02}). Probably the missing ingredient to explain Supernovae is related to the coupling between different effects, such as the one proposed in this paper, the magneto--rotational instability (Akiyama et al., \\cite{Aki03}), or a combination of all of them.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{0924fig4.ps}} \\caption{The direction of the motion for the most rapidly growing modes in a realistic model of PNS 0.5 s after birth. In the convective region (17-28 km), the instability is more effective in the polar region and, even at intermediate latitudes, vertical motions are preferred. } \\label{fig4} \\end{figure}"
    },
    "0403/astro-ph0403217_arXiv.txt": {
        "abstract": "We have made the first unambiguous detection of vibrational ground-state maser emission from $^{28}$SiO toward six evolved stars.  Using the Very Large Array, we simultaneously observed the $v=0, J=1-0$, 43.4-GHz, ground-state and the $v=1, J=1-0$, 43.1-GHz, first excited-state transitions of $^{28}$SiO toward the oxygen-rich evolved stars IRC+10011, $o$~Ceti, W Hya, RX Boo, NML Cyg, and R Cas and the S-type star $\\chi$ Cyg. We detected at least one $v=0$ SiO maser feature from six of the seven stars observed, with peak maser brightness temperatures ranging from 10000~K to 108800~K.  In fact, four of the seven $v=0$ spectra show multiple maser peaks, a phenomenon which has not been previously observed. Ground-state thermal emission was detected for one of the stars, RX Boo, with a peak brightness temperature of 200~K.  Comparing the $v=0$ and the $v=1$ transitions, we find that the ground-state masers are much weaker with spectral characteristics different from those of the first excited-state masers.  For four of the seven stars the velocity dispersion is smaller for the $v=0$ emission than for the $v=1$ emission, for one star the dispersions are roughly equivalent, and for two stars (one of which is RX Boo) the velocity spread of the $v=0$ emission is larger.  In most cases, the peak flux density in the $v=0$ emission spectrum does not coincide with the $v=1$ maser peak.   Although the angular resolution of these VLA observations were insufficient to completely resolve the spatial structure of the SiO emission, the SiO spot maps produced from the interferometric image cubes suggest that the $v=0$ masers are more extended than their $v=1$ counterparts. ",
        "introduction": "Although the study of SiO maser emission toward long-period variable stars is an active field of research, very little work has been done with regards to the vibrational ground-state ($v=0$) SiO emission toward these evolved objects.  The question of the existence of observable ground-state SiO masers was first posed by \\cite{BSLJ:75}. The ground-state lines of $^{28}$SiO and the much rarer $^{29}$SiO and $^{30}$SiO isotopic species have been observed in a number of single-dish surveys conducted primarily at 43 GHz ($v=0, J=1-0$) and 86 GHz ($v=0, J=2-1$).  In the two isotopic species, this ground-state emission has been shown to be relatively weak unsaturated maser emission \\citep{AB:92}.  For $^{28}$SiO the nature of the $v=0$ emission is not as clear.  Although \\cite{BSLJ:75} may have detected a weak $v=0, J=2-1$ maser line toward the star VY CMa, the general consensus on the 86.8 GHz line is that it is thermal circumstellar SiO emission with broad parabolic or flat-topped line profiles \\citep[see for example][]{BPGMD:86}. The $v=0,J=1-0$, 43.4-GHz line was similarly thought to be thermal in nature, however in surveys by \\cite{JWWS:85} and \\cite{JSWWG:91} several of the $v=0$ line profiles showed narrow emission spikes ($<1$~km~s$^{-1}$) superposed, in some cases, on broader pedestals of emission.  \\citeauthor{JSWWG:91} suggested that these narrow spikes could be maser emission with linewidths comparable to or narrower than the thermal Doppler width. The data of \\cite{JSWWG:91} also suggested a correlation between the intensities of the $v=0$ and $v=1$ lines providing additional evidence for ground-state maser emission and the possibility that the gas producing the two types of emission is spatially coincident.  To date, very little interferometric imaging of the $v=0$ emission toward SiO maser emitting regions has been conducted.  The more abundant $^{28}$SiO species was observed toward the Orion IRc2 star-forming region in the $J=1-0$ transition \\citep{CD:95} and the $J=2-1$ transition \\citep{WPML:95}. In each case, the detected emission was found to have both maser and thermal components.  Observations of the rarer $^{29}$SiO $v=0, J=2-1$ transition toward Orion \\citep{BHL:98} also showed this blend of thermal and maser emission. A comparison of the $^{28}$SiO $v=1, J=2-1$ and $^{29}$SiO $v=0, J=2-1$ maser spots maps lead \\cite{BHL:98} to conclude that the two isotopic species are not excited in the same gas layers.  Toward late-type stars, interferometric observations of the $v=0, J=2-1$ $^{28}$SiO emission have been conducted by \\cite{LUCAS:92} and by \\cite{SB:93}. No evidence of ground-state masers was found in either observation with the possible exception being an intense spike of emission toward the star R Cas \\cite{LUCAS:92}. Prior to the work presented here, no interferometric studies of the $v=0, J=1-0$ $^{28}$SiO transition had been conducted.  We therefore sought to detect and image the ground-state SiO maser emission toward 7 stars from the list of possible detections in \\cite{JSWWG:91}. ",
        "conclusions": "Ground-state, $^{28}$SiO maser emission is rarely addressed in the literature in which numerical models for the generation of SiO masers are presented. This may be due to the fact that to date there has been very little evidence for the existence of masers in this transition. In an early paper, \\cite{KS:74} commented that population inversions in the $v=0$ states are possible under certain conditions.  They stated that the percentage of the inversion is small due to the fact that the collision rate among the $v=0$ levels is comparable to the rate of radiative depopulation. In one model consisting of strictly radiative excitation,\\cite{KS:74} found an inversion of the $v=0, J=1-0$ line which occurs at a distance of $5\\times 10^{14} - 4\\times 10^{15}$~cm ($\\sim$9--69 $R_{\\star}$) with an emission rate down by a factor of 100 from the $v=1, J=1-0$ line.  More recent investigations addressing the excitation of the SiO masers through overlaps with optical \\citep{RKTP:96} and infrared \\citep{HB:00} lines, demonstrated possible mechanisms for the generation of $v=0$, $^{28}$SiO masers. In the case of optical pumping, \\cite{RKTP:96} found the ground-state emission to be thermal for the $v=0, J=2-1$ and $v-0, J=3-2$ transitions of $^{28}$SiO, but that maser emission was generated for the $v=0, J=1-0$ state in their simulations. The infrared pumping mechanism of \\cite{HB:00} was used to resolve some of the previously unexplained detections of certain SiO maser transitions and to predict some possible new SiO masers.  In addition, their numerical model also demonstrated that $v=0, J=1-0$ $^{28}$SiO could be generated.  In both the optical and infrared pumping cases, the $v=0$ masers are generated in regions of the CSE where the SiO abundances are lower than for the typical excited-state masers.  Since the SiO abundance decreases with distance from the star as SiO condenses onto dust grains, then the $v=0$ masers are naturally formed at distances beyond that of the vibrationally excited masers in these models.  Our qualitative comparison of the spatial distributions supports the idea that the $v=0$ SiO masers occur at distances greater than $v=1$ maser shell.  In this case the $v=0$ maser distribution is several times the diameter of the corresponding $v=1$ masers.   This contradicts the speculation by \\cite{JSWWG:91} that the correlation between the intensities of the $v=0$ and $v=1$ lines suggests that emission from the two transitions is spatially coincident.  \\cite{BHL:98}, in their interferometric observations of the $v=0, J=2-1$ $^{29}$SiO and the $v=1, J=2-1$ $^{28}$SiO masers toward Orion IRc2, determined that the ground-state and excited-state masers were not co-spatial.  They state that this is a consequence of all collisional/radiative pumping models which predict that various vibrational state masers should peak in different spatial regions of the CSE."
    },
    "0403/astro-ph0403484_arXiv.txt": {
        "abstract": "We select E+A candidates from a spectroscopic dataset of $\\sim800$ field galaxies and measure the E+A fraction at $0.3<z<1$ to be $2.7\\pm1.1$\\%, a value lower than that in galaxy clusters at comparable redshifts ($11\\pm3$\\%).  HST/WFPC2 imaging for five of our six E+A's shows they have a heterogeneous parent population: these E+A's span a range in half-light radius ($0.8<r_{1/2}<8$\\hi kpc) and estimated internal velocity dispersion ($50\\lesssim\\sigma_{est}\\lesssim220$\\kms), and they include luminous systems ($-21.6\\leq M_{Bz}-5logh\\leq -19.2$).  Despite their diversity in some aspects, the E+A's share several common characteristics that indicate the E+A phase is an important link in the evolution of star-forming galaxies into passive systems: the E+A's are uniformly redder than the blue, star-forming galaxies that make up the majority of the field, they are more likely to be bulge-dominated than the average field galaxy, and they tend to be morphologically irregular.  We find E+A's make up $\\sim9$\\% of the absorption line systems in this redshift range, and estimate that $\\gtrsim25$\\% of passive galaxies in the local field had an E+A phase at $z\\lesssim1$. ",
        "introduction": "Studies show the universal star formation rate has been declining steadily since $z\\sim1$ \\citep{lilly:95,cowie:96,madau:96}.  However, {\\it why} star formation in the universe has been decreasing for the last $\\sim8$ Gyr and why it ends for any given galaxy remain unanswered questions.  Galaxies whose star formation ends abruptly in a burst can provide insight into these issues because they may be the link between star-forming and passive systems.  Identified in significant numbers in intermediate redshift galaxy clusters \\citep[hereafter T03a]{dressler:83,couch:87,tran:03b}, ``E+A'' galaxies are characterized by strong Balmer absorption lines and weak or absent [OII]$\\lambda3727$ emission.  They have been studied extensively in clusters \\citep[T03a]{dressler:83,couch:87,franx:93,wirth:94,couch:94,belloni:95,couch:98,fisher:98,caldwell:99,dressler:99,balogh:99,poggianti:04}, and are also referred to as H$\\delta$-strong or ``k+a'' galaxies. Their distinctive spectral features are usually interpreted as being the result of a starburst ($M_{SB}\\gtrsim0.1M_{\\ast}$, where $M_{\\ast}$ is the galaxy's final total stellar mass) that ended within the last $1.5-2$ Gyr \\citep[although see Newberry et al. 1990 for an alternative scenario]{couch:87,barger:96}.  In addition to being in rich clusters, E+A galaxies have also been found in the field.  The properties of field E+A's have been studied extensively at low redshift \\citep[$z<0.2$;][]{liu:95,zabludoff:96,chang:01,norton:01,goto:03,quintero:04,yang:04}, but these low redshift field E+A's do not trace star formation beyond $z\\sim0.3$ because of the short lifetime of the E+A phase.  To determine how important a role the E+A phase has in the evolution of field star-forming galaxies into passive systems, a population of field E+A galaxies at $z>0.3$ is needed.  Finding E+A galaxies is challenging because the E+A phase is visible for only $\\sim1.5$ Gyr \\citep{couch:87,barger:96} and E+A's are only a small fraction of the total galaxy population at any given redshift. These difficulties are reflected in the highly uncertain estimates for the field E+A fraction at $z>0.3$ where values range from $\\sim1-5$\\% \\citep{hammer:97,dressler:99,balogh:99}.  Estimates of the field E+A fraction also can vary due to, e.g.  selection criteria and spectral resolution.  Assuming the E+A phase is associated with the transformation of an emission line galaxy into an absorption line system, a high E+A fraction would imply that the majority of passive galaxies in the nearby universe go through an E+A phase.  Given the formidable obstacle of even isolating the field E+A population at $z>0.3$, it is no surprise then that we know very little about their physical characteristics and their progenitors.  Studies of E+A's in the nearby universe \\citep{zabludoff:96,caldwell:99,galaz:00,norton:01,poggianti:04,yang:04} and in intermediate redshift clusters \\citep[T03a]{wirth:94,dressler:99} indicate these E+A's have a heterogeneous parent population.  By examining their physical properties, we can determine if E+A's in the intermediate redshift field are descended from a diverse population as well.  This information also enables us to determine whether field E+A's at $z>0.3$ can be the progenitors of early-type galaxies at lower redshifts.  What triggers the E+A phase remains an open question.  Earlier studies suggest the E+A phase in the field is associated with galaxy-galaxy interactions \\citep[hereafter Z96]{liu:95,yang:04,zabludoff:96}, a scenario consistent with galaxy formation models that predict interactions in the field are common \\citep{somerville:99a}, and that interactions trigger strong star formation episodes in gas-rich systems \\citep{mihos:94a,mihos:94b}.  Because morphological signatures due to interactions remain visible for the duration of the E+A phase \\citep{mihos:94a,mihos:94b}, we should find E+A's to be morphologically disturbed if the E+A phase is associated with galaxy-galaxy interactions.  To measure the field E+A fraction at intermediate redshifts, we draw from an extensive spectroscopic survey of four different fields.  Our unique dataset contains an unsually large number of spectroscopically confirmed field galaxies ($\\sim800$) gathered over a total area of $\\sim200\\Box'$; our survey is larger than that of any other previous field E+A study in this redshift range.  Using the wide-field HST/WFPC2 imaging that we have obtained as part of our cluster program, we also characterize the physical properties of the field E+A population to better understand the role E+A's have in the evolution of star-forming galaxies into passive systems.  We summarize our observations in \\S2, describe our E+A selection criteria and determine the field E+A fraction in \\S3, and detail the physical properties of the field E+A population in \\S4.  We discuss the importance of the E+A phase in \\S5 and present our conclusions in \\S6. Unless otherwise noted, we use $\\Omega_M=0.3, \\Omega_{\\Lambda}=0.7$, and $H_0=100h$\\kms~Mpc$^{-1}$ in this paper. ",
        "conclusions": "We select E+A galaxies from a spectroscopic dataset of $\\sim800$ field galaxies and measure the E+A fraction in the intermediate redshift field ($0.3<z<1$) to be $2.7\\pm1.1$\\%.  Because we select the field E+A's in the exact same manner as E+A's in galaxy clusters at comparable redshifts, we can compare directly how the E+A fractions differ between the two environments.  We find the E+A fraction in the field is a factor of four smaller than that observed in intermediate redshift galaxy clusters ($11\\pm3$\\%; T03a).  However, most E+A galaxies ($\\sim70$\\%) lie in the field because only a small fraction of galaxies are in clusters.  From HST/WFPC2 imaging, we present first results on the physical properties of field E+A's at $z>0.3$.  The five E+A galaxies that fall on the WFPC2 mosaics include luminous systems ($-21.6\\leq M_{Bz}-5\\log h\\leq -19.2$), and they span a range half-light radius ($0.8\\lesssim r_{1/2}<8.2$\\hi kpc) and estimated internal velocity dispersion ($50\\lesssim\\sigma_{est}\\lesssim220$\\kms).  These properties indicate that like their lower redshift counterparts, the E+A's in the intermediate redshift field have a heterogeneous parent population.  Despite being descended from a variety of progenitors, these E+A's have several characteristics in common.  Compared to the general field population, the E+A's are uniformly redder than the blue, star-forming galaxies that make up the majority of the field.  These E+A's also are more likely than the average field galaxy to be a bulge-dominated system, and they tend to be morphologically irregular.  These characteristics indicate the E+A phase is associated with galaxy-galaxy interactions, and that E+A's are an important link between star-forming and passive galaxies.  We find E+A's make up $\\sim9$\\% of the absorption line galaxies at $0.3<z<1$.  Combining this with timescale arguments, we estimate that $\\sim25$\\% of passive galaxies in the nearby field had an E+A phase at $z<1$.  This fraction increases to $\\sim70$\\% if we consider only galaxies that are morphologically typed as E/S0's.  Our study demonstrates that the E+A phase can have as significant a role in how galaxies evolve into early-type systems in the field as in galaxy clusters.  However, we recognize our results are based on a small sample of E+A's, and we stress that additional studies of this interesting population are needed."
    },
    "0403/astro-ph0403351_arXiv.txt": {
        "abstract": "In an alternative mechanism recently proposed, adiabatic cosmological perturbations are generated at the decay of the inflaton field due to small fluctuations of its coupling to matter. This happens whenever the coupling is governed by the vacuum expectation value of another field, which acquires fluctuations during inflation. We discuss generalization and various possible implementations of this mechansim, and present some specific particle physics examples. In many cases the second field can start oscillating before perturbations are imprinted, or survive long enough so to dominate over the decay products of the inflaton. The primordial perturbations will then be modified accordingly in each case. ",
        "introduction": "The detection of the cosmic microwave background (CMB) anisotropies indicates the presence of coherence over super horizon scales, which is a strong indication of an early inflationary stage~\\cite{inflation}. To date measurements are in agreement with the simplest prediction from inflation which is a nearly scale invariant spectrum of gaussian and adiabatic primordial perturbations~\\cite{wmap}. Nevertheless, possible experimental or theoretical deviations from this most economical possibility have been the subject of intense study.  Indeed, the increasing precision of the CMB measurements gives us the hope to raise the bar in the near future from the present confirmation of the general idea of inflation, to the study and determination of the underlying theory.  An important aspect of inflation is the stage of reheating~\\cite{reheat}, which describes all the processes from the end of the inflationary expansion to the following hot ``big-bang'' evolution. The inflaton decay is only the first, although most relevant, stage of this process. The inflaton decay is typically very quick and non-perturbative, and it occurs immediately after the end of inflation~\\cite{preheat}. Due to its efficiency, it leads to a distribution of particles which is very far from thermal equilibrium, so that reheating completes over a much longer timescale. Particle physics plays a relevant role in this period. It seems conceivable that the baryon asymmetry of the universe is generated at this stage. On the other hand, nucleosynthesis poses strong bounds on the production of gravitationally decaying relics. These requirements can be combined to constrain different models of reheating. For example, limits from nucleosynthesis are in contrast with thermal grand unified theory (GUT) baryogenesis and only marginally compatible with thermal leptogenesis, while a non-thermal origin~\\cite{non} can be more easily accounted for. However, while these considerations allow us to study the inflaton interactions for any given particle physics model, we are not yet in a stage where we can discriminate among different scenarios. As in many other areas of physics, we are more in the need of experimental evidence rather than theoretical models. It is therefore important to ask whether reheating can have some other observational consequences which can guide us discriminating among the different possibilities.  A positive answer to this question is provided by the recent observation that reheating could have played a key role also for generating the primordial perturbations leaving their imprints on the CMB, as well as on the matter power spectrum~\\cite{dgz,kofman}. This statement challenges the common assumption that the microphysics responsible for the decay of the inflaton, and the successive thermalization, cannot have an impact on the much larger scales relevant for the present observations which at that time are well outside the horizon. However, perturbations generated during inflation can be strongly affected by the following history of the universe. For example, the adiabatic mode of the perturbations is sensitive to the evolution of dominating species, and it is precisely during reheating that this evolution is mostly unknown. Although we can be sure of the starting (unstable inflaton) and final (thermal bath) points, we do not know what dominates during the intermediate stages and how its equation of state evolves. Several possibilities can occur, depending on the inflaton potential near its minimum (for example, $m^2 \\phi^2$ versus $\\lambda \\phi^4$ potential), on the possible generation of unstable heavy particles which can dominate for some time, on the possible phase transitions and on the presence of large effective masses before thermalization completes.  These processes can strongly affect primordial perturbations if they occur at (slightly) different moments in the different parts of the universe we presently observe. This can happen, for example, for the decay of the inflaton, if its coupling to matter is controlled by the vacuum expectation value (VEV) of some scalar field $X\\,$ which acquires inflationary fluctuations~\\cite{dgz,kofman}. The equation of state of the inflaton is in general different from the one of its decay products, and fluctuations in the decay rate will then give rise to fluctuations in the energy density and in the metric. Perturbations generated by this mechanism have been named {\\em modulated perturbations}. However, modulated perturbations do not need to be generated precisely at the inflaton decay. They could also arise due to fluctuations in the decay rate of some other intermediate particle, or in the rate of the processes which are more relevant for thermalization. In general, we will denote by $t_*$ the time at which this process completes, with the assumption that it can be anytime during reheating. This process changes the equation of state of the universe from an unknown one $p=w\\rho$ to $p = \\rho/3\\,$, which is typical of radiation. Due to the change in $w$, fluctuations in $t_*$ give rise to adiabatic perturbations.  A further extension on which we want to focus, noted briefly in~\\cite{kofman}-\\cite{gruzinov}, is the possibility that the field $X$ is evolving before or after $t_*\\,$. While $X$ is initially frozen at some (nearly) constant value, it will start oscillating around the minimum of its potential (with a slowly decreasing amplitude) as soon as the Hubble expansion rate drops below its mass. This can significantly affect (typically decrease) the amplitude of modulated perturbations if oscillations of $X$ start before perturbations are generated. Moreover, as any other modulus, $X$ can eventually dominate the energy density of the universe. Isocurvature fluctuations in $X$ will in this case be converted to curvature perturbations in the same way as the curvaton scenario~\\cite{curvaton,moroi}. Then perturbations in the new (dominating) thermal bath will replace modulated perturbations. A particular identification of $X$, and requirements for a successful reheating, may favour one of these cases.  The analysis of Section~\\ref{general} covers all these different possibilities. From the general discussion, we then specialize to some specific examples motivated by particle physics in Section~\\ref{examples}. The main ingredient of the modulated perturbations mechansim is the existence of a scalar field with a mass less than the Hubble expansion rate during inflation. Theories based on supersymmerty provide the best framework for implementing this mechansim as they contain many scalar fields whose mass is protected against quantum corrections. We first discuss the possibility that $X$ is a modulus field, with a mass $\\sim$ TeV and interactions of gravitational strength with the observable sector. This is probably the most natural implementation of the idea of modulated perturbations: moduli of superstring models have these properties, and their VEV's control the value of the parameters of the theory. The difficulty with this implementation, as we shall see, is that it typically results in a too small amplitude for the primordial perturbations, unless $X$ dominates. In this case, however, one has to make sure that the late (gravitational) decay of $X$ occurs before primordial nucleosynthesis. Alternative examples that we consider are the identification of $X$ with right-handed (RH) sneutrinos and with flat directions of the minimal supersymmetric standard model (MSSM), where the decay of $X$ occurs at an earlier stage. In general, other scalar fields (in particular those coupled to $X$) should have negligible fluctuations in order not to affect perturbations. In Section~\\ref{discuss} we discuss possible ways to achieve such a suppression leading us to some ``benchmark scenarios'' of modulated perturbations. The results will be summarized in the concluding Section~\\ref{end}. ",
        "conclusions": "\\label{end}  In this paper we have studied different scenarios of modulated perturbations. In this mechanism adiabatic cosmological perturbations are generated during reheating when the equation of state in the different parts of the universe changes at different moments. The inhomogeneity arises due to inflationary fluctuations of some scalar field $X$ whose VEV controls the value of parameter(s) (such as mass or coupling) involved in the relevant process. This is a generic property expected in models based on superstring and supersymmetric theories. Depending on the values of the rate of the process $\\Gamma_*$, mass of the scalar field $m_X$ and its decay rate $\\Gamma_X$, different situations can arise. We presented a general analysis of possible scenarios, with the main results summarized in fig.~(2) and fig.~(3). We also introduced some specific particle physics examples including the string moduli, RH sneutrinos and supersymmetric flat directions. These examples, in which $m_X$ is of the same order or greater than the gravitino mass $m_{3/2}$, represent a wide range of particle physics candidates of $X$.  The simplest possibility is that perturbations are generated when $X$ is frozen at an initial value $X_I$. If $m_X \\geq {\\cal O}(\\rm TeV)$ (as for the RH sneutrinos and supersymmetric flat directions), this requires a very rapid process, namely non-perturbative inflaton decay via preheating. Moreover, $X$ must decay before dominating the energy density of the universe, and the reheat temperature should be sufficiently low to avoid overproduction of gravitinos. For a slower (likely perturbative) process, $X$ can start oscillating around the minimum of its potential before perturbations are imprinted. Then the expansion of the universe redshifts fluctuations of $X$ and dampens perturbations. An interesting consequence is that the level of non-gaussianity can in this case be just below the current limits and accessible to future experiments. If $X$ is long-lived (for example, a modulus field), it can dominate the universe before decaying. In this case the situation will essentially be the same as the curvaton scenario and fluctuations in the energy density of $X$ give rise to perturbations. The $X$-dominance can be a virtue when the reheat temperature exceeds the gravitino bound necessitating late time entropy production. The amplitude of the adiabatic mode is in general very different from one case to another. It interpolates between the different values and bordering regions are charachterized by a significant isocurvature component.  In conclusion, the modulated perturbations mechanism is a viable alternative with potentially interesting observational consequences. It can be implemented in various ways and a successful particle physics identification can lead us to a clearer and more complete picture of reheating."
    },
    "0403/astro-ph0403398_arXiv.txt": {
        "abstract": "The inner annular gap (IAG), a new type of inner gap whose magnetic field lines intersect the null charge surface (NCS), is proposed to explain $\\gamma$-ray and radio emission from pulsars. The IAG can be an important source for high-energy particles. The particles can radiate between the NCS and the IAG. Some observational characteristics in both $\\gamma$-ray and radio bands, such as the $\\gamma$-ray emission beams of Crab-like, Vela-like and Geminga-like, can be reproduced by numerical method. It is predicted that the view angle $\\zeta$ should be larger than the inclination angle ($\\zeta>\\alpha$), otherwise the $\\gamma$-ray radiation will have little possibility to be observed. Whether the inner annular gap (or cap) is sparking (or free flow) depends on the surface binding energy of the pulsar. In stead of neutron star models, the scenario of the IAG is favorable for bare strange star models, because bare strange stars can easily satisfy the requisite condition to form an IAG for both pulsars ($\\vec{\\Omega}\\cdot \\vec{B}<0$) and anti-pulsars ($\\vec\\Omega \\cdot \\vec{B}>0$). ",
        "introduction": "Both polar cap and outer gap models of $\\gamma$-ray pulsars have been suggested to explain the high energy emission from pulsars \\citep{A81,H81,HM98,MH03,Z89,L90,L94,C76,R81,CHR86,C00,H00,R96,R02}. Here, we propose a new scenario of the inner annular gap (hereafter IAG) outside the conventional inner gap. The IAG is defined by magnetic field lines which intersect the null charge surface(NCS) (see Fig.\\ref{GJcg} ). As we know, at any point on the NCS the direction of local magnetic field is perpendicular to the rotation axis. So if the radiation region locates near the NCS (either inside or outside), the pulsar can easily produce a very wide pulse profile \\citep{Q03a}. This wide profile will match the observed $\\gamma$-ray pulse profiles rather well. Some authors \\citep{Q03a,Q03b,DR03} suggested that if the emission region locates inside the NCS, it would be favorable to understand the $\\gamma$-ray emission. Here we suggest that the particles flow out from the IAG could be re-accelerated and radiate $\\gamma$-rays inside the NCS, and the model to account for the $\\gamma$-ray and radio emission at the same time. In \\S 2 the basic ideas for the IAG are presented. In \\S 3 we show how $\\gamma$-ray and radio emissions can be produced. Conclusions and brief discussions are given in \\S 4. ",
        "conclusions": "We emphasize here that the IAG and the NCS play an important role in pulsar $\\gamma$-ray radiation. At least the IAG can be a source of high energy particles for pulsars emission. If the radiation region locates near the NCS, many radio and $\\gamma$-ray observational facts can be easily understood. The large phase separation of components of $\\gamma$-ray light curves and the relative compact radio phase separation reveal that the $\\gamma$-ray radiation region should locate near the NCS. The geometry calculation partly supports this radiation location.  It is indicated from the calculation above that the condition of $\\gamma$-ray pulsars which we observed is that the view angle $\\zeta$ should be large than the inclination angle, i.e. $\\beta>0$. This is because that the part of IAG between the magnetic axis and the equator is wider so that it gain higher potential drop, when the inclination angle becomes larger.  The luminosity of ${\\gamma}$-ray pulsars in the IAG model depends on the maximum potential drop across the gap, $\\Delta V=0.5\\Omega B r_{\\rm p}^2c^{-1}$. For a given  Lorentz factor we find that the $\\gamma$-ray luminosity are proportional to ${\\dot P}^{0.5} P^{-1.5}$, which is in agreement with observations.  Our IAG scenario is intrinsically different from the free-flowing type of polar cap model \\citep{MH03}. How does the IAG interact with the outer gap? Are pulsars bare strange stars or special kind of neutron stars \\cite{G02}? Does the radiation come from the re-accelerated non-stationary pair flow or free flow? All these questions need to be investigated further more. The details of the IAG model, such as the re-acceleration and radiation processes, $\\gamma$-ray luminosity, spectral behaviors, emission beam properties and so on, will be discussed in separated papers."
    },
    "0403/astro-ph0403077_arXiv.txt": {
        "abstract": "We use a high-resolution simulation of a galaxy-sized dark matter halo, published simulated data as well as four cluster-sized haloes from Fukushige, Kawai \\& Makino to study the inner halo structure in a $\\Lambda$ cold dark matter cosmology. We find that the circular velocity curves are substantially better described by \\citeauthor{Stoehr_et_al_02} (SWTS) profiles than by Navarro, Frenk \\& White (NFW) or Moore et al. profiles. Our findings confirm that no asymptotic slope is reached and that the profiles are nearly universal, but not perfectly. The velocity profiles curve at a {\\it constant} rate in $\\log(r)$ over the full converged range in radii and the corresponding extrapolated density profiles reach a finite maximum density. We find, that the claim of a strong discrepancy between the observed inner slopes of the density profiles of low surface brightness galaxies (LSB) and their simulated counterparts on the grounds of currently available observations and simulations is unfounded. In addition, if the SWTS profile turns out to be a good description of the halo profile for the regions that cannot be probed with simulations of today, then even in these regions the agreeement between simulations and observations is very reasonable. ",
        "introduction": "One of the greatest challenges to the $\\Lambda$ cold dark matter (CDM) model favoured today is probably the apparent discrepancy between the very inner parts of observed dark matter (DM) density profiles and their counterparts in cosmological N-body simulations of structure formation.  Numerical simulations seem to find nearly `universal' density profiles with cusps mostly being described by double-power-law fits with outer slope $\\beta={\\mbox d}\\log\\rho/{\\mbox d}\\log r\\approx-3$ and inner slopes between -1 and -1.5. These profiles generally match the simulated curves reasonably well for scales of $r > 0.05 \\ r_{200}$ where the virial radius $r_{200}$ is the radius of a sphere around the halo centre enclosing a density 200 times larger than the critical density (\\citealt*{Navarro_Frenk_White_97}, NFW; \\citealt*{Tormen_Bouchet_White_97}). The exact value of the inner slope, which had been constantly re-examined as computing power and thus mass resolution increased, is still a matter of debate (\\citealt{Fukushige_Makino_97,Moore_Governato_Quinn_Stadel_Lake_98,Moore_et_al_99cold,Ghigna_et_al_00,Jing_Suto_00,Fukushige_Makino_01,Klypin_et_al_01,Jing_Suto_02,Fukushige_Makino_03,Power_et_al_03}). The best resolved haloes \\citep*{Fukushige_Kawai_Makino_04} with up to $3 \\times 10^7$ simulation particles within the virial radius showed slopes at the resolution limits of the simulations that were significantly shallower than -1.5.  Extrapolating the fits to smaller distances results in a strong discrepancy with the observed rotation curves of DM-dominated low surface brightness (LSB) galaxies. Their DM haloes are consistent with profiles having constant density cores \\citep{Flores_Primack_94, Moore_94, Burkert_95, McGaugh_Blok_98, Firmani_et_al_01}.  Although the extent of this discrepancy seemed controversial \\citep[e.g.][]{vandenBosch_Swaters_01}, the most recent studies \\citep*{Blok_Gaugh_Rubin_01, deBlok_04} claim that simulations can not be reconciled with observations. Several observational effects like seeing, misalignment of the slit of the telescope, finite slit size, beam smearing due to the large beams of the radio telescopes or an offset between the dynamical and the optical centre certainly matter. However, they are too small to make the NFW profile with inner slope of -1 mimic observed rotation curves \\citep{deBlok_04}.  In this work we re-investigate the inner structure of DM haloes simulated at very high resolution and compare the results to published LSB galaxy data. In the next section we describe the simulation we have carried out and the simulated data we have used. We then give the results of the profile-fitting and discuss the implications. In the last section we summarise and discuss our conclusions.  \\begin{figure*} \\centering \\includegraphics[width=88mm]{circularvelocity.figure1.colour.eps} \\centering \\includegraphics[width=88mm]{circularvelocity.figure2.colour.eps} \\caption{{\\it Left-hand panel:} circular velocity curves for GA3new, G3-256$^3$ \\citep{Hayashi_et_al_04} and P256$^3$ \\citep{Power_et_al_03} (thick solid, red) with over-plotted SWTS (thin solid), Moore (dashed) and NFW profiles (dotted). Vertical lines indicate the radius down to which the profiles are converged. {\\it Right-hand panel:} circular velocity curves of the four cluster simulations carried out by \\citet*{Fukushige_Kawai_Makino_04}. For clarity all curves have been shifted down by multiples of a quarter decade. The $a$-values of the seven SWTS profiles are       0.150, 0.128, 0.139, 0.102, 0.0934, 0.149 and 0.135. \\label{figure:circularvelocity}} \\end{figure*} ",
        "conclusions": "We have reinvestigated the inner structure of DM haloes simulated in a $\\Lambda$CDM cosmology using the best resolved galaxy and cluster haloes existing to date. The studied haloes were produced by different groups using three different simulation codes. We find that the concept of an `inner slope' of the circular velocity curves or density profiles is not appropriate: the profiles continue to curve without reaching an asymptotic slope in $\\log(r)$.  The simplest analytical velocity profile with such property, the profile with  constantly changing slope, is the SWTS profile. It describes the analysed DM haloes typically at a level of better than one per cent over the full range from the converged to the virial radius. This is substantially better than the values we obtain for NFW or Moore profiles.  Allowing for one more free parameter in the analytic profiles, the shape parameter $a$, is necessary in order to provide fits of high accuracy. It seems plausible that the small scatter in shape parameter values reflects the differences in non-sphericity of the DM distributions. We thus find that simulated dark matter profiles are nearly universal but not perfectly so. Our $a$ values range from 0.093 to 0.15 with a median value of $a_{med}$=0.135.  It is interesting that the same simple profile accurately describes both the profiles of isolated haloes as well as those of strongly perturbed substructure haloes \\citep{Stoehr_et_al_02,Stoehr_et_al_03}.  There has been a claim of a strong apparent discrepancy between the innermost density profile slopes of observed LSB galaxies and their simulated counterparts. We find that this claim is unfounded. It relys on an extrapolation of the analytical profiles beyond the actual resolution limits of the numerical simulations. Doing so is particularily dangerous, as we have shown that NFW and Moore profiles, which were used to support the claim, show systematic deviations from the simulated profiles already in the numerically converged regions.  Indeed, when extrapolating the SWTS profile (which fits the simulated profiles substantially better than the NFW or Moore profiles) instead, we find an astonishingly good agreement with the innermost profile slopes of observed LSB galaxies.  We show in addition that then the shape of a `typical' LSB rotation curve can be very well reproduced. At least some of the remaining substantial disagreement for the shapes and slopes of the LSB galaxies from the tail of the distribution may be explained with observational effects.  \\noindent"
    },
    "0403/astro-ph0403594_arXiv.txt": {
        "abstract": "{ We give  independent proof of the deficit of stars in the in-plane central disc ($2<R<4$ kpc, $|b|\\la 3^\\circ $) with respect to the predictions of a pure exponential density distribution.\\\\ We use three different methods: 1) the inversion of the red clump giant distribution in near-infrared colour--magnitude diagrams  to obtain the star density along the line of sight; 2) the determination of the density distribution of 1612 MHz sources by means of the distance determination of OH/IR sources from their kinematical information; 3) an analysis of near- and mid-infrared star counts and comparison with models. All the tests give the same result: a deficit of stars in the inner disc with respect to an exponential disc (either with constant scaleheight or extrapolated from the outer disc), but only in near plane regions ($|b|\\la 3^\\circ $). This deficit  might be interpreted as a flare in the vertical distribution. The in-plane density is almost independent of $R$ and not an exponential law of the type $\\rho \\propto \\exp (-R/h)$. Further away from the plane, however, the density increases towards the centre due to the increase of the scaleheight. Tests also show that this result cannot be due to  extinction. This deficit affects both  the young  and the old populations, so it is probably a rather stable feature of the disc, and might be due to the existence of an in-plane bar  sweeping the near-plane stars.\\\\ An approximate expression of the disc density within $2<R<8$ kpc is: $\\rho(R,z)\\propto e^{-\\left(\\frac{R}{1970\\ {\\rm pc}}+ \\frac{3740\\ {\\rm pc}}{R}\\right)}e^{\\frac{-|z|}{h_z(R)}}$, with $h_z(R)\\approx 285[1+0.21\\ {\\rm kpc}^{-1}\\ (R-R_\\odot)+0.056\\ {\\rm kpc}^{-2}\\ (R-R_\\odot)^2] \\ {\\rm pc}$. ",
        "introduction": "L\\'opez-Corredoira et al. (2002, hereafter L02)  obtained a model of the outer parts of the stellar disc (galactocentric radius $R$ greater than 6 kpc), based on the 2MASS near infrared point sources survey. Among their findings are the stellar warp coincident with the gas warp, the flare, the absence of an external cut-off out to $R=15$ kpc and the measurement of the different parameters of the old population density distribution, generally in agreement with other independent studies. Here, we  continue the analysis of the disc structure in the inner stellar disc, examining whether or not this corresponds  to an extrapolation of the L02 disc model towards the centre and, if not, what  the differences are.  A major concern is whether  extrapolation of the exponential in the inner stellar disc is appropriate or whether there is some deficit of stars with respect to this distribution. It is well known that many barred galaxies have holes (Ohta et al.\\ 1990); up to half of them, according to Anderson et al.\\ (2002) based on  analysis of the bulge-disc decomposition carried out in Baggett et al.\\ (1998). In a large number of external galaxies, due to their more favourable orientation, the central hole in the disc has been detected in the distribution of molecular and atomic gas. See, for example, Aguerri et al.\\ (2001), who describe a hole in the HI distribution in NGC 5850 associated with the central bar region, as is the case presented by Bottema \\& Verheijen (2002) in NGC 3992. These holes have been attributed to the presence of gas velocities perpendicular to the Galactic plane (Sancisi 1999). However, there are also cases where the CO distribution, used to infer the HII density, does not reveal such a hole, as is the case, for example, of Sakamoto et al.\\ (1999), whose sample is somewhat biased towards the non-hole gas distribution because of  selection criteria. The hole does not usually consist strictly of an abrupt truncation of the disc. It can represent a gradual density decrease (for instance, Freeman type II truncation; Freeman 1970), instead of increasing towards the center within a certain radius, or even a constant density or increasing density but with a slope much lower than that produced by the extrapolated exponential law.  In our Galaxy, this hole (from now on, we will call the hole a ``deficit'' of gas or stars since, strictly speaking, the hole is not totally empty;  ``deficit'' here is always with respect to the extrapolation of the purely exponential disc towards the inner radii) is observed in the gas distribution (i.e., Robinson et al.\\ 1988). For the stellar distribution the answer is not very clear at present. Several disc models have been produced using COBE/DIRBE flux maps. Although the general features are approximately coincident, it seems that there is no agreement regarding the deficit of stars. For instance, Freudenreich (1998) and L\\'epine \\& Leroy (2000) agree that there is a such a deficit, while Bissantz \\& Gerhard (2002, hereafter B02) can reproduce quite accurately the flux maps with an exponential disc even in the inner parts. Caldwell \\& Ostriker (1981) gave a model of the disc that consists of a difference of two exponentials with different scalelengths, which gives a deficit of stars in the inner part (see fig. 15 in Caldwell \\& Ostriker 1981) based on kinematical measurements. Microlensing experiments seemed to favour a deficit of stars (Kiraga \\& Paczy\\'nski 1994), but the poorly--known interstellar extinction at optical wavelengths gives some uncertainties in their models. Moreover, Kiraga et al. (1997) recognize that their previous paper might contain some errors due to an incorrect identification of populations in the optical colour--magnitude diagrams. Analysis of OH/IR sources, both from the young and intermediate--old populations, led Baud et al.\\ (1981) to conclude that there must be a deficit of stars in the two populations; these results are more convincing, and we think that a corroboration of this result with newer surveys will be valuable (see \\S \\ref{.OHIR}). This topic has so far not received much attention in terms of star counts. L\\'opez-Corredoira et al.\\ (2001, hereafter L01) have claimed that a deficit of stars is necessary to reproduce some near-plane counts but they did not try to fit off-plane regions, and not too much attention was paid to the problem in that paper. Precisely because of this, we wish to examine the possible deficit of stars in the inner disc, using point source catalogues, which contain more information than flux-maps alone.  A further area of study is the real nature of the deficit of stars, should this be genuine. A flared stellar distribution places the stars at a higher elevation from the plane than that predicted from a pure exponential law, hence giving the impression of a hole in the plane. Flared structures are well recognized in the distribution of gas (see Burton \\& Hekkert 1986) and stars (L02; Alard 2000, and references therein) in the outer disc. The presence of that type of morphology in the central disks of spiral galaxies is not so well studied due to the inherent difficulties in separating the several structural contributions to the observed counts or flux. A good example of that difficulty can be found in Nikolaev et al. (2004), who reported a symmetric warp in the inner disk, the central 4 degrees, of the LMC  based on accurate analysis of MACHO and 2MASS counts. ",
        "conclusions": "\\begin{figure} \\begin{center} \\mbox{\\epsfig{file=0385fig26.ps,height=8cm,angle=90}} \\end{center} \\caption{Contour diagram of $\\log _{10}\\rho $ (pc$^{-3}$) in the $yz$-plane of the Galaxy (perpendicular to the line Sun--Galactic centre) with $y$ between $-$12.0 and 12.0 kpc and $z$ between $-$1.0 and 1.0 kpc (vertical scale in the plot multiplied by a factor 5). Lower contour: $\\log _{10}\\rho $ (pc$^{-3}$)=-2.1; step=0.15. Model: L02 for $R>4$ kpc, and eq. (\\ref{hole}) for $R\\le 4$ kpc. The abrupt transition (at $|y|=R=4$ kpc) is due to the discontinuous change of regime between the inner disc and the outer disc. Note also that the model of this paper is valid only for regions not very far from the plane.} \\label{Fig:contourinner} \\end{figure}  \\begin{figure} \\begin{center} \\mbox{\\epsfig{file=0385fig27.ps,height=8cm,angle=90}} \\end{center} \\caption{Contour diagram of $\\log _{10}\\rho $ (pc$^{-3}$) of a possible interpolated/extrapolated smooth model of the disc according to eq. (\\ref{expdob}) for $R<R_\\odot $ and L02 for $R\\ge R_\\odot $ in the $yz$-plane of the Galaxy (perpendicular to the line Sun--Galactic Centre) with $y$ between $-$12.0 and 12.0 kpc and $z$ between $-$1.0 and 1.0 kpc (vertical scale in the plot multiplied by a factor 5). Lower contour: $\\log _{10}\\rho $ (pc$^{-3}$) = $-$2.1; step = 0.15.} \\label{Fig:contourinner2} \\end{figure}   We have used three different methods: 1) the inversion of the red clump giant distribution in near-infrared colour--magnitude diagrams  to obtain the star density along the line of sight; 2) the determination of the density distribution of 1612 MHz sources by means of the distance determination of OH/IR sources from their kinematical information; 3) the analysis of near-infrared star counts and comparison with models. All the tests agree with the following picture for the inner disc:  \\begin{itemize}  \\item There is a deficit of stars with respect to an exponential disc in the Milky Way's inner stellar disc for 2.25 kpc $<R<4$ kpc. The density in the plane ($b\\approx 0^\\circ$, i.e. $z\\approx 0$) is almost independent of $R$ instead of being an exponential law of the type $\\rho \\propto \\exp (-R/h)$.   \\item There is a slight increase in scaleheight in the central disc, $h_z=509-48R({\\rm kpc})$ pc, so the combined effect of deficit of stars in the plane and higher scaleheight produces an important deficit of stars for low latitude regions, although not for high latitude regions. Roughly, the density follows an expression of the type $\\rho (R,z)\\approx 5.1\\pm 0.6\\rho_\\odot e^\\frac{-|z|}{509\\pm 67-48\\pm 20 R(kpc)\\rm \\ pc}$, a result that is also compatible with a flared inner disc with a fixed scalelength of 2.4 kpc: $\\rho (R,z)\\approx\\rho_\\odot e^\\frac{-(R-R_\\odot)}{2.4 {\\rm kpc}} e^\\frac{-|z|}{509\\pm 67-48\\pm 20 R(kpc)\\rm \\ pc} \\frac{285\\ {\\rm pc}}{509\\pm 67-48\\pm 20 R(kpc)\\rm \\ pc}$. It is our objective to investigate further the real nature of the inner stellar disc, with a more complete analysis of 2MASS \\emph{All Sky Release} star counts. See the graphical representation in Fig.\\ \\ref{Fig:contourinner} (both possibilities are numerically equivalent). Note that these numbers are also subject to several systematic errors ($\\la$ 15\\% in the density between latitudes 2 and 6 degrees), so the parameters might change; although not as much as to allow the fit of a purely exponential disc (which would require errors greater than 70\\%). Moreover, the deficit of stars is corroborated with the OH/IR sources and the star counts. These parameters are applicable only for regions not very far from the plane ($|b|<6^\\circ $).  \\item Both the young and old population have this deficit of stars, and at present no significant differences in their distributions have been deduced from the data. Presumably, a more accurate measure of the distribution might reveal some differences. At present we can say that, roughly speaking, the deficit of stars with respect the extrapolation of an exponential distribution is probably a rather stable feature of the disc, which might be due to the existence of an in-plane bar that sweeps out the near-plane stars.  \\item An expression that summarizes the L02 outer disc ($R>6$ kpc) and the inner disc (2.5 kpc $<R<4$ kpc) with a smooth transition between two regimes is:   \\[ \\rho(R,z)\\approx \\left[\\rho _\\odot e^{\\left(\\frac{R_\\odot}{1970\\ pc}+ \\frac{3740\\ pc}{R_\\odot}\\right)}\\right] \\]\\begin{equation} \\times e^{-\\left(\\frac{R}{1970\\ pc}+\\frac{3740\\ pc}{R}\\right)}e^{-|z|/h_z}, \\label{expdob} \\end{equation}  with \\[ h_z(R)\\approx285[1+0.21\\ {\\rm kpc}^{-1}(R-R_\\odot) \\]\\begin{equation} +0.056\\ {\\rm kpc}^{-2}(R-R_\\odot)^2] \\ {\\rm pc} .\\end{equation}  This density is plotted in Fig. \\ref{Fig:contourinner2}. Although this expression was not tested for all ranges (for instance, it is not tested around $R=5$ kpc or for $R<2$ kpc), this can be used as an approximate interpolation/extrapolation that models the disc in $R<R_\\odot $. For the outermost disc, it is better still  to use the expression of L02. The innermost disc ($R<2$ kpc) reveals with the expression (\\ref{expdob}) a very deep hole. Unfortunately, we cannot be sure of the existence of this deep hole there, since the disc and the bulge cannot be easily separated in that region. It is clear at least that between 2 and 4 kpc the deficit of stars is very significant.  \\end{itemize}  \\  Acknowledgments: We thank C. Alard, V. P. Debattista and P. L. Hammersley. This publication makes use of data products from 2MASS (which is a joint project of the Univ. of Massachusetts and the Infrared Processing and Analysis Center (IPAC), funded by the NASA and the NSF); the MSX Point Source Catalog (which was obtained from the NASA/IPAC Infrared Science Archive at Pasadena, California); DENIS (the result of a joint effort involving personnel and financial contributions from several Institutes, mostly located in Europe and supported financially mainly by the French Institute National des Sciences de l'Univers, CNRS, and French Education Ministry, the European Southern Observatory, the State of Baden-W\\\"urttemberg, and the European Commission under a network of the Human Capital and Mobility programme), the TMGS (based on observations made at Carlos S\\'anchez Telescope operated by IAC at Teide Observatory on the island of Tenerife), the  ATCA/VLA OH 1612-MHz survey (based on observations made by Sevenster et al.\\ at the Australia Telescope Compact Array and at the Very Large Array)."
    },
    "0403/gr-qc0403030_arXiv.txt": {
        "abstract": "We investigate the quasinormal mode frequencies for the massless Dirac field in static four dimensional $AdS$ space-time. The separation of the Dirac equation is achieved for the first time in $AdS$ space. Besides the relevance that this calculation can have in the framework of the $AdS/CFT$ correspondence between M-theory on $AdS_4\\times S^7$ and SU(N) super Yang-Mills theory on $M_3$, it also serves to fill in a gap in the literature, which has only been concerned with particles of integral spin $0,1,2$. ",
        "introduction": "Perturbations of black holes are dominated, over intermediate timescales, by characteristic modes known as quasinormal modes. They are similar to normal modes of a closed system, but since the field can fall into the black hole or radiate to infinity, the modes decay and the corresponding frequencies are complex \\cite{chandra}. They have been extensively studied in asymptotically flat spacetimes and good reviews on the topic can be found in \\cite{nollert}, \\cite{kokkotas}. The study of quasinormal modes in Anti-de Sitter space was first done in \\cite{mann}.  More recently quasinormal frequencies have been investigated in the context of string theory in Anti-de Sitter space \\cite{horowitz}, \\cite{lemos}, \\cite{moss}. There is a suggestion, known as $AdS/CFT$ correspondence, that string theory in Anti-de Sitter space is equivalent to a conformal field theory in one dimension fewer \\cite{maldacena}. In the framework of this conjecture the study of $AdS$ black holes has a direct interpretation in terms of the dual conformal field theory on its boundary. The duality predicts that the retarded $CFT$ correlation functions are in one to one correspondence with Green's functions on Anti-de Sitter space with appropriate boundary conditions \\cite{danielsson}, \\cite{kalyana}, \\cite{birmingham}, \\cite{son}. Furthermore, as mentioned in \\cite{horowitz}, it is assumed that a large static black hole in $AdS$ space corresponds to a thermal state in the $CFT$ on the boundary. Perturbing the black hole is equivalent to perturbing this thermal state. The perturbed system is expected, at late times, to approach equilibrium exponentially with a characteristic time-scale. This time-scale is inversely proportional to the imaginary part of the poles of the correlators of the perturbation operator. It seems that these relaxation time-scales are quite complicated to calculate in the $CFT$, therefore their computation is conveniently replaced by the evaluation of the quasinormal frequencies in the $AdS$ bulk space. In \\cite{lemos} and \\cite{moss}, the authors went beyond the scalar perturbation treated in \\cite{horowitz} and they considered electromagnetic and gravitational perturbations. In one of the most recent works on the subject, Berti and Kokkotas \\cite{berti} confirm and extend previous results on scalar, electromagnetic and gravitational perturbations of static $AdS$ black holes and analyse, for the first time, Reissner-Nordstrom Anti-de Sitter black holes and calculate their quasinormal frequencies.  The interaction of a Dirac field with a black hole has been studied by Finster and his collaborators in a series of papers. While they have found stable particle-like solutions in the Einstein-Dirac- Maxwell system \\cite{finster1, finster2}, they have also proved the non-existence of time-periodic solutions in various black hole space-times \\cite{finster4, finster5}. This means that Dirac particles, including electrons and neutrinos, cannot remain on a periodic orbit around a black hole. In \\cite{finster3}, they also showed that there are no spherical symmetric black hole solutions in the Einstein-Dirac-Maxwell system other than the Reissner-Nordstr{\\\"o}m one. This suggests that if a cloud of Dirac particles undergoes gravitational collapse, the fermionic particles either vanish inside the event horizon of a black hole or escape to infinity.  The quasinormal frequencies related to the evolution of a massless Dirac field in a Schwarzschild black hole space-time, were studied in \\cite{cho}. The quasinormal modes of the Reissner-Nordstr{\\\"o}m de Sitter black hole for Dirac fields were studied by using the P{\\\"o}shl-Teller potential approximation in Ref. \\cite{Jing}. In this paper we are interested in the analysis of the modes of vibration of a massless Dirac field in a Schwarzschild Anti-de Sitter background space and it is for this reason that we compute the related quasinormal mode frequencies. This calculation is relevant to the $AdS/CFT$ correspondence between M-theory on $AdS_{4}\\times S^{7}$ and $SU(N)$ super Yang-Mills theory on $M_3$, and  it also serves to fill in a gap in the literature, which, in $AdS$ spaces, has only been concerned with particles of integral spin $0, 1, 2.$ ",
        "conclusions": "\\label{section4} In this paper we have investigated the quasinormal modes frequencies of a spinor field interacting with a Schwarzschild black hole in four dimensional Anti-de Sitter space-time. After solving the related axial perturbative equation we have imposed on this solution Dirichlet boundary conditions and evaluated the corresponding quasinormal frequencies. This choice of boundary conditions is physically justified in the contest of the $AdS/CFT$ correspondence, in which it is suggested to exploit these frequencies in the evaluation of poles of correlation functions associated to the conformal field theory on the boundary.  For large black holes, we have numerically confirmed that the imaginary part of the frequency scales with the temperature. We have also shown that both the real and imaginary components of the frequencies are evenly spaced in mode number. The spacing between consecutive modes, which is independent of the angular quantum number $l$, behaves as in the case of scalar, electromagnetic and gravitational perturbations. The quasinormal frequencies have a number of first overtones with pure imaginary parts.  For intermediate black holes, we have found that both the real and imaginary part of the frequency scale linearly with the surface gravity. The modes are evenly spaced and the spacing, which is independent of the angular quantum number $l$, resembles the one related to other kinds of perturbations.  Our purpose is to extend this project imposing different boundary conditions on the axial or the polar Dirac equations. We could be interested, for instance, in calculating the correlation functions and associated poles of a quantity related to the axial part of the spinor field. That would suggest we set the polar component of the spinor field perturbation to vanish on the boundary and consequently we could solve either the odd parity differential equation with Dirichlet boundary conditions or the even parity equation with mixed boundary conditions due to the transformation theory between polar and axial functions."
    },
    "0403/astro-ph0403241_arXiv.txt": {
        "abstract": "{We present an observational study of the $L_{\\rm X}$ vs $L_{\\rm B} \\sigma^2$ relation for early-type galaxies in the Coma cluster based on the XMM-Newton survey data. Compared to a similar relation for a sample dominated by field early-type galaxies, the Coma cluster galaxies show a flatter slope. Our calculations show that adiabatic compression produces a flattening in the $L_{\\rm X}$ vs $L_{\\rm B}\\sigma^2$ relation that is in remarkable agreement with the observed effect. Our scenario is further supported by the observed compactness of the X-ray emission of Coma galaxies. ",
        "introduction": "X-ray emission from early-type galaxies has been subject of investigations and theoretical modeling ever since its discovery with the Einstein telescope (Forman, Jones, Tucker 1985). The X-ray emission appears in part to be diffuse in nature and to originate in the galaxy hot interstellar medium (ISM); and in part to be due to a population of X-ray point sources, namely low mass X-ray binaries characterized by a hard spectrum.  In fact, the X-ray spectra are typically well fit with a two-component model: a soft $\\sim$1 keV thermal component for the hot ISM and and a hard $\\sim$1.4 photon index power law contribution for the point sources (Irwin et al. 2003). The latter component is systematically more important for fainter objects (Matsushita 2001).   The temperature of the ISM is proportional to the stellar velocity dispersion and, although with some scatter, the coefficient of proportionality is close to unity. In fact, the main energy input to the hot ISM is thought to be due to thermalization of kinetic energy associated with stellar mass loss, in addition to energy input from explosions of SN Ia (Mathews 1989; Ciotti et al. 1991; David et al. 1991). However, in some cases the galaxy ISM is found to be considerably hotter than the stellar velocity dispersion would indicate, as for example in the case of galaxies in the center of massive groups (Brown \\& Bregman 1998; Matsushita 2001).  In this respect the relation between X-ray and optical properties is important in order to understand the processes regulating the evolution of the hot gas. In fact, early on it was realized the existence of a correlation between the X-ray and the blue band stellar luminosity. Both the log-slope of the correlation and the significant scatter about it are physically meaningful and need to be properly understood and modeled.  In addition, X-ray and optical properties of ellipticals are used in combination for testing different scenarios of galaxy formation. Thus, Kodama \\& Matsushita (2001) used the continuity of optical properties across X-ray extended and compact early type galaxies to show that, if the different X-ray emission properties of the two classes are related to the structure of the underlying dark matter potential, then the formation of stars in these galaxies must predate the epoch when the potential structure was established.  In this {\\it Letter} we address the impact of high pressure cluster environment on the diffuse X-ray emission of early type galaxies. For the purpose we use the recent X-ray catalog of Coma cluster galaxies based on XMM-Newton survey and presented in Finoguenov et al. (2004a) and compare with the results reported in Matsushita (2001). In addition, we concentrate on X-ray compact galaxies as defined in Matsushita (2001). We find that X-ray emission from early-type galaxies in Coma is systematically higher when compared to their counterparts in the field or in loose groups, such as those presented in Matsushita (2001). We show that compression induced by the high pressure ICM is sufficient to reproduce the observed effect and point out that such environmental effects should be taken into account when modeling the X-ray properties of these galaxies.  This {\\it Letter} is structured as follows. In \\S\\ref{s:data} we describe the data; in \\S\\ref{rela:se} we plot $L_{\\rm X}$ vs $L_{\\rm B}\\sigma^2$ for Coma galaxies illustrating the main result of this letter; in \\S\\ref{s:th} we attempt an interpretation of the observational results by estimating the effect of intracluster gas on X-ray luminosity; we conclude with \\S\\ref{s:con}.  We use $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$ throughout. ",
        "conclusions": "\\label{s:con}  In this paper we have studied a sample of early type galaxies from Coma cluster. We have shown that the diffuse X-ray emission of such galaxies when plotted against the blue band stellar luminosity times the stellar central velocity dispersion, follows a flatter curve than similar galaxies in the field do (Matsushita 2001). Further, we have shown that, starting from the $L_{\\rm x}$ vs $L_{\\rm B} \\sigma^2$ relation measured by Matsushita (2001), we can recover the flatter relation for the galaxies in Coma cluster by accounting for the effects of adiabatic compression from the hot ICM.  $L_{\\rm B} \\sigma^2$ represent the maximum energy input associated with stellar mass loss (Matsushita 2001, Ciotti et al. 1991).  According to our calculations the line \\mbox{$L_{\\rm X} =3.2\\times 10^{40} \\; {\\rm ergs}\\; {\\rm s}^{-1}\\; \\alpha_{\\rm x} \\;{ L_{\\rm B} \\sigma^2 \\over 10^{15}\\; L_\\odot \\; ({\\rm km} \\; {\\rm s}^{-1})^2}$} in the $L_{\\rm X}-L_{\\rm B}$ diagram, where $\\alpha_{\\rm x} \\simeq 0.4$ is the ratio of 0.5--2 keV X-ray to bolometric emissivity, corresponds to the maximum X-ray emissivity if indeed stellar mass loss is the energy source that feeds the diffuse X-ray emission. In fact, this line roughly draws an upper envelope for the compact sources in Matsushita's (2001) sample. Interestingly, Coma galaxies reported in this letter tend to lay on such curve. This is consistent with the above idea about the source of energy for the X-ray emission. In this respect, it is worth pointing out that numerous Coma galaxies in our sample can be characterized as having a ``ultracompact'' X-ray structure and, therefore, according to Matsushita, should be underluminous. Despite this, they also clump near the same theoretical curve based on the assumption that the heating of the galaxies ISM is supplied by the stellar mass loss.  It is worth pointing out that in deriving the X-ray scaling relations for the early-type galaxies, we have not included the contribution of SN Ia to gas heating (Ciotti et al. 1991).  However, according to recent stellar age determinations, field early-types have a systematically younger age than the cluster population (e.g. Terlevich, Forbes 2002), which implies a stronger SN Ia activity there (Greggio \\& Renzini 1983). Thus, if SN Ia played a dominant role in determining the X-ray luminosity of early-types, we would expect cluster galaxies to exhibit lower luminosities, which is not observed.  There are theoretical arguments and observational evidence supporting the idea that ram pressure stripping is one of the principal mechanisms affecting the morphology of X-ray galaxies (e.g. Vollmer et al. 2001, Bravo-Alfaro et al. 2001, Finoguenov et al. 2004a; see also Finoguenov et al. 2004c).  Gas stripping, however, implies a reduction of gas content of the cluster galaxies which, in absence of other effects, leads to a decrease of the level of X-ray emission. We have demonstrated that in fact X-ray emission in Coma cluster galaxies is on average higher than in their field counterparts with similar stellar properties. This is sensibly explained by ICM compression as outlined in the previous Section, which, therefore, adds to the wide variety of processes that affect the evolution of X-ray galaxies.  Finally, as Coma galaxies are more luminous than those in the field because of the high outer pressure, some galaxies in the field may likewise be underluminous, because they expanded when the energy input from mass loss is sufficiently high in comparison to the depth of their potential well.  It will be interesting to further investigate this possibility in the future especially in connection with the new developments in the study of low-sigma early-type galaxies (Burkert \\& Naab 2003)."
    },
    "0403/astro-ph0403288_arXiv.txt": {
        "abstract": "{We have carried out a detailed analysis of the ${\\it IUE}$ archival high resolution spectra of the classical nova V1974 Cyg 1992. The main UV resonance lines show P Cygni profiles in the first days, which change into symmetric pure emission lines, and then slowly become fainter and narrower. Lines of higher ionization species reach their peak luminosity later than those of low ionization.  This can be explained by a fast wind which is optically thick in the early days, when the pseudo-photosphere  is located inside the wind. As the mass loss decreases, the radius of the pseudo-photosphere schrinks. This has three effects that explain the observed changes: (1) the deeper accelerating layers of the wind become visible where the emission lines are formed by collisional excitation and/or recombination, (2) as the mass loss rate decreases the emission comes from deeper regions of the wind where the velocities are smaller, (3) the effective temperature and the degree of ionization increase.  In addition to the P Cygni and emission lines, we could identify two shortward shifted absorption systems which originate in two separate expanding shells, outside the wind layers where the emission lines are formed. The velocity of both shells increase with time.  The outer main shell, containing most of the matter ejected at the outburst, produces the so-called ``principal absorption line system'', and the inner faster moving second shell produces the so-called ``diffuse--enhanced absorption line system''.  The acceleration of the two shells is the result of increasing line-radiation pressure due to the UV-brightening of the star as the effective radius decreases.  Around day 60 the second shell has overtaken the slower moving principal system shell, and merged with it. This explains: the sudden disappearance of the diffuse line system near that date, the upward jump of ${\\Delta}v$ = 240 km~s$^{-1}$ in velocity of the principal system and the first detection of hard X--ray emission on day 63. This velocity jump indicates that the main shell is ${\\approx}$ 4 times more massive than the second shell. The deceleration suffered by the diffuse--enhanced system after the shock provides a shock temperature $T_{shock}$ $\\approx$ 1.6~ keV, in fairly good agreement with the temperature of the observed hard X--ray emission.  The UV observations are interpreted through an empirical model in which the pre--nova slow wind phase is followed by the ejection of two shells, where the principal and the diffuse--enhanced absorption systems are formed, and by a phase of fast continuous lower density wind.  Our empirical expansion velocity law for the principal system, together with H${\\alpha}$ interferometric observations of the angular radius on day 10 are used to determine the distance to the nova, which is found to be 2.9 $\\pm$ 0.2 kpc, in agreement with {\\it HST} imaging and with the absolute magnitude versus rate of decline relationship. ",
        "introduction": "\\label{sec:intro}  V1974 Cyg (Nova Cygni 1992) was discovered in outburst by Collins (1992) on 1992 February 19.07 UT (J.D. = 2448672.3) at a visual magnitude $V$ = 6.8.  The {\\it discovery time} will be taken as day zero all through this paper.  The true time of the outburst of V1974 Cyg can be roughly estimated by comparing its visual light curve, characterized by an optical decay time $t_3$ = 42 $\\pm$ 2 days (Chochol et al. 1993), with that of two novae having a very similar value of $t_3$ of about 40 days: N Lac 1910.9 and N Mon 1918.0. The rise time from quiescence to maximum is ${\\approx}$ 8 days for both objects (cf. the light curves in Cecchini \\& Gratton 1942).  We can then estimate, by analogy, that the outburst of V1974 Cyg took place about 8 days before optical maximum, i.e. about 4.6 days before the reported discovery.  An independent guess of the effective date of outburst can be made by extrapolating the pre-maximum light curve of V1974 Cyg to the magnitude at quiescence, as quoted in Andrillat \\& Houziaux (1993). We obtain in this case that the outburst took place at least 2 days before the time of discovery.  By averaging the two estimates, we obtain that the outburst has probably taken place about 3.3 days before the discovery time (our adopted day zero).  V1974 Cyg has immediately caught the attention of observers because of its exceptional brightness at maximum ($V_{max} \\approx$ +4.4 visual magnitude) reached on February 22.5 (Rosino et al. 1996) and consequently, as being an easy target for multiwavelength observations, which were soon extended to the ultraviolet, radio, infrared and X--rays. These abundant data represent a unique opportunity to get insight into important questions concerning the nova phenomenon.  In this paper we study the dynamics of the expanding envelope of V1974 Cyg through the analysis of the velocity variations of the ultraviolet absorption lines.  Our main purpose is to get insight on the velocity law of the absorption systems, on their origin and possible interaction, and on the velocity and ionization stratification within the envelope.  In general the optical spectra of classical novae show four systems of shortward shifted absorption lines (McLaughlin 1960). In order of appearance, of increasing radial velocity and degree of ionization these systems are: the pre--maximum, the principal, the diffuse--enhanced and the Orion systems.  (For a critical review of the controversial interpretation of the absorption line systems in novae see Friedjung \\& Duerbeck (1993)). In this paper we will adopt this classification scheme. ",
        "conclusions": "\\label{sec:discussion}  {The present investigation, based on high resolution ultraviolet spectroscopy of V1974 Cyg has led  to an empirical model in which the object goes through different mass loss phases. The major results of this study are:}\\\\  (1) In the earliest days the  high velocity wind is optically thick so that the pseudo--photosphere lies within the wind itself. The P Cygni profiles observed in the strong UV resonance lines are formed in this wind. The absorption components of the P Cygni profiles do not start at zero Doppler velocity, because the lower velocity or acceleration layers of the wind are located below the pseudo-photosphere. As the mass loss rate of the wind decreases, the pseudo-photosphere moves inwards, and the deeper higher density wind layers become visible, where emission components are formed by recombination and/or collisional excitation.  The receding pseudo-photosphere also results in a decrease of the effective radius, which causes an increase of the effective temperature.  This explains the increasing degree of ionization, reflected in the observed strengthening of the higher excitation emission lines at the same time that the lower excitation lines become fainter. As the mass loss rate and the wind density decrease further, the emission comes from deeper layers in the wind, where the velocity is smaller, which results in the observed narrowing of the emission lines (see Williams 1992).   (2) Superimposed on the emission lines, one observes two absorption line systems, which are formed in two different shells. The main shell is ejected first and the second shell is ejected later, around day 3 to 5 (see spectra of Andrillat \\& Houziaux, 1993). The acceleration of the two shells is remarkably similar, which suggests a common mechanism.  We propose that the shells are accelerated by radiation pressure in spectral lines (similar to the winds of hot stars) due to the increasing UV brightness of the star.  (3) Within each of the two shells the higher ionization lines have a slightly higher outflow velocity than lower ionization lines, which indicates  that the velocity of the shells  decreases slightly outwards. This is not due to the gravity of the star, nor to the interaction with material outside the shells, but most likely due to the fact that the velocity slightly increased with time during the ejection of the shells.  (4) The inner shell (second shell), which has a higher velocity than the outer shell (main shell) catches up with the outer shell at about day 50 to 70. This agrees roughly with the interaction time expected, based on the difference in velocity of the two shells (Sect. \\ref{sec:impulsive}).  Also, it explains the rather drastic disappearance of the faster shell system around that date, as well as the sudden increase in velocity of the main shell, and the appearance of hard X-ray flux. From the velocities of the faster secondary shell, of the slower main shell, and the final velocity of the merged shell we estimate that the mass ratio of the main shell to the second shell is  roughly 4/1 (Sect. \\ref{sec:impulsive}).  (5) In Sect. \\ref{sec:interpretation} we estimated the distance to V1974 Cyg by assuming a spherically symmetric expansion for the nova envelope at the time of the H${\\alpha}$ interferometric measurements of Quirrenbach et al. (1993) on day 10, and by combining these with the estimated radius of the principal system at the same date. The value so obtained, 2.9 kpc, is in good agreement with the independent estimates by Quirrenbach et al. (1993), Della Valle \\& Livio (1995) and Paresce et al. (1995).\\\\   Our model is qualitatively in agreement with the optically thick wind models of novae, by Kato \\& Hachisu (1994), first introduced by Friedjung (1966a,b,c), and further treated by Kovetz (1998) and Bath \\& Harkness (1989), who showed that the decreasing mass loss rate of the wind will result in an inward moving photosphere, an increase in the effective temperature and an increase in the degree of ionization in the wind.  Since in these models matter is accelerated deep inside the pseudo--photosphere, it is possible in this way to explain the formation of the peculiar P Cygni profiles seen in the earliest spectra and the progressive narrowing of the emission lines.  Also, because of the predicted inward recession of the pseudo--photosphere and the consequent increase of the effective temperature, it is possible to explain with an unique mechanism the acceleration of the two outward moving shells by radiation pressure in the UV lines.  The stratified structure within the expanding shells as suggested in the present model, is in agreement with the major observational features reported by Friedjung (1966b,c), i.e. the overall decrease of the expansion velocity and of the (Zanstra) temperature with increasing radial distance.   However, our observations and our empirical model require that the nova ejected two shells: one before the discovery at about day -3, and one at about day 3 to 5.  In this way it is possible to explain why the absorption of the principal system lines is superimposed on the P Cygni profiles of the wind observed on day 1.3 and why the principal and diffuse--enhanced system lines are superimposed on the wind emission lines after day 6.  The first shell contains most of the ejected matter and the second shell, ejected with a higher velocity, contains considerably less material (about 1/4 in this case).  We might suggest that the second ejection is the reaction to the main ejection. After the ejection of the main shell, the envelope is out of balance and will try to restore equilibrium, but if this does not succeed immediately, due to some sort of overshooting, a second, less massive shell might be ejected.  We like to point out that a similar effect occurs in the large eruptions of luminous blue variables (LBV), which also show that the main eruption is followed by a second one (Humphreys, Davidson  \\& Smith, 1999).  In the case of LBVs, which are very extended and have a small surface gravity, the main eruption lasts for tens of years, and the second eruption occurs about 60 years after the main one. In the case of eruptions from the surface of white dwarfs, where the gravity is very high, the time scales will be much shorter. It might be interesting to see if dynamical models of nova eruptions show this effect.  According to Friedjung (1987), the principal system shell (here the main shell) is formed by interaction of a wind continuously ejected at a decreasing rate, with material ejected before optical maximum. This model does not explain, however, the presence of the second absorption system with higher velocity, produced inside the main shell. Although Friedjung's interaction model is attractive because wind--shell interaction is likely to occur (see also McLaughlin, 1960; Friedjung \\& Duerbeck 1993), it is not clear how this could lead to {\\it two} shells with strongly different and well separated velocities. On the other hand, as discussed in Sect.  \\ref{sec:impulsive}, it is unlikely that the two absorption systems are due to flows in two directions, e.g. one in the orbital plane of the binary and one in the polar direction.  To conclude, we would like to encourage further studies of radial velocity variations in novae  (poorly pursued nowdays), especially during the first few days, since they represent a powerful means to get insight into the dynamical structure of the nova envelopes and into the origin of the X--ray emission. Such studies should also provide valuable input information for non--LTE synthetic spectra of novae (see e.g. Short et al. 2001)."
    },
    "0403/astro-ph0403131_arXiv.txt": {
        "abstract": "We present results from a JCMT/SCUBA 850~\\micron\\ search for cold dust around eight nearby young stars belonging to the \\bPic\\ ($t\\approx12$~Myr) and the Local Association ($t\\approx50$~Myr) moving groups. Unlike most past sub-mm studies, our sample was chosen solely on the basis of stellar age. Our observations achieve about an order of magnitude greater sensitivity in dust mass compared to previous work in this age range. We detected two of the three M~dwarfs in our sample at 850~\\micron, GJ~182 and GJ~803 ($\\Mstar \\approx 0.5~\\Msun$), with inferred dust masses of only $\\approx$0.01--0.03~\\Mearth. GJ~182 may also possess a 25~\\micron\\ excess, indicative of warm dust in the inner few~AU of its disk. For GJ~803 (AU~Mic; HD~197481), sub-mm mapping finds that the 850~\\micron\\ emission is unresolved.  A non-detection of the CO~3--2 line indicates the system is gas-poor, and the spectral energy distribution suggests the presence of a large inner disk hole ($\\approx$17~AU = 1.7\\arcsec\\ in radius for blackbody grains).  These are possible indications that planets at large separations can form around M~dwarfs within $\\sim$10~Myr.  In a companion paper (Kalas, Liu \\& Matthews 2004), we confirm the existence of a dust disk around GJ~803 using optical coronagraphic imaging.  Given its youthfulness, proximity, and detectability, the GJ~803 disk will be a valuable system for studying disk, and perhaps planet, formation in great detail. Overall, sub-mm measurements of debris disks point to a drop in dust mass by a factor of $\\sim$10$^3$ within the first $\\sim$10~Myr, with the subsequent decline in the masses of sub-mm detected disks consistent with $t^{-0.5}$ to $t^{-1}$. ",
        "introduction": "Circumstellar disks are common around stars at ages of a few million years. Rotationally supported disks are a natural outcome of the star formation process, because of the need to conserve angular momentum during the collapse of the natal molecular core.  In addition, disks are central to current theories of planet formation.  These young disks are composed of gas and dust, with estimated masses as large as several percent of the stellar mass.  By virtue of their large optical depths, they can be easily detected and have been intensively studied.  Over the span of $\\sim$5--10 Myr, these primordial disks are thought to dissipate \\citep[e.g.][]{1993prpl.conf..837S, 2001ApJ...553L.153H}, based on the observed decline of IR excesses, which originate from warm dust in the inner disk.  At older ages, planetesimals and planets comprise the remaining circumstellar population. Collisions between the planetesimals produce dust grains, which form a \"second-generation\" disk.  The {\\sl Infrared Astronomical Satellite} (\\IRAS) discovered the first example of a such dusty debris disk around Vega \\citep{1984ApJ...278L..23A}.  Overall, volume-limited far-IR surveys have found that about 15\\% of main-sequence stars have debris disks \\citep{1993prpl.conf.1253B, 1999ApJ...520..215F, 1999ApJ...510L.131G, 1999A&A...343..496P, 2001A&A...365..545H, 2002A&A...387..285L}.  Such disks have a small amount of dust, little evidence for gas, low characteristic temperatures indicative of grains at tens of~AU separations, and a dearth of warm material in their inner regions \\citep[e.g.,][]{2001ARA&A..39..549Z}.  These attributes make debris disks physically distinct from primordial disks and lead to much weaker observational signatures.  Loss processes for the dust in debris disks act quickly compared to the age of the host stars, and the debris must be continually replenished. Thus, studying debris disks over a wide range of ages can provide insight into their evolution and physical origin.  Most studies have focused on stars with ages of $\\gtrsim$200~Myr \\citep{1999Natur.401..456H, 2001A&A...365..545H, 2001ApJ...555..932S, 2002A&A...387..285L}, a reflection of the fact that most nearby stars are old.  The last few years have seen the discovery of many stars which are both young ($t\\approx10$--100~Myr) and close to Earth ($d\\lesssim100$~pc). These stars have largely been found in kinematically associated moving groups, by combining accurate space motions with evidence for youthfulness (\\eg, Li abundances, X-ray emission, and chromospheric activity).  They include the \\bPic\\ ($\\approx$8--20~Myr; \\citealp{2001ApJ...562L..87Z}) and the Local Association ($\\approx$20--150~Myr; \\citealp{1995MNRAS.273..559J, 2001MNRAS.328...45M}) moving groups.  These are important systems for disk studies as they span an age range which is poorly studied.  In addition, this timescale is associated with assemblage of terrestrial planets in our own Solar System \\citep[e.g.][]{2000prpl.conf..995W}, so understanding disk properties at this epoch can illuminate the overall physical context for planet formation.  Given the cold characteristic temperatures of the dust ($\\approx$40--100~K; \\eg, \\citealp{2000MNRAS.314..702D}), sub-mm continuum observations of debris disks are valuable.  The emission is optically thin and hence can be used to estimate the disk masses, with the caveat that large bodies are missed.  Previous sub-mm studies have focused on stars already known to have strong IR excesses \\citep{1993ApJ...414..793Z, 1996MNRAS.279..915S, 1998Natur.392..788H, 1998ApJ...506L.133G, 2000MNRAS.313...73S, 2001MNRAS.327..133S, 2003ApJ...582.1141H, 2003AJ....125.3334H}, with the most outstanding examples being spatially resolved at sub-mm wavelengths (\\bPic, $\\epsilon$~Eri, Fomalhaut, Vega).  Much less has been done in the way of \"blind\" searches of stars in a given age range to understand their sub-mm disk frequency and properties.  \\citet{2003MNRAS.342..876W} have recently surveyed the secondary stars of binaries with O and B-type primaries, a sample originally identified by \\citet{1986A&A...156..223L}.  By virtue of the primary stars being on the main sequence, the systems are $\\lesssim$100~Myr old. Previous searches for dust around this sample had relatively poor sensitivity \\citep{1994A&A...282..123G, 1994AJ....108..661J, 1995ApJ...440L..89R}. Of particular interest is the suggestion by Wyatt~\\etal\\ that some stars may harbor very cold ($\\lesssim$40~K) disks, which would be missed by previous far-IR searches and only detected in the sub-mm.  However, one limitation of the Lindroos sample is that most of the systems are fairly distant, which impacts the resulting disk sensitivity of sub-mm observations.  In this Letter, we present an 850~\\micron\\ search for cold dust around young stars in the \\bPic\\ and Local Association moving groups.  Unlike most previous sub-mm studies, our sample was chosen solely on the basis of stellar age.  Our sample includes three M~dwarfs (GJ~182, 799 and 803) which \\citet{1999ApJ...520L.123B} identified as the youngest of the nearby M~dwarfs, based on the stars' pre-main sequence location on the color-magnitude diagram.  The majority of debris disks from IR studies have been identified around A, F, and G-type dwarfs, simply due to observational sensitivity bias (\\eg, compilations by \\citealp{2003MNRAS.345.1212G} and \\citealp{2003astro.ph.11546Z}).  The loss processes for circumstellar dust depend on the host star luminosity and gravity.  Therefore, studying debris disks over a range of stellar masses may also advance our understanding of this phenomenon. ",
        "conclusions": "\\subsection{Sub-Millimeter Evolution of Debris Disks}  Figure~\\ref{agetrend} summarizes the known debris disk mass estimates based on sub-mm observations from our work and the published literature. Our SCUBA survey adds a significant number of stars of $\\approx$10--50~Myr.  Our raw 850~\\micron\\ sensitivity (median rms of 1.9~mJy) is poorer than the \\citet{2003MNRAS.342..876W} survey of Lindroos binaries (median rms of 1.6~mJy).  But given the much closer distances of our targets, our dust mass sensitivity is better by around an order of magnitude.  In particular, Figure~\\ref{agetrend} shows the following:  \\begin{enumerate}  \\item Compared to primordial disks around T~Tauri and Herbig~Ae/Be stars which have dust masses of $\\approx$10--300~\\Mearth\\ \\citep[e.g.][]{1995ApJ...439..288O, 2000prpl.conf..559N}, the dust masses for the detected young stars are $\\sim$10$^3$ smaller.  This points to very rapid evolution of the circumstellar dust mass within the first $\\sim$10~Myr, likely arising from grain removal (\\eg, via accretion onto the central star) and/or grain growth into larger bodies, which would greatly diminish the sub-mm/mm emission.  \\item The upper envelope for disk masses around A~stars exceeds that for later-type stars.  This may reflect the distribution of primordial disk masses; for T~Tauri stars a similar trend is seen, where the upper envelope of primordial disk masses is larger for more massive stars \\citep{2003astro.ph..4184N}.  \\item For stars of $\\approx10-100$~Myr old, nearly all the detected disks at a given spectral type (\\ie, stellar mass) are less massive than the upper limits on non-detections.  This highlights the fact that existing sub-mm measurements are only sensitive to the most massive of the young debris disks.  \\item The apparent correlation of dust mass with stellar age arises from observational bias: the young stars detected in the sub-mm are at larger distances than the old stars.  However, the decline in the upper envelope of disk masses with age is probably a real effect.  While there are no published sub-mm data in the upper right of Figure~\\ref{agetrend} (\\ie, old stars with massive disks), the existing sub-mm detections of nearby, old stars come from sources known to possess large IR excesses in \\IRAS\\ and \\ISO\\ data.  Any old stars with even more massive disks would have been detected by these IR surveys, unless the dust was unprecedentedly cold.  Based on the data in Figure~\\ref{agetrend}, the masses of sub-mm detected disks evolve roughly as $M_{dust}\\sim t^{-0.5}\\ {\\rm to}\\ t^{-1}$, with a simple unweighted fit giving $M_{dust} \\propto t^{-(0.7\\pm0.2)}$.    \\end{enumerate}     \\subsection{Debris Disks and Planets around Low-Mass Stars}  GJ~182 and GJ~803 are the first M-star debris disks detected at sub-millimeter wavelengths.\\footnote{The young M3~dwarf Hen 3-600 also shows an IR excess \\citep{1989ApJ...343L..61D}.  This binary star is a member of the $\\approx$10~Myr old TW Hydrae Association \\citep{2001ApJ...549L.233Z}.  Hen~3-600 has a very large dust optical depth, a significant 10--20~\\micron\\ excess \\citep{1999ApJ...520L..41J}, and active gas accretion \\citep{2000ApJ...545L.141M}.  These characteristics are distinct from debris disks and indicate that the Hen~3-600 disk is primordial or else in a transitional state.}  The dust lifetime is short compared to the stars' ages \\citep[e.g.][]{1993prpl.conf.1253B}.  Therefore, the dust must be replenished by collisions between planetesimals, either in situ or else from regions farther out.  Note that for these low luminosity stars (0.1--0.2~\\Lsun), radiation pressure on the grains is negligible compared to the star's gravity, unlike debris disks around A~stars, \\eg, \\bPic\\ where micron-sized and smaller grains are expelled \\citep{1988ApJ...335L..79A}.  Our work shows that most debris disks around low-mass stars probably lie at or below the sensitivity limits of the \\IRAS\\ survey and JCMT/SCUBA observations.  GJ~182 is just at the practical limit for \\IRAS\\ and SCUBA detection.  GJ~803 is the nearest object in our sample, and among the nearest known young stars, which makes detection of its very tenuous dust mass feasible.  The one undetected M~dwarf in our sample, GJ~799, is just as close and young as GJ~803.  GJ~799 is known to be a 3\\arcsec\\ binary of comparable magnitude \\citep{2001AJ....122.3466M}, which may impact its dust content.  In addition, its M4.5 spectral type corresponds to $\\Mstar\\approx$0.10--0.15~\\Msun, about 3--4$\\times$ smaller than for GJ~182 and GJ~803 (\\eg, see models presented in \\citealp{1997AJ....113.1733H} and \\citealp{1999ApJ...525..466L}). Hence, the non-detection may simply be due to insufficient sensitivity; our JCMT data for GJ~799 would have to be $\\approx$4$\\times$ deeper in order to reach the same $M_{dust}/M_{star}$ as for the two detected M~dwarfs.  Primordial disks are common around low mass stars \\citep{2001AJ....121.2065H} and even brown dwarfs \\citep{2002astro.ph.10523L} at ages of a few~Myr, and these systems may very well generate debris disks as they age.  The {\\sl Spitzer Space Telescope} will offer the requisite sensitivity for detecting and studying such disks.  For GJ~182, the 25~\\micron\\ excess suggests that dust resides in the inner regions of its disk: 150--200~K blackbody grains would lie at 1--2~AU from the star.  Given the sparse sampling of the current SED, one cannot determine if the dust distribution is continuous throughout the disk, or if there is a gap in the middle regions of the disk (\\eg, from planet formation) between the 25~\\micron\\ and 850~\\micron\\ emitting dust.  However, as shown by \\citet{2002AJ....124..514S}, single band excesses near the sensitivity limit of the \\IRAS\\ catalog may not be reliable.  Though the SCANPI detection of GJ~182 has S/N~=~4.6, higher than the S/N~=~2--4 objects examined by Song \\etal, for now we treat the 25~\\micron\\ excess of GJ~182 as tentative.  More detailed SED measurements, in particular with \\Spitzer, would be invaluable to characterize this star's disk.   The absence of warm dust in the inner regions of GJ~803 (and perhaps GJ~182, if the 25~\\micron\\ detection is spurious) is naturally explained by an unseen inner companion.  The 40~K temperature characteristic of the SED is too cold to be explained by sublimation or melting of icy grains ($\\approx$100-170~K).  Without an inner companion, the effect of Poynting-Robertson drag would cause grains to spiral inward and produce significant mid-IR emission.  Such inner holes are typical for debris disk systems.  Indeed, the few debris disks with resolved dust emission show morphologies highly suggestive of the dynamical influence of an unseen companion at mean-motion resonances \\citep{1998Natur.392..788H, 1998ApJ...506L.133G, 2000ApJ...537L.147O, 2001ApJ...560L.181K, 2002ApJ...569L.115W, 2002ApJ...578L.149Q, 2003ApJ...584L..27W, 2003ApJ...588.1110K, 2003ApJ...597..566T}.  The 40~K characteristic temperature from the SED fitting means that blackbody grains around GJ~803 would be at a distance of 17~AU, or an impressively large angular separation of 1.7\\arcsec.  (For GJ~182, 17~AU corresponds to only 0.6\\arcsec.)  This provides only a rough guide as to the angular separation of any interior companion, since it is difficult to use the SEDs of debris disks to predict the spatial dust distribution \\citep{2003astro.ph.11593S}. Shallow IR adaptive optics imaging from the Keck Observatory with 0.05\\arcsec\\ resolution shows no stellar or massive brown dwarf companion, suggesting that an even lower mass companion is present.  To date, GJ~876 is the only M~dwarf known to have extrasolar planets \\citep{1998ApJ...505L.147M, 2001ApJ...556..296M}; its two planets have semi-major axes of 0.13 and 0.21~AU.  Thus the GJ~803 inner disk hole may indicate that planets can exist at larger separations around M~dwarfs than known so far and that such planets can form within $\\approx$10~Myr. Furthermore, given the very young age of the system, a planetary companion is predicted to have significant thermal emission \\citep[e.g.][]{1997ApJ...491..856B} and hence might be detectable with deeper AO imaging.  M~dwarfs are by far the most numerous type of star, so a better understanding of their planet-bearing frequency is needed for a complete census of the extrasolar planet population.  Follow-up coronagraphic imaging by \\citet{klm03} finds that GJ~803 has a large, nearly edge-on disk seen in scattered optical light.  This disk is detected from 5--21\\arcsec\\ (50--210~AU), though the inner and outer extent are not well-constrained.  Our sub-mm imaging finds that the 850~\\micron\\ dust emission is unresolved.  This is entirely compatible with the large optical disk, given that the grains responsible for scattering may be very cold and/or small, and hence poor sub-mm emitters.  Deeper sub-mm mapping is needed for stronger constraints on any extended emission. In general, the proximity and youthfulness of GJ~803, combined with the high sub-mm and optical detectability of its disk, make this system an excellent opportunity to study disk evolution, and perhaps planet formation, in great detail.   The onset and duration of the debris disk phenomenon remain open questions.  As discussed by \\citet{2003MNRAS.345.1212G}, the disk masses might slowly decay over time, or perhaps the mass distribution is bimodal, reflecting ``on'' and ``off'' states for detectability.  The GJ~803 debris disk offers some insight into this issue. Using the {\\sl Hipparcos} catalog, which is nearly complete for early-type M~dwarfs (M0--M2) within 10~pc, \\citet{2002AJ....124..514S} found GJ~803 was the only object out of a sample of 152~M~dwarfs to possess an IR excess in the \\IRAS\\ FSC. This is not due to detection bias, as there are many early-type M~dwarfs which have closer distances than GJ~803. GJ~803 is also known to be among the youngest of the nearby M~dwarfs \\citep{1999ApJ...520L.123B, 2001ApJ...562L..87Z}. If the appearance of debris disks was an intermittent event, Song \\etal\\ should have detected older M~dwarfs with IR excesses at both 60~and 100~\\micron, if their disks were like the GJ~803 one. Either these older M~dwarfs do not have disks, or their disks are much less massive than GJ~803. Therefore, at least for the case of M~dwarfs, the GJ~803 system suggests that the time evolution of debris disks is driven primarily by age, with the youngest stars having the largest disk masses."
    },
    "0403/astro-ph0403307_arXiv.txt": {
        "abstract": "{It is generally believed that major galaxy mergers with mass ratios in the range 1:1--3:1 result in remnants that have properties similar to elliptical galaxies, and minor mergers below 10:1 result in disturbed spiral galaxies. The intermediate range of mass ratios 4:1--10:1 has not been studied so far. Using N-body simulations, we show that such mergers can result in very peculiar systems, that have the morphology of a disk galaxy with an exponential profile, but whose kinematics is closer to that of elliptical systems. These objects are similar to those recently observed by Jog \\& Chitre (2002). We present two cases with mass ratios 4.5:1 and 7:1, and show that the merging causes major heating and results in the appearance of elliptical-type kinematics, while surprisingly the initial spiral-like mass profile is conserved. ",
        "introduction": "Numerical simulations commonly show that the merging of two equal-mass spiral galaxies results in the formation of an elliptical galaxy with an $r^{1/4}$ radial profile consistent with observations (de Vaucouleurs 1953), and which is mainly supported by velocity dispersion rather than rotation (e.g., Barnes 1992). Galaxy mergers with mass ratios of 1:1 to 3:1 and 4:1 have been studied by Bendo \\& Barnes (2000), Cretton et al. (2001) and Naab \\& Burkert (2003): these lead to the formation of  boxy or disky elliptical galaxies, depending on the mass ratio. On the other hand, minor mergers with mass ratios of more than 10:1 have also been studied in numerical simulations. They correspond to the merging of a spiral galaxy with a dwarf companion, which has visible effects on the morphology and kinematics of the remnant: Walker et al. (1996) have found a significant heating and thickening of the disk. However, the remnant of a minor merger remains an exponential disk supported by rotation.  These numerical works do not account for a recent observational result: Chitre \\& Jog (2002) have analyzed a set of observed advanced merger remnants, and have classified them in two morphological classes. ''Class~I\" remnants are elliptical-type galaxies with an $r^{1/4}$ radial profile. ''Class~II\" systems have an exponential profile typical of spiral disks (Freeman 1970) ; a central bulge is present, too, but does not dominate the mass distribution, so that these systems cannot be simply regarded as elliptical galaxies with faint outer disks. Among the Class~II objects, Jog \\& Chitre (2002) have pointed out systems with kinematical properties typical of elliptical galaxies: they have velocity dispersions as large or larger than rotation velocity, while spiral disks are usually supported by rotation. We will call these merger remnants with spiral-like morphology and elliptical-like kinematics ''hybrid\" systems.  In this Letter, we test the hypothesis of Jog \\& Chitre (2002), that the unexplored range of mass ratios of 4:1 to 10:1 could lead to the formation of these hybrid systems in galaxy mergers. We present N-body simulations of galactic encounters with mass ratios 4.5:1 and 7:1, and show that the remnants reproduce the mixed properties observed in the hybrid systems. Mergers with such unequal-mass ratios are very likely to occur, since low mass galaxies are much more numerous than large-mass ones (e.g., Binney \\& Tremaine 1987). ",
        "conclusions": "We have shown that galaxy mergers in the unexplored range of mass ratios 4:1--10:1 can result in the formation of hybrid systems, that have the mass distribution as in spiral galaxies and yet have kinematical properties closer to elliptical galaxies than spiral galaxies. They correspond to the merger remnants observed recently by Jog \\& Chitre (2002). We have thus quantitatively proven their conjecture that for this mass range, the merger causes major heating and yet leaves the mass distribution exponential.  Whether such hybrid systems are formed frequently should be studied with a large set of simulations, using various mass ratios and orbital parameters. This would more generally indicate the parameter range for which a galaxy merger results in an elliptical galaxy or a disk galaxy. This, and the evolution of hybrid remnants under successive galaxy mergers or external gas accretion, will be studied in future papers."
    },
    "0403/astro-ph0403182_arXiv.txt": {
        "abstract": "We present new broad-band optical and near-infrared CCD imaging together with deep optical intermediate-resolution  spectroscopy of the Hickson Compact Group 31. We analyze the morphology and colors of the stellar populations of the galaxies, as well as the kinematics, physical conditions and chemical composition of the ionized gas in order to get a  more complete view on the origin and evolution of the system. We estimate the ages of the most recent star formation  bursts of the system, finding an excellent consistency among the values obtained with different indicators and starburst models. We find that member F hosts the youngest starburst of the group, showing a substantial population of Wolf-Rayet stars. The chemical abundances are fairly similar in all the members of the group despite their very different absolute magnitudes. We argue that the use of traditional metallicity-luminosity relations based on the absolute $B$-magnitude is not appropriate for dwarf starburst galaxies, because their luminosity is dominated by the  transient contribution of the starburst to the blue luminosity. We think that members E and F of the group are candidate tidal dwarf galaxies because of their high metallicity, their kinematics, and the absence of underlying old  stellar populations. Finally, we propose that HCG~31 is suffering several almost simultaneous interaction processes. The most relevant of these processes are: (a) the merging of members A and C, that would have produced two optical tidal tails; and (b) a fly-by encounter between G and the A+C complex, that would have produced an \\ion{H}{1} tidal tail from the stripping of the external gas of A+C, from which members F and E have originated. ",
        "introduction": "\\footnotetext{Based on observations made with several telescopes operated on the islands of La Palma and Tenerife by  the Isaac Newton Group of Telescopes, Nordic Optical Telescope and Instituto de Astrof\\'\\i sica de Canarias in the Spanish observatories of Roque de Los Muchachos and Teide of the Instituto de Astrof\\'\\i sica de Canarias.}  The starburst phenomenon in galaxies was first described by \\citet{SS70}, who noted that some galaxies seem to be experiencing strong episodes of star formation. \\citeauthor{SS70} estimated that the material available for the production of stars would be exhausted in a very short time compared to the age of the Universe. Since this  discovery, many studies have been performed trying to understand the processes that trigger starburst galaxies. Wolf-Rayet (WR) galaxies are a subset of starbursts galaxies that show broad stellar \\ion{He}{2} $\\lambda$ 4686  emission in  their integrated spectra due to WR stars, whose ages are less than 6 Myr. This spectral feature indicates the presence  of a substantial population of this sort of massive stars and offers the opportunity to study very young bursts.  Making use of population synthesis models it is possible to determine the age of the WR bursts and, consequently, the  star formation and its causes can be studied.  In dwarf galaxies, starburst phenomena cannot be understood under the wave-density theory because of their low masses,  so other mechanisms are needed. One of the proposed alternative mechanisms for large scale starburst formation is  the gas compression by shocks due to the mass lost by galactic winds, followed by the subsequent cooling of the medium. Recent  results of this theory \\citep{Hi00} predict an intermittent behavior of the bursts in dwarf galaxies, with long inactivity periods. This theory was first suggested by \\citet{Th91}. Others authors have proposed galaxy interactions  as the massive star formation triggering mechanism \\citep{S88}, although the interactions cannot be with neighbor  giant galaxies in most cases \\citep*{CA93,TT95}. In fact, they seem to be more frequent with low surface brightness  galaxies \\citep*{T93,T95,T96}. This led \\citet{ME99} to suggest that interactions with other dwarf objects could  be the main star formation triggering mechanism in dwarf galaxies and, especially in WR galaxies.  Compact groups of galaxies are systems with a high density of galaxies and a low dispersion velocity, constituting the  highest density cusps in the non-clustered large-scale structure. They were defined by Hickson (1982) using various  richness, compactness and isolation criteria. \\citet{H92} and \\citet{B96} extended the compact group  definition. The resultant compact group population represents $\\sim$1-2\\% of the galaxy field and is characterized by aggregates of 4-8 galaxies with mean projected separations of the order of a component galaxy diameter. The small  projected separations and observed signs of interaction imply physical densities as high as those found in the cores of rich clusters of galaxies. In fact, interactions between the group-members are important in some compact groups. \\citet{Mo94} found that star formation rates (SFRs) in galaxies belonging to compact groups are somewhat higher than in field galaxies. \\citet{MOH94} reported morphological interaction signs, like mergers and tidal  tails, in many galaxies from \\citet{H82} catalogue of compact groups. Recent studies of interactions in compact  groups can be found in \\citet*{VIP98, IPV99, IPV01, VM97, VM98, VM01, VM02, SRV01, PAM02}. The finding of the WR characteristics in some of the members of these compact groups allows the study of star formation and its triggering mechanism with detail. Perhaps, one of the best systems to perform this investigation is the WR galaxy NGC 1741 and the compact group where it belongs: HCG~31.   \\subsection{Brief history of the membership}   The compact group HCG~31, at a distance of 54.8 Mpc (H$_0$=75 km s$^{-1}$ Mpc$^{-1}$; \\citealt{VC92}), was  identified by \\citet{H82}. It is one of the best studied compact groups because of the peculiar morphology of its  members that includes tidal tails, a possible merger, irregular structures, and prominent starbursts in its  brightest  galaxies (see Fig.~\\ref{fig1}). \\citet{H82} described the group as four galaxies in close proximity, that he called A,  B, C and D. He noted that two of them, members A and C, are clearly interacting, and they both actually form NGC 1741  (Mrk 1089, VV524, Arp259). In fact, \\citet{MB93} also classified NGC 1741 as a double nucleus Markarian galaxy.  \\citet{KS86} detected the 4686 \\AA\\ WR feature in the spectrum of the nucleus of galaxy C. This was later  confirmed by \\citet*{R90}. NGC 1741 was included in the first catalogue of WR galaxies by \\citet{C91} as one  of the most luminous WR galaxies known. \\citet{VC92} obtained the optical spectra of A and C and found that they are dominated by strong nebular emission lines, as expected from the presence of a large number of hot, massive  stars. UV studies performed by \\citet*{C96} revealed that the centre of NGC 1741 is dominated by two main starbursts, which are composed of several intense knots of recent star formation. \\citet{J99} obtained {\\it Hubble Space Telescope} ($HST$) images of the system, dating the main starbursts (A and C) in 5 Myr,  and detecting 434 Super Star Clusters (SSCs) in the central knots. \\citet{R03} called NGC 1741 the A+C complex, because they considered that A and C are a single kinematical entity.  Member B was classified as a Sm galaxy by \\citet{H82}. It shows three different bright knots in H$\\alpha$  images (\\citealt{R90, IPV97}, hereinafter IV97; \\citealt{JC00}). Its morphology suggests that it is a dwarf spiral or irregular galaxy seen nearly edge-on. \\citet{R90} and \\citet{R03} found that it shows solid-body rotation and seems to  be kinematically distinct from the A+C complex.  \\citet{R90} identified two nearby additional objects, E and F, and the close galaxy Mrk 1090 (object G), that were  all included as new members of HCG~31, as well as the far member Q, that lies 2\\arcmin\\ to the north of the A+C complex.  These  authors also determined that D is a background galaxy at a distance of 359 Mpc. IV97 studied the ages of the bursts,  and concluded that F is as young as member C. The \\ion{H}{1} map of the group \\citep*{W91} shows that  all the galaxies, except D, are embedded in the same neutral cloud. \\citeauthor{W91} estimated a total hydrogen mass of 2.1  $\\times$ 10$^{10}$ M$_\\odot$. Maxima of \\ion{H}{1} column density are coincident with the galaxies, indicating  that they are gas rich.  \\citet*{HCZ96} proposed several tidal dwarf candidates in HCG~31: one of them corresponds to member F, and  the other ones are located at the northeast of the A+C complex. \\citet{JC00} considered E and F as tidal dwarf  galaxies, and detected new SSCs in them. Those authors  confirmed that F shows a starburst of around 4 Myr, although  it can be younger. \\citet{IPV01} noted 2 tidal objects at the northeast of the A+C complex and 7 tidal objects (including E, F1, F2, and faint H) at the southeast tidal tail, which appears delineated in the \\ion{H}{1} map. \\citet{IPV01} analyzed  those objects and concluded that only member F satisfies the escape and self-gravitation conditions to be a tidal  dwarf galaxy. \\citet{R03} argued that E is an integral part of the A+C complex or is now separating from it. For those authors, the kinematics of member E confirms that the extension of the A+C complex towards E and F is a tidal tail. \\citeauthor{R03} also suggested that galaxy F could be a tidal fragment of the A+C complex that detached some time ago.  G has a spherical shape in the broad-band images, but IV97 found that it hosts some knots of stellar formation at  the northwest. This was confirmed by \\citet{JC00}, who found a very asymmetric H$\\alpha$ emission, with the star-forming regions forming a U shape along the northwest side. \\citet{R03} described G as a small disk galaxy  seen nearly face-on because the H$\\alpha$ velocity has only a very small gradient across its disk. They also noted  that its kinematics is distinct from that of the A+C complex.  \\subsection{Structure of this paper}  In \\S~2 we present our observations and the data reduction processes. In \\S~3 we present our optical and near-infrared photometric results, studying the effect of the reddening in the final data.  In that section, we also present the results of our intermediate-resolution spectroscopy: physical conditions,  chemical abundances, and kinematics of the ionized gas. In \\S~4 we discuss our results: we analyze the ages of the bursts and the stellar and WR populations, the star formation rates, the luminosity-metallicity relation in the members of HCG~31, the possibility that members E and F are tidal dwarf galaxies, and a revision of the star formation history in HCG~31. We finally present our main conclusions in \\S~5. ",
        "conclusions": "In this paper we present results based on new broad-band optical and near-infrared imagery and optical spectroscopy  of the compact group HCG~31. Our main results can be summarized as follows: \\begin{itemize} \\item We have obtained remarkably consistent age estimations for the most recent star formation bursts of the  group using different observational indicators and theoretical models. Member F hosts the youngest starburst of the group (age between 2 and 3 Myr). It does not show evidences of an underlying old stellar population and hosts a substantial  population of WR stars. \\item We have obtained direct determinations of the oxygen abundance for most of the galaxies of the group,  finding very similar values in all the cases (12+log(O/H) = 8.0--8.2) despite their very different  absolute magnitudes. We report the first direct abundance determination for members B, E, F2, and G. \\item We detect the \\ion{C}{2} $\\lambda$4267 line in the spectrum of member C. This is the first time this  recombination line is reported in an \\ion{H}{2} galaxy. The C$^{++}$/H$^+$ ratio we derive from this line is  too high, but this could be due to our possible overestimation of the intensity of this weak line. \\item We argue that the use of traditional metallicity-luminosity relations based on the absolute  $B$-magnitude (as those of \\citealt{RM95}, constructed for non-bursting dwarf galaxies) is not appropriate for  dwarf starburst galaxies. Their luminosity is dominated by the transient contribution of the starburst to the blue  luminosity, which increases dramatically the $B$-magnitude with respect to the pre- or post-starburst quiescent phase. \\item We propose that members E and F are tidal dwarf galaxies made of gas-rich material stripped  from the A+C complex because of: (a) their high oxygen abundance (similar to that of the other galaxies of the group);  (b) the absence of substantial underlying old stellar populations; (c) their kinematics. \\item We propose that HCG~31 is a very complex interacting system of galaxies that is suffering several  simultaneous interaction processes. A and C are the dominant members of the group and they are now merging.  This merging process could be the origin of the two optical tidal tails that extend towards the northeast and  southwest of the A+C complex. There was a fly-by encounter between G and the A+C complex (or one of these galaxies  before the merging) and this encounter produced an \\ion{H}{1} tidal tail from the stripping of the external gas of  A+C.  This tail moves in or very close to the direction of the plane of the sky.   Members E and F could have formed from instabilities produced in this gas-rich structure. Finally, member B is in  interaction with the A+C complex. This is indicated by: (a) the strong star formation that is taking place along the   tilted disk of B; (b) the presence of apparent streaming motions in the ionized gas between both  objects; (c) the clear faint bridge of matter that connects B and the A+C complex. \\end{itemize}"
    },
    "0403/astro-ph0403527_arXiv.txt": {
        "abstract": "We reiterate that there is no evidence that BL Lacs are sources of ultrahigh energy cosmic rays. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403711_arXiv.txt": {
        "abstract": "s{We study the two-dimensional structure of thin transonic accretion discs in the vicinity of a non-spinning black hole within the framework of hydrodynamical version of the Grad-Shafranov equation. Our analysis focuses on the region inside the marginally stable orbit (\\MSO), $r < \\rms$. We show that all components of the dynamical force in the disc become significant near the sonic surface and (especially) in the supersonic region. Under certain conditions, the disc structure can be far from radial, and we review the affected disc properties, in particular the role of the critical condition at the sonic surface. Finally, we present a simple model aimed at explaining the quasi-periodical oscillations that have been observed in the infra-red and X-ray radiation of the Galactic Centre.% } ",
        "introduction": "\\label{Sec.intro} The investigation of accretion flows near black holes (BHs) is undoubtedly of great astrophysical interest. Substantial energy release must take place near BHs, and general relativity effects, attributable to strong gravitational fields, must show up there. Depending on external conditions, both quasi-spherical and disc accretion flows can be realized. The structure of thin accretion discs has been the subject of many papers. Many results were included in textbooks~\\cite{sha83,lip92}. Lynden-Bell \\cite{lyn69} was the first to point out that supermassive BHs surrounded by accretion discs could exist in galactic nuclei. Subsequently, a theory for such discs was developed that is now called the standard model, or the model of the $\\alpha-$disc~\\cite{sha72,sha73,nov73}.  Since then the standard disc thickness prescription has been widely used, \\begin{equation} H \\approx r\\frac{\\cs}{v_{\\rm K}}, \\label{2} \\end{equation} where the disc thickness $H$ is assumed to be determined by the balance of gravitational and accreting matter pressure forces with the dynamical force neglected. This relation was later used in the renowned approach where all quantities were averaged over the disc thickness~\\cite{pac81}, with a lot of such one-dimensional models following~\\cite{abr88,pap94,rif95,che97,nar97,pei97,bel98,gam98a,gam98b,art01}. As for the two-dimensional structure of accretion discs, it was investigated mostly only nume\\-rical\\-ly and on\\-ly for thick discs~\\cite{pap94,igu97,bal98,kro02}.  Even though standard disc thickness prescription \\eqref{2} and the averaging procedure are likely to be valid in the region of stable orbits \\hbox{$r > \\rms$}~\\cite{sha83}, they in our view require a more serious analysis. It is the assumption that the transverse velocity $v_\\theta$ may \\emph{always} be neglected in thin accretion discs all the way up to horizon~\\cite{abr97}, \\ie\\ that the disc thickness is always determined by (\\ref{2}), that is the most debatable~\\cite{bes04}. This assumption is widely used, explicitly or implicitly, virtually in all papers devoted to thin accretion discs~\\cite{abr81,cha96}.  There is a brief discussion on thin discs in this volume that covers theory shortcomings and a brief description of our approach (\\cf\\ Sec.\\ named Thin disk in~\\cite{bes04}). In this paper we briefly describe our study of subsonic and transonic regions of thin discs followed in Sec.\\hbox{}~\\ref{Sec.SupersonicFlow} by the elaborate discussion of the supersonic flow. Finally, in Sec.\\hbox{}~\\ref{Section.Observ} we develop a toy model for explaining the observed quasi-periodical oscillations detected in the infra-red and X-ray observations of the GC.  \\section[]{Basic equations} \\label{Sec.BasicEqs} We consider thin disc accretion on to a BH in the region where there are no stable circular orbits. The contribution of viscosity should no longer be significant here~\\cite{bes04}. Hence we may assume that an ideal hydrodynamics approach is suitable well enough for describing the flow structure in this inner area of the accretion disc. Below, unless specifically stated, we consider the case of non-spinning BH, \\ie\\ use the Schwarzschild metric, and use a system of units with $c=G=1.$ We measure radial distances in the units of $M$, the BH mass.  In \\hbox{Boyer-Lindquist} coordinates the Schwar\\-z\\-schild metric is~\\cite{lan87a} \\begin{equation} {\\rm d}s^{2}=-\\alpha^{2}{\\rm d}t^{2} + g_{ik}{\\rm d}x^{i}{\\rm  d}x^{k}, \\label{a1} \\end{equation} where \\begin{gather} \\alpha^2 = 1-2/r, \\quad g_{rr} = \\alpha^{-2}, \\quad g_{\\theta \\theta} = r^2, \\quad g_{\\varphi \\varphi} = \\varpi^{2} = r^2\\sin^2\\theta.\\label{a2} \\end{gather}  We reduce our discussion to the case of axisymmetric stationary flows. For an ideal flow there are three integrals of motion conserved along the streamlines, namely entropy, $S$, energy  $E = \\mu\\alpha\\gamma$, and $z$-component of angular momentum $L = \\mu\\varpi u_{\\hat\\varphi}$, where $\\mu=(\\rho_{m}+P)/n$ ($\\rho_{m}$ is internal energy density, $P = n T$ is pressure) is relativistic enthalpy. The relativistic Bernoulli equation $\\up[2] = \\gamma^2-u_{\\hat\\varphi}^2 - 1$, where \\up\\ is the physical poloidal $4$-velocity component~\\cite{bes04,bes97,bes02}, now becomes \\begin{align} \\up[2] &= \\frac{E^2-\\alpha^2L^2/\\varpi^2 -\\alpha^2\\mu^2}{\\alpha^2\\mu^2}. \\label{up2full} \\end{align}  Below we use another angular variable $\\Theta = \\pi/2 - \\theta$ and for the sake of simplification we adopt the polytropic equation of state $P = k(S)n^{\\Gamma}$ so that temperature and sound velocity can be written as~\\cite{sha83} \\begin{equation} T = k(S)n^{\\Gamma - 1}; \\qquad \\cs[2] = \\frac{\\Gamma}{\\mu}k(S)n^{\\Gamma - 1}. \\label{Tc} \\end{equation}  \\section[]{Subsonic flow} \\label{Sec.SubsonicFlow}  Following Sec.\\hbox{}~\\ref{Sec.intro}, we assume that the $\\alpha$-disc theory holds outside the \\MSO. We adopt the flow velocity components, which this theory yields on the \\MSO\\ $r=\\rms,$\\footnote{A nearly parallel inflow with a small radial velocity $v_r \\approx \\alpha_{\\rm SS} \\cs[2]/v_{\\rm K} \\ll \\cs \\ll 1$; $\\rms = 3 r_{\\rm g}$, where $r_{\\rm g} = 2M$ is the gravitational radius of the BH of mass $M$.} as the first three boundary conditions for our problem. For the sake of simplicity we consider the radial velocity, which is responsible for the inflow, to be constant at the surface $r = \\rms$ and equal to $u_0$ and the toroidal velocity to be exactly equal to that of a free particle revolving at $r=\\rms$.% \\footnote { For a free particle revolving at $r=\\rms$ around a non-spinning BH we have \\cite{lan87a} $u_{\\hat\\varphi}(\\rms) = 1/\\sqrt{3}, \\alpha_0 = \\alpha(\\rms) = \\sqrt{2/3}, \\gamma_0 = \\gamma(\\rms) = \\sqrt{4/3}.$ \\label{Ftn.MS} } We also assume the speed of sound to be constant at the \\MSO, $\\cs = c_0 = \\mathrm{const}$. Having introduced $\\Thms$ \\cite{bes02} --- the Lagrange coordinate of streamlines at the \\MSO\\ --- for $\\cs \\ll 1$, \\ie\\ non-relativistic temperature, we obtain from (\\ref{up2full}) and (\\ref{Tc}), \\begin{equation} \\up[2]  = u_0^2 +w^2 + \\frac{2}{\\Gamma - 1}\\left(c_0^2 - \\cs[2]\\right) + \\frac{1}{3}\\left(\\Thms[2]-\\Theta^2\\right) + \\dots \\label{uup} \\end{equation} The quantity \\begin{equation} w^2(r) = \\frac{e_0^2-\\alpha^2l_0^2/r^2-\\alpha^2}{\\alpha^2} \\equiv \\frac{1}{\\alpha^2}\\frac{(6-r)^3}{9r^3}, \\label{w2} \\end{equation} where $e_0 = E_0/\\mu_0$ and $l_0 = L/(\\mu_0 \\cos \\Thms)$, is the poloidal four-velocity of a free particle having zero poloidal velocity at the \\MSO.  In the extreme subsonic case, $\\up \\ll \\cs$, the Grad-Shafranov hydrodynamic equation is significantly simplified~\\cite{bes04,bes02}. The numerical results are shown in Fig.\\hbox{}~\\ref{Fig.global}. In the subsonic region, $r_* < r \\le \\rms \\equiv 3 r_{\\rm g}$, the disc thickness rapidly diminishes, and at the sonic surface we have $H(r_*) = u_0/c_0 H(\\rms)$, so that you cannot neglect the dynamical force there.  We stress that taking the dynamical force into account is indeed extremely important. This is because, unlike zero-order standard disc thickness prescription \\eqref{2}, the Grad-Shafranov equation has second order derivatives, \\ie\\ contains two additional degrees of freedom. This means that the critical condition only fixes one of these degrees of freedom (\\eg\\ imposes some limitations on the form of the flow) rather than determines the angular momentum of the accreting matter~\\cite{bes04,bes02}.  \\section[]{Transonic flow} \\label{Sec.TransFlow}  In order to verify our conclusions we consider the flow structure in the vicinity of the sonic surface in more detail. Since the smooth transonic flow is analytical at a singular point $r = r_*, \\Theta = 0$~\\cite{lan87b}, it is possible to express the quantities via a series of powers of $h = (r-r_*)/r_*$ and $\\Theta$. Substituting these expansions into equations of motion we get a set of equations on the coefficients which allows us to reconstruct the flow structure in the vicinity of the sonic point (see Fig.\\hbox{}~\\ref{Fig.transonic}), in particular \\begin{align} \\up[2] &= c_*^2 \\left[1-2\\eta_1 h+ \\frac{1}{6}(\\Gamma - 1) \\, \\frac{a_0^2}{c_0^2}\\Theta^2 + \\frac{2}{3}(\\Gamma + 1)\\eta_1^2\\Theta^2 \\right], \\nonumber \\\\ \\cs[2] &= c_*^2 \\biggl[1+\\left(\\Gamma-1\\right)\\eta_1 h + \\frac{1}{6}(\\Gamma -1 ) \\, \\frac{a_0^2}{c_0^2}\\Theta^2 - \\frac{1}{3}(\\Gamma - 1)(\\Gamma+1) \\eta_1^2 \\Theta^2 \\biggr], \\label{cs} \\end{align} where $a_0 = \\left[2/(\\Gamma+1)\\right]^ {(\\Gamma+1)/2(\\Gamma-1)}c_0/u_0$ gives the compression of streamlines, $a_0 = H(\\rms)/H(r_*)$, and $\\eta_1 \\sim u_0^{-1}$. Equation \\eqref{cs} yields shape of the sonic surface, $\\up = \\cs$; it has the standard parabolic form $h = (\\Gamma+1)\\eta_1\\Theta^2/3$.  Since the transonic flow in the form of a nozzle (see Fig.\\hbox{}~\\ref{Fig.transonic}) has longitudinal and transversal scales of one order of magnitude~\\cite{lan87b}, near the sonic surface we have $\\delta r_\\parallel \\approx \\delta r_\\bot,$ \\ie\\ $\\delta r_\\parallel \\approx H(r_*)$. Hence for thin discs (\\ie\\ for $c_0 \\ll 1$) this longitudinal scale is always much smaller than the distance from the BH, $\\delta r_\\parallel/r_* \\approx H(r_*)/r_* \\ll 1$. Only by taking the transversal velocity into account do we retain the small longitudinal scale $\\delta r_\\parallel \\ll r_{\\rm g}.$ This scale is left out during the standard one-dimensional approach.  \\section[]{Supersonic flow} \\label{Sec.SupersonicFlow} Since the pressure gradient becomes insignificant in the supersonic region, the matter moves here along the trajectories of free particles. Neglecting the $\\nabla_{\\theta}P$ term in the $\\theta$-component of relativistic Euler equation~\\cite{fro98}, we have~\\cite{abr97} \\begin{equation} \\alpha u_{\\hat r} \\frac{\\partial (r u_{\\hat \\Theta})} {\\partial r} + \\frac{(r u_{\\hat \\Theta})}{r^2} \\frac{\\partial (r u_{\\hat \\Theta})} {\\partial \\Theta} + (u_{\\hat\\varphi})^2 \\tan\\Theta = 0. \\label{theuler} \\end{equation} Here, using the conservation law of angular momentum, $u_{\\hat\\varphi}$ can be easily expressed in terms of radius: $u_{\\hat\\varphi} = 2 \\sqrt{3}/r$. We also introduce dimensionless functions $f(r)$ and $g(r)$: $\\Theta f(r) = r u_{\\hat \\Theta}$ and $g(r) = -\\alpha u_{\\hat r} > 0$. Using (\\ref{theuler}) and the definitions above, we obtain an ordinary differential equation for $f(r)$ which could be solved if we knew $g(r)$: \\begin{equation} \\frac{{\\rm d} f}{{\\rm d} r} = \\frac{f^2 + 12}{r^2 g(r)}. \\label{Eq.f} \\end{equation}  From (\\ref{up2full}) we have \\hbox{$\\up[2] \\rightarrow w^2$} as \\hbox{$r \\rightarrow r_{\\rm g}$}. On the other hand, $\\up \\approx c_* \\approx c_0$ for $r \\lesssim r_*$. Therefore, the following approximation should be valid throughout the \\hbox{$r_{\\rm g} < r < r_*$} region, $g(r) \\approx \\sqrt{(\\alpha w)^2 + (\\alpha c_*)^2}$.  Equation \\eqref{Eq.f} governs the supersonic flow structure for the case of non-spinning BH. To get a better match with observations (\\cf\\hbox{} Sec.\\hbox{} \\ref{Section.Observ}), we also consider a more general case of spinning BH, \\ie\\hbox{} a Kerr BH with non-zero specific angular momentum $a$. After some calculation, equation \\eqref{Eq.f} can be generalized to the Kerr metric with a strikingly simple form, \\begin{equation} \\frac{{\\rm d}f}{{\\rm d}r} = \\frac{f^2 + a^2 \\left(1-e_0^2\\right)+l_0^2}{r^2 \\tilde{g}\\left(r\\right)}, \\label{Eq.fKerr} \\end{equation} where $\\tilde{g}\\left(r\\right)$ is a straightforward generalization of $g(r)$ to the Kerr case; we omit it here due to space limitations. For the Schwarzschild BH ($a=0$, $e_0=\\sqrt{8/9}$, $l_0 = 2\\sqrt{3}$ \\cite{lan87b}) equation \\eqref{Eq.fKerr} reduces back to \\eqref{Eq.f}.  Integrating (\\ref{Eq.fKerr}), we obtain \\begin{equation} f\\left(r\\right) = \\kappa \\tan \\left[ \\kappa \\int_{r_*}^r \\frac{d\\xi}{\\xi^{2}g(\\xi)} + \\frac{\\pi}{2} \\right], \\end{equation} where $\\kappa = \\sqrt{a^2\\left(1-e_0^2\\right)+l_0^2}$ and $\\pi/2$ has been to a good accuracy substituted for the integration constant $\\arctan \\left[f(r_*)/\\sqrt{3}\\right]$.\\footnote{For $r$ just below $r_*,$ the function $f$ should be positive to reflect the fact that the flow diverges. Then, $f=0$ corresponds to the point where the divergency finishes, and the flow starts to converge.}  The results of numerical calculations are presented in Fig.\\hbox{}~\\ref{Fig.global}. \\begin{figure}% \\begin{minipage}[b]{0.65\\linewidth} \\begin{center} \\epsfig{figure=Fig1.eps,width=\\linewidth} \\caption{The structure of a thin accretion disc (actual scale) for $c_0 = 10^{-2}$, $u_0 = 10^{-5}$ after passing the \\MSO\\ $r = 3r_{\\rm g}$ ($a=0$, Schwarzschild case). As sufficient dissipation can take place in the vicinity of the first node $r = r_{\\rm 1}$, we do not prolong the flow lines to the region $r < r_{\\rm 1}$.} \\label{Fig.global} \\end{center} \\end{minipage}\\hfill \\begin{minipage}[b]{0.3\\linewidth} \\begin{center} \\epsfig{figure=Fig2.eps,width=\\linewidth} \\caption{Schematics of thin disc stre\\-am\\-lin\\-es pro\\-file around the sonic point. The flow has the form of the standard nozzle. Here \\hbox{$x = -h$}, and \\hbox{$y = \\Theta$}.} \\label{Fig.transonic} \\end{center} \\end{minipage} \\end{figure} In the supersonic region the flow performs transversal oscillations about the equatorial plane, their frequency independent of their amplitude. We see as well that the maximum thickness of the disc in the supersonic (and, hence, ballistic) region, which is controlled by the transverse component of the gravitational force, actually coincides with the disc thickness within the stable orbits region, $r > \\rms$, where standard estimate (\\ref{2}) is correct.  Once diverged, the flow converges once again at a `nodal' point closer to the BH. The radial positions the nodes are given by the implicit formula $f(r_n)=\\pm\\infty,$ \\ie\\ \\begin{equation} \\kappa \\int_{r_n}^{r_*} \\frac{d\\xi}{\\xi^{2}g(\\xi)} = n \\pi, \\end{equation} where $n$ is the node number; the node with $n=0$ corresponds to the sonic surface. In this formula the sonic radius $r_* \\equiv r_0$ can be to a good accuracy approximated by $\\rms = \\rms(a)$ the expression for which can be found in most textbooks~\\cite{sha83}. Figure~\\ref{Fig.nodespos} shows the positions of nodes for different values of $c_0$ (the positions do not depend on $u_0$ for $u_0 \\ll c_0$) and the BH spin parameter $a$. \\begin{figure} \\begin{minipage}[b]{0.64\\linewidth} \\begin{center} \\epsfig{figure=Fig3.eps,width=\\linewidth} \\caption{Radial positions (in the units of $M$) of the nodes for a range of initial sound velocities. Dotted, dashed, and solid curves correspond to the cases $a=-0.5$, $a=0$, and $a=0.5$ respectively. Each curve relates the radial position of a node to a value of the initial sound velocity. Intersection points of these curves with the line $c_0 = {\\rm const}$ give the the nodes' radial positions for that particular value of $c_0$.} \\label{Fig.nodespos} \\end{center} \\end{minipage}\\hfill \\begin{minipage}[b]{0.32\\linewidth} \\begin{center} \\epsfig{figure=Fig4.eps,width=\\linewidth} \\caption{Relation of $a$ vs.\\ $M$. Dashed, dash-dotted, and solid lines come from matching $\\overline T$ ($700 \\pm 100$ s), $T_\\ms$ ($2200 \\pm 300$ s), and $T_0$ ($1100 \\pm 100$ s) respectively. The resulting error polygon is bolded.} \\label{Fig.aM} \\end{center} \\end{minipage} \\end{figure}% The matter travel time between the nodes has weak dependence not only on $u_0$ but also on $c_0$ as well. This provides a means for testing the theory via observations, and we do this in the following section. ",
        "conclusions": ""
    },
    "0403/astro-ph0403705_arXiv.txt": {
        "abstract": "{ The limited angular resolution of $\\gamma$--ray telescopes prevents the straight identification of the majority of the sources detected so far.  This is particularly true for the low latitude, probably galactic ones, only 10 \\% of which has been identified.  Most of the counterparts are {\\it Isolated Neutron Stars}, both radio--loud and radio--quiet, characterised by an extremely high value of $f_{X}/f_{opt}$.  The best way to search for INSs in the error boxes of unidentified EGRET sources is to perform the X--ray coverage of the $\\gamma$--ray field, followed by the optical characterization of each X--ray source.  We applied this procedure to two EGRET sources, which could belong to a local galactic population:  3EG J0616--3310 and 3EG J1249--8330.  Here we report on the analysis of about 300 X--ray objects, as well as on their optical study. } \\authorrunning{N. La Palombara et al.} \\titlerunning{3EG J0616-3310 and 3EG J1249-8330} ",
        "introduction": "The third EGRET catalogue (\\citet{Hartman}) lists 271 high--energy $\\gamma$--ray sources: the high--latitude ones ($|b|\\ge 10^{\\circ}$), presumably extragalactic, are 183, while the remaining 88 are at low latitudes ($|b|\\le 10^{\\circ}$) and should belong to our Galaxy.  Source identification has been hampered by several problems:  the EGRET poor angular resolution ($\\simeq1^\\circ$); the source confusion; the variety of potential emitters; the low count statistics.  Up to now, only about 100 sources have been identified:  67 are blazars, 27 are candidate blazars, while only 7 sources have been associated to $\\gamma$--ray pulsars.  Therefore about 170 EGRET sources are still unidentified.  In particular, this is true for more than 90\\% of the low--latitude, presumably galactic ones (\\citet{Caraveo 2001}).  At low/medium latitude Isolated Neutron Stars (INSs) are the most promising candidates as $\\gamma$--ray source counterparts.  Up to now, INSs have provided the only confirmed identifications, while various attempts to associate the low--latitude unidentified EGRET sources with different classes of galactic objects have not yielded conclusive results (\\citet{Romero}; \\citet{Geherels}).  Unfortunately, the identification of an INS is very difficult:  the limited statistics prevents the direct detection of a $\\gamma$--ray time signature and may render difficult the search for an already known radio signal.  The situation is even worse for radio--quiet INSs, where the absence of a reference radio pulsation makes it even more difficult to detect a significant time signature.  In order to search for both {\\em radio--loud} and {\\em radio--quiet} INSs we can take advantage of one of their characteristics:  the extremely high value ($>$100) of the X--ray--to--optical ratio ($f_{X}/f_{opt}$).  Thus, INSs can be found through a 2--step procedure: \\begin{enumerate} \\item each $\\gamma$--ray error--box is covered by dedicated X--ray observations, which will yield an harvest of new X--ray sources \\item the error--box of each X--ray source is scrutinized in the optical waveband, in order to find its counterpart or absence thereof and to single out potential neutron stars using the $f_{X}/f_{opt}$ values \\end{enumerate}  Such an investigation strategy was devised for the identification of {\\em Geminga} (\\citet{Bignami}) and is being used for the study of other EGRET sources (\\citet{Mirabal}; \\citet{Caraveo 2001}).  In this work we applied it to two middle--latitude EGRET sources:  3EG J0616--3310 ($l=240^\\circ$ 20$'$ 45.4$''$, $b=-21^\\circ$ 14$'$ 31.7$''$) and 3EG J1249--8330 ($l=302^\\circ$ 51$'$ 35.2$''$, $b=-20^\\circ$ 37$'$ 44.8$''$).  They have been selected owing to their relatively good positional accuracy, spectral shape, galactic location and lack of candidate extra--galactic counterparts.  In Fig.~\\ref{galacticposition} we outline their celestial position compared to the other unidentified {\\em EGRET} sources.  \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics[angle=+90,clip=true]{3eg_fig1.ps}} \\caption{Sky  distribution  of  the  unidentified  {\\em EGRET}  sources:  the two red circles show the celestial  position of 3EG0616--3310 and 3EG1249--8330 (courtesy of S.Vercellone)}\\label{galacticposition} \\end{figure} ",
        "conclusions": "Thanks to the {\\em XMM-EPIC} observations, we have detected about 150 X--ray sources in both the error--boxes of 3EG J0616--3310 and 3EG J1249--8330.  We have found that $\\simeq$ 50\\% of the sources have no candidate optical counterpart, up to a limiting magnitude of $F \\simeq$ 20.  Since they are often high $f_{x}/f_{opt}$ objects, this sub--sample seems very promising in order to find the counterpart of the two unidentified EGRET sources.  Using the newly acquired WFI data, we will proceed towards the classification of the optical objects with an automatic alghorytm (\\citet{Hatziminaoglou}) and to complete the optical identification of our X-ray sources."
    },
    "0403/astro-ph0403533_arXiv.txt": {
        "abstract": "{ We study the binary corrected X-ray light curve of Cygnus X-3 using the RXTE-ASM (All Sky Monitor) data and RXTE-PCA data. We find that the X-ray template derived from the EXOSAT observations adequately explains the binary variations. Using the recently determined quadratic binary ephemeris we correct the RXTE-PCA lightcurve and obtain the power spectrum which shows distinct features of shifting towards the low frequency regime vis-a-vis the Galactic X-ray binaries. The power  spectrum doesn't vary from the hard to the soft state. We also  obtain the binary corrected RXTE-ASM monitoring lightcurve and examine the previously obtained correlation between soft X-ray (2 -- 12 keV from RXTE ASM), hard X-ray (20 -- 100 keV from CGRO -- BATSE) and radio (2.2 G Hz from GBI) after the binary correction. We find that the correlation time scale is less than a day. ",
        "introduction": " ",
        "conclusions": "In this paper we have suggested a method to correct the  X-ray light curve of Cygnus X-3 for the binary modulation. The binary period shows consistency over a long period of 25 years, with minor variation in the asymmetric binary phase profile. An X-ray halo consisting of photoionized plasma originating from the wind of the companion Wolf-Rayet type companion star reprocesses the X-ray emission, while any eccentricity of the binary orbit which may cause asymmetry in the X-ray binary profile is minimal because of absence of any detectable apsidal motion over the long period of $\\sim$ 25 years.  We have also examined the binary corrected light curve for short term variabilities and noted that the shape of the PDS follows the power law behavior in both low and high states, similar to other Galactic black hole sources. Most significantly there exists no power at frequency $>$ 0.1 Hz, with the PDS being shifted to the low frequency regime vis-a-vis the Galactic X-ray binaries, both black hole and neutron stars. This is a unique feature of this enigmatic source.  Using the corrected light curve we have examined the X-ray radio correlations at various time scales and have concluded that the time scale is shorter than a day. The soft and hard X-ray correlation need to be examined using the binary corrected RXTE pointed (PCA and HEXTE) data covering a time scale of a few hours or longer to unambiguously determine the time scale of anti-correlation, and hence determine the detailed physical as well geometrical structure of the disk."
    },
    "0403/astro-ph0403643_arXiv.txt": {
        "abstract": "We are using the NRAO VLBA at 5.0~GHz to image about 200 FIRST sources stronger than 10~mJy at 1.4~GHz in the NDWFS. ",
        "introduction": "Traditional continuum surveys using Very Long Baseline Interferometry (VLBI) have had limited astrophysical impact, because they suffer from a paucity of redshifts and from biases introduced by targeting bright and, often, flat-spectrum radio sources.  To overcome these limitations, we are using the NRAO Very Long Baseline Array (VLBA) at 5.0~GHz to image about 200 compact FIRST sources stronger than 10~mJy at 1.4~GHz (White et al.\\ 1997) in the NOAO Deep Wide-Field Survey (NDWFS).  Optical identifications to 26 magnitudes (mag) are becoming available from the NDWFS (Jannuzi \\& Dey 1999) and spectroscopic follow-up has commenced for many of the FIRST sources. ",
        "conclusions": ""
    },
    "0403/astro-ph0403196_arXiv.txt": {
        "abstract": "The age of the universe is obtained in a subset of Cardassian models by using WMAP data.  Cardassian expansion is a modification to the Friedmann equation that allows the universe to be flat, matter dominated, and accelerating, without a vacuum component. Since this model changes the evolution of the universe, we should not \\textit{a priori} expect the Cardassian age to be the same as the WMAP Friedmann derived result of $13.7 \\pm 0.2$ Gyrs.  However, in the subset of Cardassian models we consider, we discover that the age of the universe varies from 13.4-13.8 Gyr over the range of parameter space we explore, a result close to that of the standard $\\Lambda$ Cold Dark Matter model. The Hubble constant $h$, which may also vary in these models, likewise varies little from the Friedmann result. ",
        "introduction": "Introduction}  Over the past five years, Cosmic Microwave Background (CMB) data have shown the universe to be flat \\cite{omega1, wmap} and supernova data have indicated that it is accelerating \\cite{SNCP:HighZ,SNST:Accel}.  The matter content of the universe falls well short of the necessary energy density to provide a flat curvature in the standard cosmological model, the Friedmann model.  The most popular interpretation of this mismatch is that the ``missing'' density is assumed to be present in the form of a vacuum (or dark) energy that provides a pressure leading to the acceleration of the universe.  Alternatively, general relativity may need to be modified. Modification of the standard cosmological model, as in the Cardassian model \\cite{Freese:2002sq}, allows both the flatness and acceleration to be accounted for solely due to the matter content of the universe. This modification changes the evolution of the universe.  Hence the CMB anisotropy spectrum will imply a different age of the universe \\cite{wmap}.  The Cardassian model \\cite{Freese:2002sq,Freese:2002gv,Gondolo:2002fh, Wang:2003cs,Sen:2003cy} % modifies the Friedman equation by adjusting the right hand side to be a more general function of the energy density.  This function returns to the usual Friedmann equation for the early history of the universe, so that ordinary nucleosynthesis takes place. However, in the recent past, beginning at a redshift $O(1)$, this function drives the universe to accelerate.  Such modifications to the Friedmann equation may arise if our universe exists on a brane embedded in a higher dimensional universe \\cite{Chung:XDimensions}. Other proposed modifications to the Friedmann equations include \\cite{other}.  In section~\\ref{sec:Friedmann}, we will review the Friedmann model and describe how to determine the age of the universe.  In section~\\ref{sec:WMAP}, Wilkinson Microwave Anisotropy Probe (WMAP) results are reviewed and the choice of appropriate parameters is discussed.  In section~\\ref{sec:Cardassian}, we will expand these techniques to the Cardassian model. We note that, since the customary parameters $h$, $\\Omega_b$, and $\\Omega_m$ are not orthogonal and indeed have values that depend upon the cosmological model itself, more suitable parameters must be extracted from the data in order to obtain the universe age in any model. And finally, results for the Cardassian model are discussed in section \\ref{sec:Results}. In this paper, we obtain the age of the universe only for the Modified Polytropic (MP) Cardassian models defined below. ",
        "conclusions": ""
    },
    "0403/hep-ph0403230_arXiv.txt": {
        "abstract": "In a recent paper by Link, it was pointed out that the the standard picture of the neutron star core, composed of a mixture of a neutron superfluid and a proton type-II superconductor, is inconsistent with observations of long period precession in isolated pulsars. In the following we will show that intervortex force between the magnetic flux tubes may be attractive resulting in a type-I (rather than type-II) superconductor. In this case the magnetic field cannot exist in the form of magnetic flux tubes, supporting Link's observation. This behavior of the system is due to the strong interaction between the proton-neutron Cooper pairs, which was previously ignored. We also calculate the critical magnetic fields $H_c$ and $H_{c2}$ for type-I/II superconductors. These results also support our claim of type-I superconductivity in the cores of neutron stars. ",
        "introduction": "Recently, it was pointed out by Link \\cite{link} that observations of long period precession \\cite{precession} may be inconsistent with the standard picture of the interior of a neutron star. In the conventional picture, the extremely dense interior is mainly composed of neutrons, with a small amount of protons and electrons in beta equilibrium. The neutrons form $^3P_2$ Cooper pairs and Bose condense to a superfluid state, while the protons form $^1S_0 $ Cooper pairs and Bose condense to give a superconductor as well (see e.g. Ref. \\cite{review} for a recent review). It is generally assumed that the proton superconductor is a type-II superconductor, which means that it supports a stable lattice of magnetic flux tubes in  the presence of a magnetic field. This belief is based on  simple estimations of the coherence length and the London penetration depth which  ambiguously imply a type-II superconductivity. In addition, the rotation of a neutron star causes a lattice of quantized vortices to form in the superfluid neutron state, similar to the observed vortices  that form when  superfluid $He$ is rotated fast enough. The axis of rotation and the axis the magnetic field are not aligned. This, coupled with the fact that the two different types of vortices interact quite strongly lead to the suggestion that the observed long period precession may require reexamining the picture of type-II superconductivity inside neutron stars \\cite{link} that follows from the standard analysis when only a single proton field is considered. A possible resolution of this ``apparent contradiction\" was suggested in our recent letter \\cite{typeI}, where we argued that if one takes into account that the Cooper pairs of neutrons are also present in the system and they strongly interact with proton Cooper pairs, the superconductor may in fact be type-I, even when a naive analysis that only includes the proton degrees of freedom  seems to indicate type-II behavior. If this scenario is realized in nature, this means that the interior of neutron stars would exhibit the Meissner effect and therefore would not support a stable lattice of magnetic flux tubes. This would resolve the apparent discrepancy \\cite{link} between the observation of long period precession \\cite{precession} and the typical parameters of the neutron stars which naively suggest type-II superconductivity in neutron stars. The main goal of this paper is to provide a  detailed analysis   with complete calculations supporting the claim  of  our short letter \\cite{typeI}.  It is well known that type-II superconductors have magnetic flux tubes in the presence of a magnetic field. In the interior of neutron stars, which is a mixture of neutron and proton superfluids, the proton superfluid will form a vortex lattice of magnetic flux tubes if the superconductor is type-II. Inside the core of these vortices, the proton condensate vanishes, and the core is filled with normal protons resulting in the restoration of the broken \\UEM symmetry. For accepted estimates of proton correlation length and London penetration depth, the distant proton vortices repel each other leading to formation of a stable vortex lattice. This is the standard picture realized in conventional type-II superconductors. However, there are many situations where this picture will be qualitatively modified. For example, if there is a second component (such as a neutron component in our specific case), it may be energetically favorable for the cores of vortices to be filled with a nonzero condensate of this second component, as it was originally suggested in the cosmological context by Witten \\cite{witten,defectsbook}. There are numerous examples of physical systems where this phenomena occurs: superconducting cosmic strings in cosmology, magnetic flux tubes in the high $T_c$ superconductors, Bose-Einstein condensates, superfluid $~^3He$, and high baryon density quark matter \\cite{witten,defectsbook,volovik,BEC,highTc,krstrings,superk}. Given this, one might guess that such nontrivial vortex structure may occur in the core of neutron stars where we have another example of a two component system. We shall argue in what follows that if the interaction between proton and neutron Cooper pairs is approximately equal (a precise  condition of this ``approximately\" will be derived below), the vortex-vortex interaction will be modified and the system will be a type-I superconductor where the magnetic field is completely expelled from the bulk.\\footnote{In reality, the magnetic field must be present in the neutron star interior. The way this picture may be realized in nature is through the formation of domains of superconductor matter and normal matter.} The main assumption  that we are making is that the interactions between the proton and neutron Cooper pairs are approximately equal, leading to an approximate $U(2)$ symmetry. We believe that this assumption is justified by the original isospin $SU(2)$ symmetry of the   neutrons and protons. The result of this is that the proton vortices or magnetic flux tubes have nontrivial core structure. The superfluid density of the neutrons is larger in the vortex core than at spatial infinity. In addition, the size of the vortex core and the asymptotic behavior of the proton condensate is modified due to the additional neutron condensate. The most important result of these effects is that the interaction between distant proton vortices may be attractive in a physically realizable region of parameter space leading to type-I behavior: destruction of the proton vortex lattice and expulsion of the magnetic flux from the superconducting region of the neutron star.  This paper is organized as follows. In Sec. II we will introduce the free energy that describes a two component (neutron-proton) superfluid/superconducting system and show that the addition of the second component (neutron) leads to nontrivial core structure, which alters the properties of the magnetic flux tube. In Sec. III, we will give two different calculations of the interaction between two widely separated vortices. These two different methods lead to the same conclusion, that the interior of a neutron star may be a type-I superconductor. In the subsequent section, Sec. IV, we will calculate the critical magnetic fields $H_c$ and $H_{c2}$ (the magnetic fields above which superconductivity is destroyed in type-I and type-II superconductors, respectively). These results confirm our findings in Sec. III that type-I superconductivity may occur in the interior of neutron stars. In Sec. V we will end with concluding remarks on possible implications of our results. Specifically, we will comment on the possible nature of glitches (observed in many systems) when  the environment of the neutron star core is a type-I superconductor. ",
        "conclusions": "In this paper we have demonstrated using various calculations that the proton superfluid present inside a neutron star may in fact be a type-I superconductor \\cite{typeI}. This supports the observation made my Link \\cite{link} that the conventional picture of type-II superconductivity may be inconsistent with the observations of long period precession in isolated pulsars \\cite{precession}. The most important consequence of this paper is that whether the proton superconductor is type-I or type-II depends strongly on the magnitude of the $SU(2)$ asymmetry parameter $\\epsilon$. Specifically, we find that the superconductor is type-I when $\\kappa_{np} = \\sqrt{\\epsilon} \\lambda/\\xi < 1/\\sqrt{2}$, and type-II otherwise. This result is quite generic, and not very sensitive to the specific details  of the interaction potential $V$. In particular, when $\\epsilon \\rightarrow 0$ the superconductor is type-I. The parameter $\\epsilon$ is not known precisely; the corresponding  microscopical calculation would require the analysis of the scattering lengths of Cooper pairs for different species. We can roughly estimate this parameter  as being related to the original $SU(2)$ isospin symmetry breaking $\\epsilon \\sim (m_n - m_p)/m_n \\sim 10^{-2}$. If we assume a typical value for $\\lambda/\\xi \\sim 3$ and $\\epsilon \\sim 10^{-2}$, we estimate $\\kappa_{np} =\\sqrt{\\epsilon} \\lambda/\\xi \\sim 0.3 < 1/\\sqrt{2}$, which corresponds to a type-I superconductor. From these crude estimates, we see that it is very likely that neutron stars are type-I superconductors with the superconducting region devoid of any magnetic flux, as was originally suggested in \\cite{link} to resolve the inconsistency with observations of long period precession \\cite{precession} in isolated pulsars. Many mechanisms have been proposed to explain glitches \\cite{glitch1}. If the proton superfluid exhibits type-I superconductivity then some explanations of glitches \\cite{glitchflux} that assume type-II superconductivity would have to be reconsidered, as suggested in Ref. \\cite{link}. It might be interesting to consider how the presence of a nonzero proton condensate affects the characteristics of the  neutron vortices that carry angular momentum. In particular, as we have demonstrated for magnetic flux tubes in this paper, the neutron vortices might have an enhanced proton superfluid density inside, as well as a coherence length $\\xi_n$ (the approximate size of the vortex) which is much larger than originally expected. In this case, the pinning force that is related to the size of the vortex core \\cite{pinning} could be very different from the simple estimates when the strong interaction between the neutron and proton Cooper pairs is ignored.  If the core is indeed a type-I superconductor, the magnetic field could exist in macroscopically large regions where there are alternating domains of superconducting (type-I) matter and normal matter. In this case, neutron stars could have long period precession. Such a structure follows from few different arguments. First of all, as has been estimated \\cite{Baym}, it takes a very long time   to expel a typical magnetic flux from the neutron star core. Therefore, if the magnetic field existed before the neutron star became a type-I superconductor, it is likely that magnetic field will remain there. Another argument suggesting the same outcome follows from the fact that topology (magnetic helicity) is frozen in the environment with high conductivity; therefore, the magnetic field must remain in the bulk of the neutron star. The last argument supporting the same picture is due to Landau \\cite{LandauLam} who argued that if a body of arbitrary shape (being a type-I superconductor) is placed under influence of the external magnetic field with strength $H < H_c$, then   the magnetic field in  some parts of the body may reach the critical value $H_c$, while in other parts of the body it may remain smaller than the critical value, $H < H_c$. In this case, a domain structure will  be formed, similar to ferromagnetic systems.  Specifically, on a macroscopic distance scale, the magnetic flux must be embedded in the superconductor. This would mean that the superconductor is in an intermediate state as opposed to the vortex state of the type-II superconductor, which was assumed to be realized up to this point. The intermediate state is characterized by alternating domains of superconducting and normal matter. The superconducting domains will then exhibit the Meissner effect, while the normal domains will carry the required magnetic flux. The pattern of these domains is usually strongly related to the geometry of the problem. The simplest geometry, originally considered by Landau \\cite{LandauLam}, is a laminar structure of alternating superconducting and normal layers. However, it has been also argued \\cite{Sedrakian} that due to the geometry of the neutron vortex lattice, the normal proton domains will be in the shape of cylinders concentric with the rotational neutron vortices.  While some precise calculations are required for understanding of the magnetic structure in this case, one can give some simple  estimation of  the size of the domains using the calculations Landau presented for a different geometry. His formula \\cite{LandauLam} suggests that the typical size of a domain is \\beq \\label{domain} a\\sim 10 \\sqrt{R \\Delta}, \\eeq where $R$ is a typical external size identified with a neutron star core ($R\\sim 10$~km), while $\\Delta$ is a typical width of the domain wall separating normal and superconducting states. We estimate $\\Delta\\sim \\delta =\\xi/\\sqrt{\\epsilon}$ as the largest microscopical scale of the problem. Numerically, $a\\sim 10^{-1}$~cm which implies that a typical domain with size of order $\\sim 10^{-1}$~cm can accommodate $\\sim 10 $ neutron vortices separated by a distance $\\sim 10^{-2}$~cm.  The consequences of this picture still remain to be explored. In particular, it might be of importance in the explanation of glitches. It may be also important in analysis of the cooling properties of the neutron stars.  It would be very interesting to test the ideas outlined in this paper by doing laboratory experiment in the spirit of the Cosmology in the Laboratory (COSLAB) program. In particular, it would be interesting to find a condensed matter system (high $T_c$ superconductor?) where the very interesting feature discussed in this paper can be tested. Namely, the core of the vortices and their interactions could be very different from the simple estimates. This might happen if there is a condensate of another field with energy scales almost degenerate with energies determined by the  Landau-Ginsburg complex $\\psi$ field describing a superconductor. Over the last few years several experiments have been done to test ideas drawn from cosmology and astrophysics (see Ref. \\cite{volovik} and the web page \\cite{COSLAB} of the latest COSLAB meeting for further details)."
    },
    "0403/astro-ph0403313_arXiv.txt": {
        "abstract": "Most astronomers now accept that stars more massive than about $9\\,M_{\\odot}$ explode as supernovae and leave stellar remnants, either neutron stars or black holes, with neutron stars being more prevalent.  Recent modeling of the explosions suggests a significant diversity in the key natal properties --- rotation rate, velocity, and magnetic field strength --- of the resulting neutron stars that account for the association of active radio pulsars, pulsar wind nebulae, and magnetars with supernova remnants (SNRs). The discovery of a central X-ray source in Cas~A, the youngest known Galactic SNR, dramatized the expected diversity. However, less than half of the SNRs within 5 kpc have identified central sources, and only three are identified as the remnants of Type~Ia SNe.  Here, we report a systematic effort to search for compact central sources in the remaining 23 SNRs of this distance limited sample.  Our search was inspired, on empirical considerations, by the enigmatic faint X-ray source in Cas~A; motivated, on theoretical grounds, by the expectation that young neutron stars emit cooling X-ray emission; and made possible by the superb angular resolution offered by the {\\em Chandra} X-ray mission and the sensitivity of the \\textit{XMM-Newton} mission.  In this first paper we report {\\em Chandra} observations of four SNRs (\\Ga, \\Gb, \\Gc, and \\Gd).  We have undertaken a systematic optical/IR identification program of the X-ray sources detected in the field of each SNR. Foreground (flare stars, active stars) and background (active galactic nuclei) sources have identifiable IR/optical counterparts. In contrast, the counterparts of neutron stars (or black holes) are expected to be very faint.  We are able to account for all the well-detected X-ray sources and thus able to state with some confidence that there are no associated central sources down to a level of one tenth of that of the Cas~A central source, $L_{X}\\lsim 10^{31}\\mbox{ ergs s}^{-1}$.  We compare our limits with cooling curves for neutron stars and find that any putative neutron stars in these SNRs must be cooling faster than  expected for traditional $1.35\\,M_{\\odot}$ neutron stars and that any putative pulsar must have low spin-down luminosities $\\dot E \\lsim 10^{34}\\mbox{ ergs s}^{-1}$. However, our limits are unable to constrain the presence or absence of more unusual options, such as relatively more massive neutron stars with $M\\gsim 1.45\\,M_\\odot$, neutron stars with exotic interiors,  or quiescent black holes. In subsequent papers, we will report on the X-ray and optical/IR observations of the remaining members of the 5-kpc sample. ",
        "introduction": "Understanding the deaths of massive stars is one of the frontiers of modern astrophysics.  Considerable observational evidence substantiates the idea that stars below $M_{w}\\sim 8\\,M_\\odot$ end their lives as white dwarfs \\citep{weidmann87}, while the detection of neutrinos from SN~1987A dramatically illustrated that more massive stars undergo core collapse \\citep{hkk+87}.  The outcome of core collapse can either be a neutron star or a black hole \\citep*{whw02}.  However, there are great uncertainties in the mapping between initial mass of the star and the end product, and even more uncertainties in the natal properties of the stellar remnant.  It is these uncertainties that give observers opportunities to make new discoveries and theorists to predict or ``postdict'' these discoveries.  The first issue --- the state of the star prior to collapse --- is very dependent on the mass loss history of stars \\citep{hfw+03}, a phenomenon that is poorly understood and can easily be modified by the presence of a binary companion.  For solar metallicity (the situation relevant to this paper), stars with between $M_l \\sim 9\\,M_\\odot$ and $M_u\\sim 25\\,M_\\odot$ are expected to form a neutron star \\citep{hfw+03}, while stars above $M_u$ are expected to form a black hole either by fall-back of material which transmutes the neutron star to a black hole or by direct collapse.  As an aside we note that progenitors with masses between $M_{w}$ and $M_l$ may form an O-Ne-Mg white dwarf that may collapse to neutron stars or simply explode \\citep{mnys80,nomoto84,nomoto87}.    The second issue --- the natal properties of the stellar remnant --- apparently involves delicate physics but has strong observational ramifications.  The gravitational binding energy of a neutron star is $10^{53}\\,$erg, of which only 1\\% appears to be coupled to the ejecta (which ultimately powers the SNR).  Even more minuscule fractions go into rotational energy, kinetic energy (bulk motion) and magnetic fields.  It is now generally agreed that three dimensional effects in the explosion determine the natal properties \\citep{burrows00,kpj+03}.  So far the discussion has assumed that the only gross parameter of interest is the mass of the star. However, it is likely that rotation of the progenitor can also profoundly affect the outcome.  These two issues are now propelling two different areas of inquiry. The relationship between the progenitor properties (mass, rotation) and the gross outcome of core collapse (neutron star or black hole) is observationally being determined by systematic studies of supernovae and GRBs and their interconnection. The second area is in understanding the natal properties of neutron stars, which is the main focus of the paper.  The discovery of pulsars in the Vela SNR \\citep*{lvm68} and the Crab Nebula \\citep{sr68} made concrete the suggestion that core collapse results in neutron stars \\citep{bz34}, some of which manifest themselves as radio pulsars.  Young pulsars, in addition to pulsing in the radio, can and usually do power synchrotron nebulae \\citep{wp78} that are indirect markers of pulsars.  These synchrotron nebulae are commonly called pulsar wind nebulae, or PWNe.  Over the following two decades, the notion that neutron stars resemble the Crab pulsar guided the search for central objects as well as intensive radio mapping of SNRs. As a result of these efforts, the term ``composite'' SNR (PWN + shell) was added to the SNR lexicon (see \\citealt{mgh+79}).  However, over the last five years there have been three developments that have severely revised our picture of young neutron stars.  First, astronomers have come to accept of tremendous diversity in the natal properties of young neutron stars.  Anomalous X-ray pulsars (AXPs), soft $\\gamma$-ray repeaters (SGRs), nearby thermal and radio quiet neutron stars, long period radio pulsars with high inferred magnetic fields (HBPSR) are now routinely found in astronomical literature. These new classes of neutron stars have primarily come from high energy (X-ray and $\\gamma$-ray) observations.  Second, there is appreciation that the radio luminosities of PWNe is poorly dependent on the spin-down luminosity of their central pulsars. For example, energetic pulsars may have  faint PWNe (e.g., PSR~J1119$-$6127; \\citealt{gs03}), and very sensitive imaging of the regions around identified central sources has frequently yielded only upper limits to the surface brightness of putative PWNe (e.g., \\citealt*{gbs00}). Third, seven nearby cooling NSs have been identified \\citep{haberl03} through \\rosat\\ observations. Many of these neutron stars do not appear to be evolved versions of standard radio pulsars; e.g., RX~J0720.4$-$3125 has a period longer than almost any known radio pulsar but has a typical $B$-field strength \\citep{kkvkm02,zhc+02}.  While this diversity is clearly demonstrated observationally, theory and simulation cannot yet constrain the fundamental birth properties of neutron stars \\citep[e.g.,][]{bom+03}.  Models still have difficulties achieving explosions, much less following the activity in the post-collapse object in any detail.  Three years ago we began a program of observationally investigating the stellar remnants in a volume-limited census of Galactic supernova remnants. The approach we took was inspired by the first light picture of \\chandra, the discovery of a central X-ray source in the well-studied and youngest known supernova remnant in our Galaxy, Cassiopeia~A \\citep{t99}.  The nature of the object continues to be debated \\citep{pza+00,cph+01,mrj+02}.  However, one conclusion is crystal clear: the X-ray source is not a standard radio pulsar (unbeamed or otherwise).  The basis of our effort is that observationally, all central sources in SNRs known to date, regardless of the band of their initial identification ($\\gamma$-ray, X-ray, or radio) appear to possess detectable X-ray emission. Theoretically, we expect thermal X-ray emission from young neutron stars. Thus, on both counts the search for central sources in young remnants is very well motivated. However, a follow-up program is essential since many other foreground sources such as flare stars, young stars, and accreting sources and background sources such as AGN dominate the source counts \\citep{hg88,mcba00}. Fortunately, the sub-arcsecond spatial resolution of \\chandra\\ allow efficient filtering of such contaminating objects.  To this end, we have identified a sample of SNRs within 5\\,kpc of the Sun which do not have known radio pulsars or PWNe and have not been associated with Type~Ia explosions (Table~\\ref{tab:snrs}.  We successfully proposed for a ``large'' \\chandra\\ effort in AO-3 to image nine SNRs. This initial allocation has been supplemented with additional time in AO-4 of \\chandra\\ and AO-2 of \\xmm.  Followup of the X-ray sources has been undertaken with a plethora of ground based telescopes in the optical and near-IR bands (Palomar 60-inch, Palomar 200-inch, Las Campanas 40-inch, ESO 3.5-m, Magellan 6.5-m and Keck 10-m). Here, we report the first analysis of four SNRs for which the followup is now complete. The analysis for the remaining remnants will be reported in future papers.  The organization of the paper is as follows. In \\S~\\ref{sec:class} we summarize the empirical X-ray properties of the known sample of young neutron stars. Such a summary is essential since the guiding principle of our effort is to place our search for central objects against the framework of existing classes of sources. Specifically, we are not entirely guided by the relatively poorly understood cooling of neutron stars. Our search has been designed to find objects as faint as one tenth of the central X-ray source in Cassiopeia~A.  In \\S~\\ref{sec:sample} we present a complete sample of cataloged SNRs within 5\\,kpc. Of these, 18 have an identified central source or are known to be a composite remnant, while three are thought the be the results of Type~Ia SNe. The remaining 23 seemingly ``hollow'' SNRs form our primary sample.  By ``hollow'', we refer to SNRs that have distinct shells but no obvious indication of central neutron stars --- see \\citet{vg97}.  Section~\\ref{sec:obs} has a general overview of our observations and analysis techniques: in \\S~\\ref{sec:Chandra-survey}--\\ref{sec:timing} we present a summary of the details of our \\chandra\\ observations and data reduction, and a likewise global description of the extensive multi-wavelength followup is given in \\S~\\ref{sec:cpt}.  We follow this by detailed descriptions of each of the first four SNRs, its observations, and identification of counterparts to its X-ray sources in \\S\\S~\\ref{sec:snra}, \\ref{sec:snrb}, \\ref{sec:snrc}, and \\ref{sec:snrd} for SNRs \\Ga, \\Gb, \\Gc, and \\Gd, respectively.  It must be appreciated that complete identification of all X-ray sources is essential, given the small sample size.  We hope that our detailed cataloging will be of help to efforts such as {\\em ChaMPlane} \\citep{gzh+03}.  Finally, in \\S~\\ref{sec:limits} we discuss what limits our data can place on the existence of central sources in the four SNRs, and we conclude in \\S~\\ref{sec:conclusions}.  \\begin{deluxetable}{l c c l r c c c c l} \\setlength{\\tabcolsep}{0.03in} \\tablecaption{SNRs Within 5 kpc\\label{tab:snrs}} \\tablewidth{0pt} \\tabletypesize{\\scriptsize} \\tablehead{ \\colhead{SNR G} & \\colhead{Other Name} & \\colhead{Dist.} & \\colhead{Distance Method} & \\colhead{Size\\tablenotemark{a}} & \\colhead{Type\\tablenotemark{b}} & \\colhead{Central Source/} & \\colhead{X-ray?\\tablenotemark{d}} & \\colhead{Sample\\tablenotemark{e}} & \\colhead{Refs\\tablenotemark{f}} \\\\ &  & \\colhead{(kpc)} & & \\colhead{(arcmin)} & & \\colhead{Ia} \\\\ } \\startdata 004.5+6.8   &Kepler\\tablenotemark{g} & 4.5 & Optical expansion/\\scrhi & 3 & S & Ia?\\tablenotemark{c} &\\nodata &\\nodata & 1\\\\ 005.4$-$1.2 & Duck & 5 & \\scrhi\\ absorption & 35 & C? & pulsar?  & yes & \\nodata & 2,3 \\\\ 006.4$-$0.1 &W 28\\tablenotemark{g} & 3.0 & OH masers & 42 & S & \\nodata &\\nodata &\\nodata & 4\\\\ 011.2$-$0.3 & & 5 & \\scrhi\\ absorption & 4 & C & pulsar  & yes &\\nodata& 5 \\\\ 013.3$-$1.3 & & 2.4 & CO absorption & 70 & S & \\nodata  & \\nodata & \\xmm-AO2 \\\\[0.1in] 034.7$-$0.4 & W44 & 2.5 & \\scrhi\\ absorption & 35 & C & pulsar  & yes& \\nodata& 6,7,8\\\\ 039.7$-$2.0 & W50 & 5& VLBI & 120 & ? & SS 433  & yes &\\nodata& 9,10 \\\\ 053.6$-$2.2 &3C 400.2 & 2.8 & \\scrhi\\ association & 33& S & \\nodata &\\nodata & \\cxo-AO3 \\\\ 054.4$-$0.3 &HC 40 & 3.3 & CO association & 40 & S & \\nodata  &\\nodata & \\cxo-AO3 \\\\ 065.3+5.7   & & 0.8 & Optical velocity & 310 & S & \\nodata  &\\nodata & \\cxo-AO4 \\\\[0.1in] 069.0+2.7 & CTB 80 & 2 & \\scrhi\\ absorption & 80 & ? & pulsar  & yes &\\nodata& 11,12\\\\ 074.0$-$8.5 & Cygnus Loop & 0.44 & Optical proper motion & 230 & S & X-ray source?  & \\nodata & \\cxo-AO4 & 13\\\\ 078.2+2.1 & $\\gamma$ Cygni & 1.2 & OH/CO association & 60 & S & \\nodata  &\\nodata & \\xmm-AO2 & 14 \\\\ 084.2$-$0.8 & & 4.5 & \\scrhi/CO association & 20 & S & \\nodata  &\\nodata & \\cxo-AO3 \\\\ 089.0+4.7 & HB 21 & 0.8 & OB association & 120 & S & \\nodata  &\\nodata & \\nodata\\\\[0.1in] 093.3+6.9   &DA 530 & 3.5 & \\scrhi\\ absorption/X-ray fitting & 27 & S & \\nodata &\\nodata  & \\cxo-AO3 \\\\ 111.7$-$2.1 &Cas A & 3.4 & Optical expansion & 5 & S & CCO & yes & \\nodata & 15 \\\\ 114.3+0.3 & & 3.5 & \\scrhi\\ association & 90 & S & pulsar & yes &\\nodata & 16,17,18 \\\\ 116.5+1.1 & & 4 & \\scrhi\\ association & 80 & S & \\nodata &\\nodata  & \\nodata \\\\ 116.9+0.2   &CTB 1 & 3.1 & Optical lines & 34 & S & \\nodata &\\nodata& \\cxo-AO3& \\\\[0.1in] 119.5+10.2 & CTA 1 & 1.4 &\\scrhi\\ association  & 90 & S & $\\gamma$-/X-ray neutron star &\\nodata  & \\nodata & 19,20,21  \\\\ 120.1+1.4 & Tycho & 2.4 & Proper motion/shock vel. & 8 & S & Ia & \\nodata & \\nodata & 22 \\\\ 127.1+0.5   &R5 & 1.3 & Assoc.\\ with NGC 559 & 45 & S & \\nodata &\\nodata  & \\cxo-AO3 \\\\ 130.7+3.1 & 3C 58 & 3.2 & \\scrhi\\ absorption & 9 & F & pulsar &yes &\\nodata& 23,24\\\\ 132.7+1.3 & HB 3 & 2.2 & Interactions with ISM & 80 & S & \\nodata &\\nodata  & \\xmm-AO2 \\\\[0.1in] 156.2+5.7 & & 1.3 & NEI fits & 110 & S & \\nodata &\\nodata  & \\cxo-AO4& \\\\ 160.8+2.6 & HB9 & $<4$ & \\scrhi\\, optical vel. & 140 & S & \\nodata &\\nodata  & \\cxo-AO4 & \\\\ 166.0+4.3   &VRO 42.05.01 & 4.5 & \\scrhi\\ association & 55 & S & \\nodata &\\nodata  & \\cxo-AO3 \\\\ 166.2+2.5 & & 4 & \\scrhi\\ interaction & 90 & S & \\nodata &\\nodata  & \\nodata & \\\\ 180.0$-$1.7 &S147 &1.2 & pulsar dispersion measure & 180 & S & pulsar & yes &\\nodata& 25,26,27\\\\[0.1in] 184.6$-$5.8 & Crab& 2 & proper motion/radial vel.\\ & 7 & F & pulsar & yes &\\nodata& 28 \\\\ 189.1+3.0   &IC 443 & 1.5 & Opt.\\ vel./assoc.\\ with S249 & 45 & S & neutron star & yes & \\nodata & 29 \\\\ 203.0+12.0 & Monogem Ring & 0.3 & \\scrovi\\ absorption/modeling & 1500 & S & pulsar & yes & \\nodata & 30\\\\ 205.5+0.5 & Monoceros & 1.2 & Optical velocity & 220 & S & \\nodata &\\nodata  & \\cxo-AO4\\\\ 260.4$-$3.4 & Puppis A & 2.2 & \\scrhi\\ association  & 60 & S & CCO & yes & \\nodata & 31,32 \\\\ 263.9$-$3.3 & Vela & 0.3 & pulsar parallax & 255 & C & pulsar & yes &\\nodata& 33 \\\\[0.1in] 296.5+10.0 & PKS 1209$-$51/52 & 2.1 & \\scrhi\\ association  & 90 & S & CCO/X-ray pulsar & yes & \\nodata & 34,35\\\\ 309.8+0.0   & & 3.6 & \\scrhi\\ absorption & 25 & S & \\nodata &\\nodata  & \\cxo-AO3 \\\\ 315.4$-$2.3 &RCW 86 & 2.8 & Optical lines & 42 & S & \\nodata &\\nodata  & \\cxo-AO3 & 36\\\\ 327.4+0.4 & Kes 27 & 4.3 & \\scrhi\\ absorption/interact. & 30 & S & \\nodata &\\nodata  & \\nodata & \\\\ 327.6+14.6 & SN~1006 & 2.2 & Spectra/proper motion & 30 & S & Ia & \\nodata & \\nodata &  37,38\\\\[0.1in] 330.0+15.0 & Lupus Loop & 0.8 & \\scrhi\\ & 180 & S & \\nodata &\\nodata  & \\cxo-AO4&  \\\\ 332.4$-$0.4 & RCW 103 & 3.3 & \\scrhi\\ absorption & 10 & S & CCO & yes &\\nodata& 39 \\\\ 343.1$-$2.3 & & 2.5 &\\scrhi\\ absorption  & 32 & C & pulsar &yes& \\nodata& 40,41\\\\ 354.1+0.1 & & 5 & recombination lines & 15 & C & pulsar? & no & \\nodata& 42,43\\\\ \\enddata \\tablenotetext{a}{Major axis.} \\tablenotetext{b}{Types are: S (shell), C (composite), and F (filled, or PWN), as determined by \\citet{green01}.  A question mark indicates that the type is poorly determined.} \\tablenotetext{c}{Indicates if a central source is known, or if the SNR is thought to be of type Ia.  If a central source is known and it falls into one of the classes from \\S~\\ref{sec:class}, it is labeled accordingly.} \\tablenotetext{d}{Indicates whether the central source has been detected in X-rays.  See Table~\\ref{tab:psrs}.} \\tablenotetext{e}{Refers to X-ray samples of shell SNRs described in \\S~\\ref{sec:Chandra-survey}.} \\tablenotetext{f}{References deal only with the central source or Ia classification. General SNR properties were taken from \\citet{green01} and references therein.} \\tablenotetext{g}{Has already been observed with \\chandra.}  \\tablerefs{1: \\citet{kt99}; 2: \\citet{fk91}; 3: \\citet{mjk+91}; 4: \\citet{yswk00}; 5: \\citet{ttdm97}; 6: \\citet*{wcd91}; 7: \\citet{hhs+97}; 8: \\citet{pks02}; 9: \\citet*{cgc75}; 10: \\citet{wwgs83}; 11: \\citet{kcb+88}; 12: \\citet{mgb+02}; 13: \\citet{mtt+98}; 14: \\citet{bkc+96}; 15: \\citet{t99}; 16: \\citet{kpha93}; 17: \\citet*{bbt96}; 18: \\citet*{frs93}; 19: \\citet*{sss95}; 20: \\citet{ssb+97}; 21: \\citet{szh+04}; 22: \\citet{baade45}; 23: \\citet{mss+02}; 24: \\citet{csl+02}; 25: \\citet{acj+96}; 26: \\citet{rn03}; 27: \\citet{klh+03}; 28: \\citet{ccl+69}; 29: \\citet{ocw+01}; 30: \\citet{tbb+03}; 31: \\citet*{pzt99}; 32: \\citet*{ztp99}; 33: \\citet{lvm68}; 34: \\citet{hb84}; 35: \\citet{zpst00}; 36: \\citet{vbdk00}; 37: \\citet{fwlh88}; 38: \\citet{apg01}; 39: \\citet*{gph97}; 40: \\citet*{mop93}; 41: \\citet*{bbt95}; 42: \\citet{fgw94}; 43: \\citet{bt97}. } \\end{deluxetable} ",
        "conclusions": "\\label{sec:conclusions} There are 45 known SNRs that are reliably within 5 kpc of the Sun.  Most of these SNRs are expected to contain central neutron stars: $\\lsim 20$\\% are expected to result from Type Ia SNe and thus not contain a central compact source, while  $\\sim 20$\\% (dependent on the stellar initial mass function, the limiting mass for black holes, and binary evolution; \\citealt{hfw+03}) are expected to host a central black hole that may not be easily identified as such.  Thanks to the persistent efforts by astronomers over the past four decades, central sources have been detected in the X-ray and/or radio bands in 18 of these SNRs, and three have been identified as probable Type~Ia SNe. In some cases, only a centrally located PWN is detected, but in those cases (i.e.\\ IC~443) it is reasonably assumed that the PWN is powered by a central compact source.  We have begun a program of searching for compact sources in the remaining 23 SNRs. The program has been motivated by the discovery of a point-like X-ray source at the very center of the youngest known Galactic SNR, Cas~A.  The program has two observational components: imaging with \\chandra\\ or \\xmm\\ in the X-ray band followed by ground-based optical/IR followup. The latter is essential because of the high incidence of of interlopers (foreground stars and background galaxies). Such a comprehensive program is possible thanks to the astrometric accuracy of the X-ray missions combined with high sensitivity.  In this paper we report on \\chandra\\ ACIS-I observations of four shell remnants (\\Ga, \\Gb, \\Gc, \\Gd).  We undertook optical and IR observations of every X-ray source detected with more than 10 counts. For all detections, we found, within astrometric errors, a counterpart in  one or more bands. These counterparts were consistent with either foreground (stars) or background (galaxies) sources.  In particular, the X-ray flux and/or the X-ray to optical (IR) ratio of the detected X-ray sources were not as extreme as all known classes of neutron stars: accreting neutron stars, radio pulsars, AXPs and SGRs.  In \\S~\\ref{sec:limits} we discuss reasons why standard neutron stars were not found in these SNRs: they could have fallen in the gap between the chips, they could have escaped our field of view due to a very high velocities, they could be undetectable black holes, or they could not exist owing to the SNRs being the results of Type Ia explosions.  All of these scenarios are unlikely for a single source, and even more so for all four, but are technically possible.  If, on the other hand, these scenarios do not apply, then four remnants contain neutron stars that are fainter than our X-ray detection limit (typically, $L_X \\lsim 10^{31}\\,$erg s$^{-1}$ in the 0.5--10~keV band).  We now consider this last (and most interesting) possibility. In the absence of other forms of energy generation (accretion, rotation power, magnetic field decay) the minimum X-ray flux one expects is set by the cooling of the neutron star.  From Section~\\ref{sec:cool} and Figure~\\ref{fig:cool}  we immediately see that the central neutron stars in these four remnants must be cooler than those present for example in the similarly-aged Puppis~A, PKS~1205$-$51/52 and RCW~103.  Our knowledge of the physics of cooling is by no means firm. There is considerable debate among theorists as to which of the multitude of physical processes can significantly affect the cooling output and as to which of the physical parameters (mass, rotation rate, magnetic field) controls these processes \\citep{ygkp02}.  Nonetheless, there is agreement that more massive neutron stars (with their larger mean densities) cool more rapidly than those with smaller mass; this expectation is illustrated in Figure~\\ref{fig:cool}.  Thus our upper limits can be made consistent with the cooling possibility provided the central neutron stars in these four SNRs are more massive than known cooling neutron stars. Indeed, the known examples of radio-quiet objects could well result from a strong selection effect, namely the earlier X-ray observations by \\textit{Einstein} and \\rosat\\ detected the warmer cooling neutron stars (ignoring the neutron stars detected because of non-thermal emission).  The existing data may already hint at a parameter affecting cooling, as exemplified by PSR~J0205+6449 \\citep{shm02} and the Vela pulsar \\citep{pzs+01}, but also possibly by PSR~B1853+01 \\citep*{pks02} and RX~J0007.0+7302 \\citep{szh+04}.  Of course, we also do not see rotation-powered pulsars such as the majority of the objects in Table~\\ref{tab:psrs}.  Recent observations are finding pulsars with lower radio luminosities and values of $\\dot E$ than ever before, and our limits would only be consistent with these newer sources.  One might then ask why we see neither a standard cooling neutron star nor a standard active pulsar, assuming that there is no  intrinsic correlation between these properties.  It is possible that there truly are no neutron stars in these SNRs, allowing one to speculate wildly about what actually is there.   In subsequent papers we plan to report the X-ray observations and ground-based follow up of the remaining 19 SNRs.  If no central sources are identified then the hypothesis proposed here, namely that there is a parameter that determines cooling of neutron star will be strengthened.  The simplest (and most reasonable) suggestion is that this second parameter is the mass of neutron stars, so that cooling observations could be used to ``weigh'' isolated neutron stars (as discussed by \\citealt*{khy01}). This hypothesis is at odds with the strong clustering of binary neutron star masses \\citep{tc99}. We note, though, that mass determination is only possible for neutron stars in compact binary systems where significant interaction with the companion may have taken place. Thus it is possible that neutron stars resulting from single stars are systematically more massive than those which evolve in compact binary systems (and have gained mass through accretion)."
    },
    "0403/astro-ph0403639_arXiv.txt": {
        "abstract": "We present new spectra of WR 140 (HD 193793) in the $JHK$ bands with some covering the 1.083-$\\mu$m He\\,{\\sc i} emission line at higher resolution, observed between 2000 October and 2003 May to cover its 2001 periastron passage and maximum colliding-wind activity.  The WC7 + O4-5 spectroscopic binary WR 140 is the prototype of colliding-wind, episodic dust-making Wolf-Rayet systems which also show strong variations in radio and X-ray emission. The $JHK$ spectra showed changes in continuum and in the equivalent widths of the WC emission lines, consistent with formation of dust starting between 2001 January 3 and March 26 (orbital phases 0.989 and 0.017) and its subsequent fading and cooling.  The 1.083-$\\mu$m He\\,{\\sc i} line has a P-Cygni profile which showed variations in both absorption and emission components as WR 140 went through periastron passage. The variation of the absorption component of the profile yielded tight constraints on the geometry of the wind-collision region, giving $\\theta = 50^{\\circ}\\pm8^{\\circ}$ for the opening semi-angle of the interaction region `cone', indicating a wind-momentum ratio $(\\dot{M} v_{\\infty})_{\\rm O} / (\\dot{M} v_{\\infty})_{\\rm WR} = 0.10$, about three times larger than previously believed. As the system approached periastron, the normally flat-topped emission component of 1.083-$\\mu$m line profile showed the appearance of a significant sub-peak. The movement of the sub-peak across the profile was seen to be consistent with its formation in wind material flowing along the contact discontinuity between the WC7 and O4-5 stellar winds and the changing orientation of the colliding wind region as the stars moved in their orbits. The flux carried in the sub-peak was significant, exceeding the X-ray fluxes measured at previous periastron passages. This additional source of radiative cooling of the shock-heated gas probably causes it to depart from being adiabatic around periastron passage, thereby accounting for the departure of the X-ray flux from its previously expected $1/d$-dependency. ",
        "introduction": "Stars with  M$_{\\ast} \\geq$ 30M$_{\\odot}$ go through a Wolf-Rayet (WR) phase in their evolution, during which they undergo large scale mass loss ($\\dot{M} \\sim$ 10$^{-5}$ $M_{\\odot}$ year$^{-1}$).  Consequently they display broad emission line spectra and large infrared excesses originating in their accelerated stellar winds (wind terminal velocity v$_{\\infty}$=750--5000 km s$^{-1}$). They are post-main sequence stars which have lost most of their hydrogen envelopes through stellar winds in the pre-WR stage. Based on the emission lines seen, they are classified as WN stars (lines of He, N and H), WC and WO stars (lines of He, C and O). Late WC stars (WC7--9) show evidence of circumstellar dust, mainly of amorphous carbon (Williams 1995).  Some of them are persistent dust makers which show continuous dust formation and others are episodic dust makers, where we see periodic rises in their IR flux due to episodic formation of dust. It is believed, now, that episodic dust making WC stars are massive binary systems in which the stellar winds of two massive stars interact. WR 140 (HD 193793) is the prototype of colliding-wind, episodic dust forming WC stars (Williams 1995).  From the photometry of the infrared maxima in 1977 and 1985 and earlier data, Williams et al. (1987) found that the occurence of the IR maxima and the dust formation were periodic and derived a period of 7.9 years. Studying the existing radial velocity data, from which no orbit had previously been derived, they showed that WR 140 was indeed a spectroscopic binary with high eccentricity. This was extended to a multi-wavelength study, including 1.65--12.6 $\\mu$m light curves, radio and X-ray data by Williams et al. (1990, hereafter W90), who derived a photometric period $P=2900 \\pm 10$ days (7.94 y) and a spectroscopic orbit from published radial velocities giving elements $e=0.84\\pm0.04$, $\\omega$=32$^{o}\\pm$8$^{o}$ and T$_{0}$=JD 2446160$\\pm$29 (1985.26$\\pm$0.08)). The system formed dust for about four months, 4\\% of the period, during periastron passage. This is perhaps one of the least well understood of the phenomena shown by WR 140.  The most recent periastron passage and dust-formation episode occurred in early 2001, for which a multi-wavelength observing campaign was planned and undertaken.  Marchenko et al. (2003) carried out optical photometric and spectroscopic observations and derived a spectroscopic period $P=2899\\pm1.3$ days, in excellent agreement with the photometric period and new orbital elements, $e$=0.881$\\pm$0.005, $\\omega$=46.7$^{o}\\pm$1.6$^{o}$ and T$_{0}$=JD 2446147.4$\\pm$3.7, which we use in this paper.  They found that at phases very close to periastron,  the profiles of the 5696\\AA  $\\>\\>$C\\,{\\sc iii} and 5876\\AA $\\>\\>$He\\,{\\sc i} lines developed extra emission components at their `blue' ends which rapidly moved to the red.  Panov \\& Dinko (2002) observed a  minimum in UBV photometric bands in 2001 June (phase = 0.038-0.046). This minimum was  observed at the same phase, but  deeper than the optical minimum observed in 1993 (Panov, Altmann \\& Seggewiss 2000). By 2001 July, the optical light almost reached the pre-eclipse level. They interpreted this eclipse-like phenomenon as due to the carbon dust envelope, the formation of which was triggered at the periastron passage. The difference in the rising part of the eclipse light curves of 1993 and 2001 could be due to the differences in clumping in the expanding dust shell from cycle to cycle.  Aperture masking interferometric observations by Monnier, Tuthill \\& Danchi (2002) using the Keck Telescope showed the first direct images of dust in WR 140. Comparison of their pre-periastron observation in 1999 June and post-periastron observation in 2001 July at 2.2 $\\mu$m shows expanding clumpy dust. They resolved five clumps in the shell,  expanding with velocities 0.46--1.29 milliarcsec. day$^{-1}$.  \\begin{table} \\begin{centering} \\caption{Log of Spectroscopic Observations} \\label{obslog} \\begin{tabular}{@{}lllll} \\hline \\multicolumn{5}{c}{\\underline{UKIRT + CGS4 (150 lines/mm grating, R=4700)}}\\\\ UT Date   \t&JD\t\t&Bands/\t\t\t&Phase  &$d/a$\t    \\\\ {\\tiny yyyymmdd.dddd}  \t&2450000+\t&$\\lambda (\\mu m)$\t&\t&\t\t\\\\ \\hline 20001013.3495  &1830.85\t&1.07,1.1 \t\t&0.9605\t&0.5092     \\\\ 20001225.1921  &1903.69\t&1.083\t\t\t&0.9856\t&0.2506     \\\\ 20001226.2004  &1904.7\t&1.083\t\t\t&0.9860 &0.2465     \\\\ 20010318.654   &1987.15\t&1.083\t\t\t&1.0144 &0.2507     \\\\ 20010331.6208  &2000.12\t&1.083 \t\t\t&1.0189 &0.3018     \\\\ \\hline  \\multicolumn{5}{c}{\\underline{UKIRT + CGS4 (40 lines/mm grating)}} \\\\ \\multicolumn{5}{c}{\\underline{R=800 in J band, R=400 in H and K bands}}\\\\ \\hline 20010331.6618  &2000.16\t&J,H,K\t     &0.0189\t&0.3020\t\\\\ 20010428.6164  &2028.12\t&J,K\t     &0.0285\t&0.4045\t\\\\ 20010521.5254  &2051.03\t&J,H,K\t     &0.0364    &0.4810\t\\\\ 20010605.4691  &2065.97\t&J,K\t     &0.0416    &0.5277 \\\\ 20010613.5553  &2074.06\t&J           &0.0444    &0.5521 \\\\ 20010704.6004  &2095.1\t&J,H,K       &0.0516    &0.6129 \\\\ 20010812.2456  &2133.75\t&J,H,K       &0.0650    &0.7156 \\\\ 20010909.3325  &2161.83\t&J,H,K       &0.0747    &0.7841 \\\\ 20011006.2138  &2188.71\t&J,H,K       &0.0839    &0.8457 \\\\ 20011122.2333  &2235.73\t&J,K         &0.1001    &0.9452 \\\\ 20011226.1985  &2269.7\t&J,K         &0.1119    &1.0115 \\\\ 20020403.6638  &2368.16\t&J,H,K       &0.1458    &1.1816 \\\\ 20020630.3636  &2455.86\t&J,H,K\t     &0.1761    &1.3104 \\\\ 20020717.5517  &2473.05\t&J,H,K       &0.1820    &1.3335 \\\\ \\hline \\multicolumn{5}{c}{\\underline{UKIRT + UIST (IJ, HK \\& short$\\_$J grisms)}} \\\\ \\multicolumn{5}{c}{\\underline{R=3000-short$\\_$J, R=950-IJ \\& R=800-1000-HK grism}}\\\\ 20030524.5496  &2784.05\t&short$\\_$J  &0.2893    &1.6542 \\\\ 20030524.5640  &2784.06\t&IJ\t     &0.2893\t&1.6542\t\\\\ 20030524.5740  &2784.07\t&HK          &0.2893    &1.6542 \\\\ \\hline \\multicolumn{5}{c}{\\underline{Mt. Abu (R=1000)}}\\\\ 19980508.9457  &0942.45\t&JHK\t     &0.6540\t&1.7620 \\\\ 20001221.575   &1900.08\t&JH          &0.9844    &0.2650\t\\\\ 20001222.5658  &1901.07\t&J           &0.9847\t&0.2610 \\\\ 20001223.5958  &1902.10 \t&J           &0.9851    &0.2570 \\\\ 20010102.5861  &1912.09\t&J           &0.9885    &0.2168 \\\\ 20010103.6090  &1913.11\t&J           &0.9889    &0.2127 \\\\ 20010326.9995  &1995.50\t&J           &0.0173    &0.2836 \\\\ 20010328.9988  &1997.50\t&J           &0.0180    &0.2914 \\\\ 20010402.9167  &2002.42\t&J           &0.0197    &0.3106 \\\\ 20010403.8743  &2003.37\t&J           &0.0200    &0.3145 \\\\ \\hline \\end{tabular} \\end{centering} \\end{table}  The programme of near-infrared spectroscopy described in the present paper was undertaken to search for changes in the infrared line spectrum, particularly the profile of the He\\,{\\sc i} line at 1.083 $\\mu$m, around the time of maximum colliding wind activity. Previous 1-micron and $JHK$ spectroscopy, e.g. Vreux, Andrillat \\& Bi\\'emont (1990), Eenens, Williams \\& Wade (1991), and observations of the 1.083-$\\mu$m He\\,{\\sc i} line profile (Eenens \\& Williams 1994), were made at phases far from periastron. In each observation, the 1.083-$\\mu$m emission-line profile showed a flat top, characteristic of formation in the asymptotic region of the Wolf-Rayet wind, with no evidence of a sub-peak which could be formed in a wind collision region (Stevens \\& Howarth 1999). Nevertheless, given the very high eccentricity of the binary system and short duration of maximum colliding-wind activity such as dust formation, we considered it worth reobserving the 1.083-$\\mu$m profile as close to the periastron as possible. ",
        "conclusions": "Infrared spectroscopy of WR 140 during and after its 2001 periastron passage showed dilution of the emission lines consistent with the formation of circumstellar dust between 2001 January 3 and March 26, phases 0.989 and 0.017. The He\\,{\\sc i} 1.083-$\\mu$m line showed a strong P-Cygni profile. Variation of the strength of the absorption component measured from high-resolution spectra was interpreted in terms of the passage of different parts of the colliding wind structure in our line of sight to the two stars. This allowed us to set fairly tight limits on the opening half angle of the `cone' modelling the interaction region: $\\theta = 50^{\\circ}\\pm 8^{\\circ}$. This, in turn, indicated a wind-momentum ratio $\\eta = 0.10$, signifcantly higher than the $\\eta = 0.034$ found from the mass-loss rates and terminal velocities of the WC7 and O components. We suggested  that the mass-loss rate of the WC7 star, $5.7 \\times 10^{-5} M_{\\odot}y^{-1}$ derived from the 5-GHz flux density at radio minimum was overestimated because the wind is clumped.  At phases ($\\Delta\\phi \\simeq 0.1P$) close to periastron passage, the normally flat-topped emission component of the 1.083-$\\mu$m line showed the appearance of a sub-peak which moved from the blue to the red end of the profile during periastron. The evolution of the position and width of the sub-peak were found to be reasonably conistent with the flow of the emitting material along the surface of the interaction region modelled by a cone having $\\theta = 50^{\\circ}$. The maximum radiative flux in the sub-peaks is greater than the 1--10 keV X-ray flux at periastron and is a significant source of cooling of the shocked WC wind. We suggest that the shocked WC wind is not totally adiabatic near periastron, thereby accounting for the departure of the X-ray flux from the previously expected $1/d$-dependence and, as suggested by Marchenko et al., allowing dust formation."
    },
    "0403/astro-ph0403255_arXiv.txt": {
        "abstract": "Galaxy-galaxy lensing has emerged as a powerful probe of the dark matter halos of galaxies, but is subject to contamination if intrinsically aligned satellites of the lens galaxy are used as part of the source sample.  We present a measurement of this intrinsic shear using 200,747 lens galaxies from the Sloan Digital Sky Survey (SDSS) spectroscopic sample and a sample of satellites selected using photometric redshifts. The mean intrinsic shear at transverse separations of 30--446$h^{-1}$~kpc is constrained to be $-0.0062<\\Delta\\gamma<+0.0066$ (99.9 per cent confidence, including identified systematics), which limits contamination of the galaxy-galaxy lensing signal to at most $\\sim 15$ per cent on these scales.  We present these limits as a function of transverse separation and lens luminosity.  We furthermore investigate shear calibration biases in the SDSS, which can also affect galaxy-galaxy lensing, and conclude that the shear amplitude is calibrated to better than 18 per cent.  This includes noise-induced calibration biases in the ellipticity, which are small for the sample considered here, but which can be more important if low signal-to-noise or poorly resolved source galaxies are used. ",
        "introduction": "Weak gravitational lensing has emerged as a powerful tool for directly measuring the matter distribution in the universe (e.g. \\citealt{2001PhR...340..291B}; \\citealt{2003ARA&A..41..645R}). One of its applications has been the measurement of the projected matter density of galaxies (\\citealt{1996ApJ...466..623B}; \\citealt{1998ApJ...503..531H}; \\citealt{2000AJ....120.1198F}; \\citealt{2001ApJ...551..643S};  \\citealt{2001astro.ph..8013M}; \\citealt{2002MNRAS.335..311G}; \\citealt{2003MNRAS.340..609H}; \\citealt{2003astro.ph.12036S}; \\citealt{2004ApJ...606...67H}), groups, and clusters (\\citealt{2001ApJ...548L...5H}; \\citealt{1998ApJ...504..636H}; \\citealt{2001ApJ...554..881S}).  In these measurements, the observable quantity is the tangential shear $\\gamma_t$ of distant source galaxies induced by the lens. It is however possible that some of the ``source'' galaxies whose shears are used to measure lensing may in fact be physically associated with the lens; in this case they may be intrinsically aligned with the lens, producing a false lensing signal. This false signal may be quantified by the ``intrinsic shear'' $\\Delta\\gamma$, which is the spurious estimate of the tangential shear that would be obtained by applying a shear estimator to a population of physically associated galaxies (see Appendix~\\ref{app:dgdef} for a more precise definition).  This intrinsic shear is essentially a type of galaxy density-shear correlation, i.e. it measures the correlation between the shears of some galaxies and the positions of others; it is distinct from the shear-shear correlations that contaminate the gravitational shear autocorrelation measured by lensing surveys in the field (\\citealt{2000A&A...358...30V}; \\citealt{2000MNRAS.318..625B}; \\citealt{2001ApJ...552L..85R}; \\citealt{2002ApJ...572...55H}; \\citealt{2002A&A...393..369V}; \\citealt{2003AJ....125.1014J}; \\citealt{2003MNRAS.341..100B}).  Several investigators have investigated this potential contaminant of the lensing signal.  \\citet{2001ApJ...555..106L} measured the density-shear correlation using 12,122 spiral galaxies.  These authors were primarily interested in the rotation axes of galaxies and their possible use as tracers of the large-scale tidal field (and hence the cosmic density perturbations; \\citealt{2000ApJ...532L...5L}), and so they measure position angles rather than shear.  A crude (and model-dependent) conversion of their measurements into an intrinsic shear suggests $\\Delta\\gamma = -0.0037\\pm 0.0025$ at separations $r_{3D}=500h^{-1}$~kpc (see Appendix~\\ref{app:pac}).  \\citet{2002AJ....124..733B} used 1,819 satellite galaxies selected spectroscopically from the Two Degree Field (2dF) Galaxy Redshift Survey \\citep{2001MNRAS.328.1039C} with ellipticities measured in the Automatic Plate Measuring (APM) survey \\citep{1990MNRAS.243..692M} to set an upper limit on the intrinsic shear of the satellite galaxies of $|\\Delta\\gamma|\\le 0.01$ (at 95 per cent confidence)\\footnote{\\citet{2002AJ....124..733B} reported an upper limit on the ellipticity of 0.02, which we have converted here to a shear.} for primary-satellite pairs separated by a transverse distance of $<350h^{-1}$~kpc.  However, this upper limit is not presented as a function of transverse separation or lens luminosity. Circularization of the galaxy images due to the APM point-spread function (PSF) may weaken the upper limit slightly (by $\\sim 10$--$20$ per cent; \\citealt{2002AJ....124..733B}) but this does not alter the conclusion that intrinsic alignment contamination of the galaxy-galaxy lensing signal observed by \\citet{2000AJ....120.1198F} and \\citet{2001astro.ph..8013M} is limited to $\\la 20$ per cent.  There have been other observational studies of intrinsic alignments, including reported detections of higher-order density-shear statistics \\citep{2002ApJ...567L.111L} and shear-shear correlations (\\citealt{2000ApJ...543L.107P}; \\citealt{2002MNRAS.333..501B}); unfortunately there is no simple conversion from these measurements into an intrinsic shear $\\Delta\\gamma$.  Theoretical work on galaxy intrinsic alignments has been motivated both by inherent interest and (more recently) by contamination of cosmic shear surveys for which the intrinsic-alignment parameter of interest is the shear-shear correlation instead of the density-shear correlation.  A number of estimates of the shear-shear correlations -- both in $N$-body simulations and in analytical models -- have been made (\\citealt{2000ApJ...545..561C}; \\citealt{2000MNRAS.319..649H}; \\citealt{2000ApJ...532L...5L}; \\citealt{2001ApJ...555..106L}; \\citealt{2001MNRAS.320L...7C}; \\citealt{2001ApJ...559..552C}; \\citealt{2002astro.ph..5212H}; \\citealt{2002MNRAS.335L..89J}).  Usually the galaxy is approximated as a disk aligned perpendicular to the halo angular momentum vector (for spirals) or as an ellipsoid homologous with the ellipticity of its halo (for ellipticals), although comparison to observations \\citep{2003astro.ph.10174H} and simulations of galaxy formation \\citep{2002ApJ...576...21V} suggest that this picture is too simple and overestimates the shear-shear correlations.  Of these authors, only \\citet{2001ApJ...555..106L} and \\citet{2002astro.ph..5212H} estimate the density-shear correlation (and even then only for spirals).  Using the rough conversion of Appendix~\\ref{app:pac}, the \\citet{2001ApJ...555..106L} prediction corresponds to $\\Delta\\gamma=-0.004$ for the spiral sources, the $-$ sign indicating radial alignment.  \\citet{2002astro.ph..5212H} present some predictions for the two-dimensional (projected) density-ellipticity correlation as a function of angular separation, but do not provide a numerical estimate $\\Delta\\gamma(r)$.  In this paper, we use SDSS spectroscopic galaxies (as lenses) and photometric galaxies (as sources) in a luminosity-dependent study of the shear around lens galaxies due to intrinsic alignments.  The large number of galaxies in the SDSS photometric sample allows tighter constraints on $\\Delta\\gamma$ than were obtained by \\citet{2002AJ....124..733B}.  It also provides the statistical power to compute these upper limits as a function of transverse separation and lens luminosity.  To make the interpretation for weak lensing as simple as possible, we directly compute $\\Delta\\gamma$ using a lensing estimator (including PSF corrections) rather than the position-angle statistics of \\citet{2001ApJ...555..106L} or the unweighted moments computed by the APM survey.  Correct calibration of the shear estimator is necessary in order to interpret either an intrinsic alignment or a weak lensing-induced shear, particularly if the shear is detected at high signal-to-noise.  Errors in shear calibration can come from a variety of effects such as incomplete PSF correction, selection effects, noise-induced biases, uncertainties in PSF reconstruction, and incomplete knowledge of the ellipticity distribution of the source population (\\citealt{2000ApJ...537..555K}; \\citealt{2002AJ....123..583B}; \\citealt{2003MNRAS.343..459H}).  Since we do not detect an intrinsic alignment signal, the only effect of calibration uncertainty is a minor degradation of our upper limits on $\\Delta\\gamma$.  Nevertheless, the calibration results presented here are of direct interest for ongoing galaxy-galaxy weak lensing studies with the SDSS; calibration errors at the $\\le$16--18 per cent level found here are potentially important in those cases where the lensing signal is detected at $\\ga 6\\sigma$.  The noise-induced calibration biases have not been explicitly studied previously; while we find that this effect is insignificant for our sample of relatively bright ($r<21$) galaxies, the calibration bias can be large for galaxies detected at low signal-to-noise ratio, especially if they are also not well-resolved.  This paper is organized as follows: Sec.~\\ref{sec:lens} includes a description of the lens catalog, and Sec.~\\ref{sec:source} includes discussion of the source catalog, including details of the apparent shear measurement, the associated calibration uncertainties, and the photometric redshifts.  In Sec.~\\ref{sec:uncertainty}, we describe the methods used to compute uncertainties on the shear. Sec.~\\ref{sec:results} includes results for intrinsic and gravitational shear, and the results of tests for systematic errors affecting these quantities.  We conclude in Sec.~\\ref{sec:discussion} with a discussion of the implications of these results.  The noise-induced calibration bias derivation is presented in Appendix~\\ref{app:kn}.  A note about the cosmological model and units used in this paper: we compute the proper transverse separation of the lens and source using the physical angular diameter distance $d_A(z_l)$ to the lens.  This distance is determined assuming a flat $\\Lambda$CDM universe with $\\Omega_m=0.27$ and $\\Omega_\\Lambda=0.73$, in accord with cosmological parameter determinations from the {\\slshape Wilkinson Microwave Anisotropy Probe} \\citep{2003ApJS..148..175S}.  We present distances in $h^{-1}$~kpc, where $h$ is the reduced Hubble parameter: $H_0=100h$~km/s/Mpc.  When appropriately scaled by $h$, the angular diameter distance is only weakly sensitive even to extreme variations in cosmology: for a lens at redshift $z=0.3$ (among the most distant in our sample; see Fig.~\\ref{fig:zdist}), we find that $d_A(z_l)$ is reduced by 7 per cent if we switch from $\\Lambda$CDM to an open universe ($\\Omega_m=0.27$, $\\Omega_\\Lambda=0$), and by 12 per cent if we switch to an Einstein-de Sitter universe ($\\Omega_m=1$, $\\Omega_\\Lambda=0$).  Given that we see no detection of intrinsic alignments, the dependence of our results on $\\Omega_m$ and $\\Omega_\\Lambda$ is negligible for the purpose of constraining contamination of the weak lensing signal. ",
        "conclusions": "\\label{sec:discussion}  In this paper, we have set stringent new limits on intrinsic alignments as a contaminant of the galaxy-galaxy lensing signal.  For example, averaging together all lens galaxies at radii out to 446$h^{-1}$~kpc we derive $-0.0062<\\Delta\\gamma<+0.0066$ (99.9 per cent confidence including identified systematics; see Table~\\ref{tab:shearlimit}) or $-0.0018<\\Delta\\gamma<+0.0021$ (95 per cent confidence, statistical errors only).  The statistical constraint is a factor of $\\sim 5$ ($\\sim 1.5$ for our more conservative error bar) tighter than $|\\Delta\\gamma|\\le 0.01$ (95 per cent confidence) as obtained from APM/2dF data by \\citet{2002AJ....124..733B}.  It also extends to smaller radii than the measurements of nearby spirals by \\citet{2001ApJ...555..106L}. Furthermore, our study of the luminosity and radius dependence of the intrinsic shear is useful for weak lensing studies that divide lens galaxies into luminosity subsamples.  Our limits have been obtained by using ``lens'' galaxies with spectroscopic redshifts and ``source'' galaxies with photometric redshifts, a different strategy than that selected by previous observers (\\citealt{2001ApJ...555..106L}; \\citealt{2002AJ....124..733B}) who used lens-source pairs selected entirely on the basis of spectroscopic redshift information. This strategy has dramatically increased our sample size and reduced our statistical uncertainties for the brighter lenses ($\\mrp<-20$) because these are typically at redshift $z>0.1$, where most of their satellites are too faint to be targeted by SDSS spectroscopy.  This method does however have disadvantages. For our fainter-luminosity lenses, the photometric-redshift source method gives very large error bars (Table~\\ref{tab:shearlimit}) because there are few physically associated lens-source pairs and hence any intrinsic alignment signal will be drowned out by the large number of unassociated pairs.  The use of photometrically selected pairs also comes with the caveat that the limits on $\\Delta\\gamma$ could be circumvented by a small class of galaxies that have strong intrinsic alignments but always produce incorrect photometric redshifts.  It is thus recommended that the values in Table~\\ref{tab:shearlimit} with redshift slice half-width $\\epsilon=0.1$ be used to constrain intrinsic alignments in galaxy-galaxy lensing studies, since in this case only about 20 per cent of the physically associated source galaxies are scattered out of the sample by photometric redshift errors.  Despite the recent interest in intrinsic alignments as a contaminant of cosmic shear (i.e. shear-shear correlations), there has been comparatively little theoretical work on the density-shear correlation. \\citet{2001ApJ...555..106L} do provide a prediction for the density-shape correlation for spiral ``source'' galaxies, which we convert (see Appendix~\\ref{app:pac}) to $\\Delta\\gamma=-0.004$ at small radii $r<1h^{-1}$~Mpc.  This is marginally consistent with our results (it is within our more conservative confidence interval of $-0.0062<\\Delta\\gamma<+0.0066$), although the comparison is imperfect as the model applies only to spiral sources.  More generally, measurements of the density-shear correlation $\\Delta\\gamma$ provide a constraint on intrinsic alignment models in addition to that provided by the shear autocorrelation measurements (\\citealt{2000ApJ...543L.107P}; \\citealt{2002MNRAS.333..501B}).  It would thus be useful to compute the predictions of other intrinsic alignment models for $\\Delta\\gamma$ and determine which models agree with the intrinsic shear measurements presented here and in \\citet{2002AJ....124..733B}.  We conclude by roughly estimating the intrinsic alignment contamination to the gravitational lensing shear, $(1-R_p^{-1})\\Delta\\gamma$.  For example, if our source catalog restricted to $z_s>z_l+0.1$ were used in a lensing study of the halos of SDSS spectroscopic galaxies, we would have a contamination fraction $1-R_p^{-1}\\sim 0.01$ in the $100<r<300 h^{-1}$~kpc bin and intrinsic alignment contamination $|\\Delta\\gamma|<0.0095$ (see Table~\\ref{tab:shearlimit}), corresponding to a contamination to the gravitational lensing shear at the level of $<9.5\\times 10^{-5}$.  This is approximately equal to the size of the $1\\sigma$ error bars at these radii for our full sample ($1.0\\times 10^{-4}$; see Fig.~\\ref{fig:e.and.bmode}) and a factor of six less than the shear signal, $\\sim 5.8\\times 10^{-4}$. Slightly weaker constraints are placed on the intrinsic alignments around the brighter lenses, but in these cases the lensing signal is a factor of several higher, e.g. for the $-21>M_{r}>-22$ magnitude range we have $|\\Delta\\gamma|<0.0124$ at $100<r<300h^{-1}$~kpc, $1-R_p^{-1}\\sim 0.02$, and shear signal $\\gamma_t= 2\\times 10^{-3}$, so the contamination is at worst about a factor of eight smaller than the signal. These results indicate at 99.9 per cent confidence that, except possibly for the faintest lens galaxies, intrinsic alignments are only contaminating the sub-Mpc scale galaxy-galaxy lensing signal at the $\\la 15$ per cent level."
    },
    "0403/astro-ph0403580_arXiv.txt": {
        "abstract": "{ We review various prescriptions which have been proposed for the turbulent transport of matter and angular momentum in differentially rotating stellar radiation zones. A new prescription is presented for the horizontal transport associated with the anisotropic shear turbulence which is  produced by the differential rotation in latitude; this `$\\beta$-viscosity' is drawn from torque measurements in the classical Couette-Taylor experiment (Richard \\& Zahn 1999). Its implementation in a stellar evolution code leads to enhanced mixing, as illustrated by models of a rotating main-sequence star of 1.5 solar mass. ",
        "introduction": "In standard  models of stellar interiors, radiation zones, which are convectively stable, are postulated to be without motion other than rotation, except in the vicinity of convection zones, where often convective overshoot is accounted for in a crude way. But it is well known, since Eddington (1925) and Vogt (1925), that in rotating stars the centrifugal force breaks the radiative equilibrium, and thus causes a slow large-scale circulation, which transports matter and angular momentum. A similar circulation is induced locally in the vicinity of a differentially rotating convection zone, in the so-called tachocline. As a result, radiation zones rotate non-uniformly, and therefore they are liable to various instabilities which may generate turbulence.  This turbulence contributes to the transport of matter and angular momentum, and to model these in a stellar evolution code, one needs reliable prescriptions for the turbulent transport. To derive such prescriptions remains today one of the most challenging problems, not only in stellar physics, but also in oceanography and in atmospheric sciences. The reason is that turbulence is still poorly understood, particularly when it departs from ideal conditions: homogeneity, isotropy, stationarity. ``Turbulence is an order of magnitude subject; there is no rigorous theory of turbulence'' warned M. Hoffert at the Aspen Global Change Institute in 1997. And he concluded:  ``The way turbulence is represented in current (global circulation) models is not realistic and this has important ramifications for the climate sensitivity predicted by these models.''  We could use the same words to describe the situation when modeling stellar interiors, replacing ``climate sensitivity'' by ``evolutionary sequences''. Therefore it takes some courage to - nevertheless - engage in this task, as we shall attempt in this article. We shall restrict our scope on turbulence produced by shear instabilities, which are likely to be the most powerful, at least in a non-magnetic star.  Why do we believe that stellar radiation zones are turbulent? Some velocity profiles are known to be linearly unstable, as soon as the Reynolds number characterizing the flow exceeds some critical value. (We recall that this number involves the typical size $L$ and and the typical velocity contrast $V$ of the flow: $Re = V L/ \\nu$, where $\\nu$ is the kinematic viscosity; it measures the relative strength of advection and viscous transport.) But flows that are linearly stable are found turbulent in laboratory experiments, again above some critical Reynolds number $Re_c$, due to non-linear processes which are still not completely elucidated (see for instance Richard \\& Zahn 1999).  Owing to their huge size and their low viscosity, stellar radiation zones have a Reynolds number which easily exceeds the critical values measured in the laboratory, $Re \\gg Re_c$, and therefore they should be turbulent, according to this simple criterion.  What seems more powerful in preventing a shearing flow from becoming unstable is the buoyancy force in a stable stratification, or the Coriolis force for fast enough rotation. We therefore expect these forces to play an important role in controlling the turbulence in stellar radiation zones; in particular, the turbulent motions will probably be anisotropic, with stronger transport in the horizontal than in the vertical direction.  Several papers have been devoted to the description of such shear turbulence, and to the derivation of prescriptions for the turbulent transport; to quote the most recent ones we refer to Zahn (1992),  Maeder (1995, 1997), Talon and Zahn (1997), Schatzman et al. (2000), Maeder (2003). Here we shall present an update of the subject, and recall the main results. We shall also introduce a new prescription for the horizontal transport, which is drawn from laboratory experiments, and discuss its impact on a stellar model, namely that of an evolving 1.5 M$_\\odot$ star. ",
        "conclusions": "We have presented here a new prescription, the  `$\\beta$-viscosity',  for the horizontal component $D_h$ of the turbulent diffusivity due the differential rotation in latitude, which is drawn from Couette-Taylor experiments. According to this prescription (eq. \\ref{new-nuh}),  $D_h$  exceeds by 3 or 4 orders of magnitude that which was initially suggested (eq. \\ref{eq:nuh-92}) and which has been used until now. Remarkably enough, the prescription established by Maeder (2003)  yields a value which is very similar, although it is based on quite different arguments. We have analyzed the effect of this dramatic change in $D_h$ on the structure of a 1.5~M$_\\odot$ star: it increases the vertical transport, and therefore enhances the mixing, particularly as the star evolves and a molecular weight gradient is building up.  The main incentive of trying to improve this prescription came from the evolutionary calculations performed by Maeder and Meynet, who found that the turbulent transport, as it was modeled initially, was insufficient to account for the observed properties of massive stars.  This new prescription may provide a better agreement with the observations, but we can hardly consider it as the final answer, given the uncertainties involved in applying to stellar interiors a result obtained in the laboratory, in cylindrical geometry and for extreme differential rotation. However we believe that this is the best we can offer today, until numerical simulations allow us to model  that type of turbulence in a more realistic way."
    },
    "0403/astro-ph0403063_arXiv.txt": {
        "abstract": "To probe the magnetic activity levels of asymptotic giant branch (AGB) stars, we used XMM-Newton to search for X-ray emission from two well-studied objects, TX Cam and T Cas. The former star displays polarized maser emission indicating magnetic field strengths of $B\\sim5$ G; the latter is one of the nearest known AGB stars. Neither star was detected by XMM-Newton. We use the upper limits on EPIC (CCD detector) count rates to constrain the X-ray luminosities of these stars, and derive $L_X < 10^{31}$ erg s$^{-1}$ ($<10^{30}$ erg s$^{-1}$) for an assumed X-ray emission temperature $T_X = 3\\times10^6$ K ($10^7$ K). These limits represent $\\lesssim 10$\\% ($\\lesssim 1$\\%) of the X-ray luminosity expected under models in which AGB magnetic fields are global and potentially play an important role in collimating and/or launching AGB winds. We suggest, instead, that the $B$ field strengths inferred from maser observations are representative of localized, magnetic clouds. ",
        "introduction": "There has been considerable recent debate over the potential role of magnetic fields in launching and/or shaping winds from asymptotic giant branch (AGB) stars (Soker \\& Kastner 2003 and references therein). Of special interest are detections of maser polarization around some AGB stars, which are indicative of the presence of relatively strong magnetic fields (e.g., Zijlstra et al.\\ 1989; Kemball \\& Diamond 1997; Miranda et al.\\ 2001; Vlemmings, Diamond, \\& van Langevelde \\ 2002; Bains 2004). SiO maser polarization measurements are particularly important in this regard, as these observations probe the important transition region between the AGB stellar photosphere and the wind. Based on high-resolution polarization maps of SiO masers, Kemball \\& Diamond (1997) deduce a magnetic field of $B \\simeq 5-10 \\G$ at a radius of $\\sim 3.5 R_\\ast$ around the AGB star TX Cam, a Mira variable. Here $R_\\ast$ is the stellar radius, which we take to be $\\sim 2 \\AU$ (see discussion in Kemball \\& Diamond 1997). With a dependence on radius of $B \\propto r^{-2}$ (Vlemmings et al. 2002) this suggests a stellar surface magnetic field of $B_\\ast \\simeq 50-100 \\G$. Miranda et al.\\ (2001) find polarization in the 1,665-MHz OH maser line that indicates the presence of $\\sim 10^{-3} \\G$ magnetic fields at $\\sim 10^{16} \\cm$ from the central star of the young planetary nebula (PN) K3-35.  Miranda et al.\\ (2001) claim that their results favor magnetic collimation models of outflows in PNs. The results for TX Cam, which will presumably undergo a PN phase, could be similarly interpreted to suggest that such magnetic collimation begins well before ionization of the circumstellar envelope. Indeed, in summarizing water maser polarization measurements for several giants, Vlemmings et al.\\ (2002) conclude that magnetic fields are strong enough to drive and shape winds from AGB stars.  However, these observations of polarized maser emission also could indicate the presence of localized, highly magnetized wind clumps, analogous to magnetic clouds in the solar wind (Soker \\& Kastner 2003), rather than large-magnitude global magnetic field strength (Soker 2002, 2003; Soker \\& Kastner 2003; Soker \\& Zoabi 2002). In that respect we note the recent results of Murakawa et al. (2003), who find that the H$_2$O maser clouds around the red supergiant VX Sgr are $\\sim 300$ times denser than the surrounding wind. On the theoretical side, based on a dynamo model for the cool supergiant Betelgeuse, Dorch (2003) finds that the magnetic structure has a typical scale of $\\sim 0.15 R_\\star$,  smaller than the the giant convection cells.  Soker \\& Zoabi (2002) summarize possible problems in models where the magnetic field plays a dynamical role in shaping the AGB wind. Among others, they consider the X-ray luminosity ($L_X$). As in the sun, they argue, globally strong magnetic fields will violently reconnect, generating flares that lead to strong X-ray emission. Such a close coupling between stellar magnetic flux and X-ray luminosity has been demonstrated to extend over 12-13 orders of magnitude in $L_X$ (Pevtsov et al.\\ 2003). Soker \\& Zoabi find that for $B_\\ast \\gtrsim 1 \\G$ on AGB stars, the expected X-ray luminosity is $\\gtrsim 10^4$ times stronger than that of the Sun, if the reconnection rate per unit surface area is similar to that in the Sun. Likewise, if the X-ray luminosity is proportional to the optical luminosity, the same scaling factor (from solar to AGB X-ray luminosity) holds.  One might even expect that AGB star X-ray surface fluxes are disproportionately larger than solar. This expectation is based on the fact that for the Sun --- where the mass loss rate is governed by magnetic activity --- the average X-ray luminosity is of the same order of magnitude as the rate of kinetic energy carried by the wind. The X-ray luminosity in the soft X-ray band (that of the ROSAT Position Sensitive Proportional Counter [PSPC], i.e., 0.1--2.4 keV) is in the range $\\sim 3 \\times 10^{26}$ to $\\sim 5 \\times 10^{27} \\erg \\s^{-1}$, at minimum and maximum, respectively (Peres et al.\\ 2000). The solar wind's kinetic energy falls between these values. If this is the case for AGB stars, as proposed in the dynamic-magnetic models, then the X-ray luminosity should be a factor of $\\sim 10^6-10^8$ stronger than in the Sun (Soker \\& Zoabi 2002). Soker \\& Zoabi conclude that this expectation, of $L_x \\sim 10^{31} - 10^{35} \\erg \\s^{-1}$, is in sharp contradiction with most ROSAT observations, which demonstrate that the X-ray luminosities of red giant stars typically are only marginally larger than the solar X-ray luminosity (Schr\\\"oder {\\it et al.} 1998) and that $L_x$ likely further decreases in late giant evolution (H\\\"unsch \\& Schr\\\"oder 1996). On the other hand, some red giants are known to be relatively luminous in X-rays (i.e., $L_x \\sim 10^{30} \\erg \\s^{-1}$; e.g., H\\\"unsch \\& Schr\\\"oder 1996; H\\\"unsch 2001).  In Soker \\& Kastner (2003) we discuss flares on AGB stars from locally, rather than globally, strong fields. Such fields should result in much weaker X-ray emission (Soker \\& Zoabi 2002; Pevtsov et al.\\ 2003). It is possible that such weak coronal X-ray emission from heavily obscured, mass-losing AGB stars has escaped detection thus far, due to ROSAT's lack of hard X-ray sensitivity and the large distances to these short-lived, luminous stars. Indeed, Mira itself is a weak ROSAT X-ray source ($L_x \\sim 2\\times10^{29} \\erg \\s^{-1}$; Karovska et al.\\ 1996; Soker \\& Kastner 2003), although the origin of its X-ray emission is uncertain given the presence of a close companion.  In an attempt to further constain models for magnetic fields in AGB stars, we are conducting observations of selected AGB stars with the XMM-Newton X-ray observatory, which features sensitivity and energy coverage far superior to that of ROSAT. In this paper we report on and discuss results from XMM-Newton observations of the polarized maser source TX Cam and the nearby AGB star T Cas. Like TX Cam, T Cas is a relatively strong SiO maser source (Herpin et al.\\ 1998), although we are unaware of measurements of the polarization of its maser emission or of high-resolution imaging of its maser spots. ",
        "conclusions": "Given the above considerations, we tentatively place upper limits of $L_X < 10^{31}$ erg s$^{-1}$ on the X-ray luminosities of TX Cam and T Cas (Fig.\\ 1). If the X-ray emission due to magnetic activity is relatively hard, then the upper limits become far more stringent: e.g., $L_X < 10^{30}$ erg s$^{-1}$ for $T_X = 10^7$ K and $N_H < 10^{22}$ cm$^{-2}$.  In light of the large uncertainties, these results do not rule out the claim that magnetic fields shape the winds from AGB stars. However, it seems that the nondetection of TX Cam, in particular, strongly constrains such models. Given the evidence for relatively strong magnetic fields around this star, and the tight coupling between stellar magnetic flux and $L_X$ (e.g., Pevtsov et al.\\ 2003), the upper limit on TX Cam's X-ray luminosity indicates that the magnetic field behavior is much quieter and/or that its active regions are far less widespread than on the Sun. Yet AGB stars, like TX Cam, have very strong convective envelopes, strongly pulsate, and lose mass at high rates. All these characteristics, we expect, should lead to a much more chaotic magnetic field structure than on the Sun, with more violent reconnection and liberation of magnetic energy (Soker \\& Kastner 2003).  Our conclusion --- although not firm, due largely to lingering uncertainties in $T_X$ and $N_H$ --- is that TX Cam and T Cas are not strong X-ray sources, and hence their average magnetic fields are much weaker than would be required to collimate and/or drive their AGB winds. Instead, the magnetic field must be concentrated in small regions, some of which give rise to the polarized maser emission from TX Cam (Soker \\& Kastner 2003). Note that some magnetic activity, but with much weaker X-ray luminosity than that expected for a global magnetic field, is still expected in our local field model (Soker \\& Kastner 2003).  Clearly more work, on both theoretical and observational aspects of AGB star magnetic fields, is needed. High-resolution polarization mapping of SiO emission from T Cas and other relatively nearby, bright maser sources is called for, as are additional, sensitive X-ray observations of AGB stars with moderate $B$ fields, as inferred from maser polarization and other methods.  We also suggest more attention be paid to post-main sequence magnetic activity models in which the fields are local, and have no global dynamical role."
    },
    "0403/astro-ph0403549_arXiv.txt": {
        "abstract": "From both photometric and broadband spectral monitoring of $\\gamma$-ray burst (GRB)\\index{GRB}\\index{Gamma-ray Burst} lightcurve ``bumps,'' particularly in GRB\\,011121, a strong case grew for a supernova (SN) origin. The GRB-SN connection was finally solidified beyond a reasonable doubt with the discovery that the bump in GRB\\,030329\\index{030329} was spectroscopically similar to a bright Type Ic SN. In light of this result, I redress the previous SN bump claims and conclude that 1) the distribution of GRB-SN bump peak magnitudes is consistent with the local Type Ibc SNe peak distribution and suggest that 2) the late-time bumps in all long-duration GRBs are likely supernovae\\index{supernova}. ",
        "introduction": "The connection of GRBs to supernovae (SNe) -- and hence the death of massive stars -- was explored theoretically \\cite{col68} even before the discovery of GRBs. In the subsequent discovery paper, Klebesadel, Strong, \\& Olson reported on a catalog search for SNe coincident in time and position with the first 13 GRBs known \\cite{kso73}.  In retrospect, the lack of a clear association is easy to explain: not only were SNe sparsely discovered, but the highest redshift SN discovered by 1973 (SN\\,1968P; $z = 0.11$) was less than the redshift of all $\\sim$35 classic GRBs redshifts known today. With the development of a model for the production of a GRB from a collapsing massive star \\cite{woo93}, came renewed interest in the possibility of a GRB--supernova connection. The modern form of and new spins on Woosley's collapsar was presented in great detail by Andrew MacFadyen (this meeting).  The discovery of the low redshift SN\\,1988bw, a bright type Ic associated with GRB\\,980425 \\cite{gvv+98} implied that massive stars were capable of producing at least some incarnation of a GRB: the low energy release of GRB\\,980425\\index{GRB 980425} compared with all other known cosmological GRB energies suggests it may be only a sub-class of GRBs \\cite{bkf+98}.  Attempts to theoretically unify GRB\\,980425 with the majority of the long-duration cosmological ($z > 0.1$) bursts fall short of explaining all the observational differences (see papers by Iwamoto, Patat this meeting).  Our discovery of a faint red light curve ``bump'' superimposed on the afterglow of GRB\\,980326 \\cite{bkd+99} was the first suggestion of a SN associated with a cosmological GRB. Supernovae-like features were later found photometrically in a number of bursts, some more plausibly than others; the principal difficulty is that SNe peak at a brightnesses comparable to their respective host galaxies (see Fig.~\\ref{fig:1}; Refs.~\\cite{blo02,blo03}). Our multicycle effort on HST (Kulkarni, PI) was designed to search for late-time afterglow features and discriminate between other physical models for such emission. While a few bursts showed no evidence for a SN (e.g., ref.~\\cite{pks+03}), the late-time multicolor HST light curve and broadband spectra of GRB\\,011121 was reliably modeled by a bright Type Ic supernova that occurred nearly contemporaneously with the GRB \\cite{bkp+02}. Moreover, this was the first burst for which the data were of high enough quality to statistically demonstrate that the afterglow propagated in a wind-stratified medium rather than a constant-density ISM \\cite{pbr+02}. I recently reviewed the pre-2003 evidence for the GRB-SN connection \\cite{blo03}.  \\label{sec:intro} \\begin{figure}[bht] \\centering \\includegraphics[height=3.6in]{Bloom_fig1.ps} \\caption{The difficulty of detecting supernovae associated with GRBs: the magnitude of the host galaxy versus the peak magnitude of the putative bump (in whichever optical filter the peak bump magnitude was reported). Above the dashed line, bump magnitudes are fainter than the integrated light of the host galaxy. The circles are measurements where only ground-based imaging was used and squares are measurements derived also using {\\it Hubble Space Telescope} (HST) \\index{HST}\\index{Hubble Space Telescope} imaging. Relative significances of the reliability of the bump detection were assigned by the author based roughly on the number of detections of the bump above the host light and the photometric error reported on the bump peak. Updated (from Ref.\\cite{blo03}) with new bumps in 030329 and 021211, reported in this meeting.} \\label{fig:1}       % \\end{figure} ",
        "conclusions": ""
    },
    "0403/astro-ph0403512_arXiv.txt": {
        "abstract": "We have identified a new Lyman limit absorption system towards PKS1937-1009, with log$N$(\\hi) $=18.25 \\pm 0.02$ at $z=3.256$ that is suitable for measuring D/H.  We find a $68.3$\\% confidence range for D/H of $1.6^{+0.25}_{-0.30} \\times 10^{-5}$, and a $95.4$\\% range of $1.6^{+0.5}_{-0.4} \\times 10^{-5}$. The metallicity of the cloud where D/H was measured is low, [Si/H] $= -2.0 \\pm 0.5$. At these metallicities we expect that D/H will be close to the primordial value. Our D/H is lower than the D/H value predicted using the $\\Omega_b$ calculated from the cosmic background radiation measured by WMAP, $2.60^{+0.19}_{-0.17} \\times 10^{-5}$.  Our result also exacerbates the scatter in D/H values around the mean primordial D/H. ",
        "introduction": "There are interesting disagreements between the various ways of calculating the baryon density of the universe ($\\Omega_b$).  The most precise $\\Omega_b$ measurement is from the recent WMAP observations of the CMB power spectrum \\citep{Spergel03}.  Another way to measure $\\Omega_b$ is by measuring the relative abundances of the light elements produced in big bang nucleosyntheis (BBN), since in standard BBN theory $\\Omega_b$ is the only free parameter determining these abundances.  Attempts have been made to measure the primordial $^4$He, $^7$Li and deuterium (D) abundances.  In principle, $\\Omega_b$ is most reliably determined by D, due to its apparently uncomplicated evolution with time \\citep*{Epstein74} and because it depends more sensitively on $\\Omega_b$ than $^4$He and $^7$Li. There is broad agreement between the D $\\Omega_b$ and CMB $\\Omega_b$, but the $^4$He and $^7$Li measurements predict a lower $\\Omega_b$ \\citep{Coc04, Cyburt03}.  Either there is something wrong with our interpretation of the $^4$He and $^7$Li measurements, or the standard BBN theory may need to be modified \\citep[e.g. see][]{Barger03, Dmitriev04, Ichikawa04}.  The mean of the primordial D values as measured in quasar absorption systems is consistent with the WMAP CMB $\\Omega_b$.  However, there is statistically significant scatter of the individual D/H measurements about the mean.  The metallicity of the absorption clouds used to make the D/H measurements is low, [M/H]$< -1.5$, so that no astration (D destruction due to stellar nucleosynthesis) is thought to have occurred \\citep{Fields96}.  In addition, the primordial D/H is thought to be isotropic and homogeneous, so the scatter is hard to explain. One explanation is that the systematic errors in measuring D/H have been underestimated by the authors.  However, with only half a dozen quasar absorption system D/H measurements other explanations, such as some early mechanism for astration or a non standard BBN, cannot be ruled out.  More D/H measurements are needed to resolve this problem. It is exciting to see further results being added to this short list: in the last year two D/H measurements have been published: one towards another QSO absorption line system \\citep{Kirkman03} and another in a low metallicity gas cloud near our Galaxy \\citep{Sembach03}.  Here we present an analysis of D/H in a newly discovered $z=3.256$ absorption system towards PKS1937-1009 (see Fig.~\\ref{img}, Fig.~\\ref{ll}). Coincidentally, \\citet*{Burles98} measured D/H in a different absorber at $z=3.572$ towards this same QSO.  The new absorber is not connected with the $z=3.572$ absorber, and we expect it to give an independent measurement of D/H. ",
        "conclusions": "We have identified a new Lyman limit absorption system towards PKS1937-1009, with log$N$(\\hi) $=18.25 \\pm 0.02$ at $z=3.256$.  The \\lya\\ and \\lyb\\ transitions are suitable for measuring D/H, and we find a $68.3$\\% range for D/H of $1.6^{+0.25}_{-0.30}\\times 10^{-5}$, and a $95.4$\\% range of $1.6^{+0.5}_{-0.4} \\times 10^{-5}$.  The metallicity of the cloud where D/H was measured is low, [Si/H] $= -2.0 \\pm 0.5$.  At this metallicity we expect that the D/H value will be primordial.  This value disagrees at the $99.4$\\% level with the predicted D/H value using the $\\Omega_b$ calculated from the WMAP cosmic background radiation measurements, $2.60^{+0.19}_{-0.17} \\times 10^{-5}$ \\citep{Coc04}, and it exacerbates the scatter in D/H measurements around the mean D/H value. It is consistent with the D/H values measured in the ISM using FUSE, \\citep[e.g.][]{Hebrard03}.  However, it is not consistent with the idea of a single primordial D/H value which stays constant until the metallicity reaches $\\sim 1/10$ solar, then begins to drop off, reaching the ISM D/H at solar metallicity. Some early mechanism for D astration may be the cause for the scatter in D/H values \\citep{Fields01}, although in this case we would expect to see a correlation between D/H and [Si/H], which is not observed.  Other authors \\citep{ Pettini01a, Kirkman03} note that there is a trend of decreasing D/H with increasing $N$(\\hi) of the absorbers in which they are measured.  The D/H value in this paper does not follow this trend, suggesting that the trend is not real and only due to the low number of D/H measurements. \\\\ \\\\ We would like to thank Antoinette Songaila for making the Keck data used in this analysis available and Max Pettini for his helpful correspondence. We also thank the anonymous referee for their constructive comments that helped improve the paper.  AOG and AFS acknowledge support from the Generalitat Valenciana ACyT project GRUPOS03/170, and Spanish MCYT projects AYA2002-03326 and AYA2003-08739-C02-01."
    },
    "0403/astro-ph0403038_arXiv.txt": {
        "abstract": "s{Neutrino processes in semi-transparent supernova matter, opaque to semi-transparent protoneutron star matter, and catalyzed neutron stars  are discussed.} ",
        "introduction": "In recent years, the study of neutrino emission, scattering, and absorption in matter at high density and/or temperature has gained prominence largely due to its importance in a wide range of astrophysical phenomena.  Energy loss in degenerate helium cores of red giant stars,\\cite{Gross,Raffelt(2000)} cooling in pre-white dwarf interiors,\\cite{O'BrienKawaler(2000)} the short- and long-term cooling of neutron stars,\\cite{Prak02,Yak02} the deflagration stages of white dwarfs which may lead to type Ia supernovae,\\cite{Wolfgang,Iwamoto} explosive stages of type II (core-collapse) supernovae,\\cite{Bur00} and thermal emission in accretion disks of gamma-ray bursters,\\cite{Matteo,Kohri} are examples in which neutral and charged current weak interaction processes that involve neutrinos play a significant role.  Since this is a vast subject, we will highlight some recent developments in the context of core-collapse supernovae and neutron stars from their birth to old age. ",
        "conclusions": "\\begin{figure}[htb] \\centerline{\\epsfxsize=3.7in\\epsfbox{korea_thank-you_4.eps}} \\centerline{\\epsfxsize=3.7in\\epsfbox{flag.eps}} % \\label{fig:thanku} \\end{figure}"
    },
    "0403/astro-ph0403454_arXiv.txt": {
        "abstract": "Ultracompact HII regions are signposts of massive star formation and their properties provide diagnostics for the characteristics of very young O stars embedded in molecular clouds.  While radio observations have given us a good picture of the morphology of these regions, they have not provided clear information about the kinematics.  Using high spectral resolution observations of the $12.8\\mu$m [NeII] line, it has been possible for the first time to trace the internal kinematics of several ultracompact HII regions.  We find that the motions in the cometary ultracompact HII regions MonR2 and G29.96-0.02 are highly organized.  The velocity patterns are consistent with parabolic ionized flows along a neutral boundary layer. ",
        "introduction": "An understanding of the formation of massive stars and of the effects of these stars on their environments is critical to an understanding of star formation overall.  Most stars of all masses form in clusters containing massive stars \\cite{ladalada}. High mass stars and the regions of the ISM that they excite are what we use to trace star formation in other galaxies. The formation of massive stars ultimately destroys the star forming clouds.  Ultracompact HII (UCHII) regions are one of the earliest manifestations of massive stars. High extinction and substantial emission by overlying dust make it impossible to detect the photospheres of very young O stars directly at any wavelength. The UCHII regions, however, become detectable in the radio continuum and in mid-IR fine-structure lines while young O stars are still embedded in their natal envelopes, as soon as the stars have begun to emit any significant amount of Lyman continuum radiation.  We have been using the $12.8\\mu$m fine-structure line of [NeII] to study the morphology and kinematics of a sample of UCHII regions.  Our ultimate goal is to understand the natures of the regions and how they depend on the age and mass of the exciting star, as well as on the properties of any disk or envelope.  We would also like to learn how long the UCHII regions last and how the evolution of these regions affects the surrounding cloud. In this paper, we summarize results from our initial studies of UCHII regions with an apparently cometary morphology \\cite{jaffe} \\cite{zhu}. ",
        "conclusions": ""
    },
    "0403/astro-ph0403381_arXiv.txt": {
        "abstract": "We estimate the contribution of AGNs and of their host galaxies to the infrared background. We use the luminosity function and evolution of AGNs recently determined by the hard X-ray surveys, and new Spectral Energy Distributions connecting the X-ray and the infrared emission, divided in intervals of absorption. These two ingredients allow us to determine the contribution of AGNs to the infrared background by using mostly observed quantities, with only minor assumptions. We obtain that AGN emission contributes little to the infrared background ($<$5\\% over most of the infrared bands), implying that the latter is dominated by star formation. However, AGN host galaxies may contribute significantly to the infrared background, and more specifically 10--20\\% in the 1--20$\\mu$m range and $\\sim$5\\% at $\\lambda<60\\mu m$. We also give the contribution of AGNs and of their host galaxies to the source number counts in various infrared bands, focusing on those which will be observed with Spitzer. We also report a significant discrepancy between the expected contribution of AGN hosts to the submm background and bright submm number counts with the observational constraints. We discuss the causes and implications of this discrepancy and the possible effects on the Spitzer far-IR bands. ",
        "introduction": "Active Galactic Nuclei (AGN) are known to be the main contributors to the cosmic X-ray background. Synthesis models suggested that the shape of the X-ray background requires a mixture of unabsorbed and absorbed AGNs (Setti \\& Woltjer 1989; Comastri et al. 1995; Gilli, Salvati, \\& Hasinger 2001). This scenario has been verified by the recent Chandra and XMM surveys which have resolved most of the X-ray background up to energies of 10~keV (Brandt et al. 2001; Giacconi et al. 2001; Hasinger et al. 2001). Indeed, the optical identification of the sources making the X-ray background, as well as their X-ray spectra, revealed both obscured and unobscured AGNs, though unusual AGNs were also identified (e.g. Mainieri et al. 2002; Barger et al. 2003; Fiore et al. 2003). The X-ray surveys have also provided evidence that the evolution of AGNs is different with respect to that known from previous optical surveys. In particular, the evolution of low-luminosity AGN (in the Seyfert range) peaks at much lower redshift (z$<$0.5) than for more luminous, QSO-like AGNs (Hasinger 2003; Fiore et al. 2003; Ueda et al. 2003). This luminosity dependent density evolution (LDDE) seems to apply both to unobscured and obscured AGNs.  The circumnuclear gas responsible for the X-ray absorption is associated with dust which absorbs the optical-UV radiation from the AGN and reprocesses it thermally into the infrared. Dust infrared emission has been observed in all AGNs (except in some BL Lacs and radio-loud quasars where the IR radiation may be dominated by the relativistically boosted synchrotron radiation). At least in the near and mid-IR, it has been recognized that a significant fraction of the dust radiation is due to heating by the AGN (e.g. Granato, Danese, \\& Franceschini 1997; Maiolino et al. 1998a; Oliva et al. 1999; Glass 2004; Minezaki et al. 2004). For this reason it has been argued that AGNs may contribute significantly also to the IR cosmic background. This issue is of utmost importance, since it would also have implications on the fraction of IR background which should be ascribed to star formation.  Some previous studies have estimated the contribution of AGN to the IR background or to the source number counts (Granato et al. 1997; Matute et al. 2002; Risaliti, Elvis, \\& Gilli 2002; Andreani, Spinoglio, \\& Malkan 2003).   The main ingredients for this estimate are the luminosity function (LF) and evolution of AGNs and their spectral energy distribution (SED). Previous works have adopted the observationally available, still largely incomplete, LF of AGNs, that only very recently has been better defined through Chandra and XMM surveys (Ueda et al. 2003). Generally, a single template for the AGN SED has been assumed.  In this paper we investigate the contribution of AGNs to the IR background and counts by using the most recent luminosity functions and evolution obtained by the recent hard X-ray surveys. We also use IR spectral energy distributions accurately derived by means of nuclear IR data for a sample of Seyfert galaxies (both obscured and unobscured) and which allow us to disentangle the AGN IR emission from the contribution of the host galaxy. These data also allow us to directly connect the IR emission of AGNs to their intrinsic X-ray emission. Moreover, we use the novel approach of dividing the observed IR SEDs into intervals of different AGN obscuration, in terms of N$_H$ along the line of sight. This is the same method adopted in the synthesis of the X-ray background. These various pieces of information combined together allow us to {\\it directly connect the X-ray background (dominated by AGNs) to the IR background without strong assumptions, by using only observed and measured quantities.}  Throughout the paper we adopt the following cosmological parameters: $\\rm H_0 =70~km~s^{-1}~Mpc^{-1}$, $\\rm \\Omega _m = 0.3$ and $\\rm \\Omega _{\\Lambda} = 0.7$. ",
        "conclusions": "We have estimated the contribution to the infrared background and infrared number counts from AGNs by using the most recent AGNs luminosity functions and evolution from the hard X-ray surveys (Ueda et al. 2003) and by connecting their X-ray and IR emission by means of new detailed AGN SEDs derived by us. This approach allow us to derive the contribution of AGN to the infrared background by using mostly observed and measured quantities, with little additional assumptions.  The main results are the following: \\begin{enumerate} \\item The AGN nuclear IR radiation contributes little ($<$5\\%) to the IR background at most wavelengths. \\item AGNs along with their host galaxies contribute significantly ($\\sim$~5--20\\%) to the IR background at $\\lambda < 60 \\mu m$. \\item AGNs and their hosts probably dominate the mid-IR number counts at bright fluxes ($>$1~mJy). \\item The predicted contribution of AGNs and of their host galaxies to the mid-IR background and number counts are in good agreement with the observational constraints available so far. \\item However, we find a strong discrepancy between the expected contribution of Seyfert host galaxies at high redshift to the submm background and number counts at bright fluxes. We ascribe this discrepancy mostly to the shortage of Seyferts at high redshift (z$\\sim$2) in the X-ray luminosity functions which, at these high redshifts and low luminosities, are still poorly constrained by the current hard X-ray surveys. This issue may also yield to underestimate the AGN hosts contribution to the far-IR bands of Spitzer. A more detailed investigation of this issue and of the possible implications for the Seyferts evolution at high redshift will be discussed in an upcoming paper (Silva et al. in prep.). \\end{enumerate}"
    },
    "0403/astro-ph0403662_arXiv.txt": {
        "abstract": "The Far Ultraviolet Spectroscopic Explorer (\\fuse) has surveyed a large sample ($ > 100$) of active galactic nuclei in the low redshift universe ($z < 1$).  Its response at short wavelengths makes it possible to measure directly the EUV spectral shape of QSOs and Seyfert 1 galaxies at $z < 0.3$. Using archival \\fuse\\ spectra, we form a composite extreme ultraviolet (EUV) spectrum of QSOs at $z < 1$ and compare it to UV/optical composite spectra of QSOs at higher redshift, particularly the composite spectrum from archival Hubble Space Telescope spectra. ",
        "introduction": "The ubiquity with which QSOs display spectral properties such as power-law continua and broad emission lines over wide ranges in luminosity and redshift has led to the use of composite spectra to study their global properties. Information about the continuum in the rest-frame ultraviolet is particularly critical for understanding the formation of the emission lines, for characterizing the Big Blue Bump, and for determining the ionization state of the intergalactic medium (IGM). Composite QSO spectra covering the rest-frame ultraviolet have been constructed for objects  with $0.33 < z < 3.6$ from \\hst\\ (Zheng et al.\\ 1997, Telfer et al.\\ 2002, T02 hereafter), and at $z > 2$ from ground-based samples like the SDSS (Vanden Berk et al.\\ 2001), the First Bright Quasar Survey (Brotherton et al.\\ 2001) and the Large Bright Quasar Survey (Francis et al.\\ 1991).  The \\fuse\\ bandpass, 905-1187 \\AA, allows us to examine the EUV properties of local AGN. We can therefore study the same rest-frame wavelength region covered by the \\hst\\ composite spectra, at redshifts less than $0.33$. The low redshifts of these AGN ensure that, although the \\fuse\\ aperture limits it to observing relatively bright AGN, our sample contains a large fraction of intrinsically low-luminosity objects. An added advantage to working with low redshift spectra is that the determination of the mean EUV spectral index requires a less significant correction for IGM absorption than was required for  the \\hst\\ sample. ",
        "conclusions": "We summarize our results as follows:  (1) We construct a composite EUV (630-1155 \\AA) AGN spectrum of objects with $z < 0.67$ from archival \\fuse\\ data. (2) We fit a power law continuum and Gaussian profiles to the emission lines in the composite spectrum, and we find that O~{\\sc vi}/Ly$\\beta$ and Ne~{\\sc viii} emission  are enhanced in the \\fuse\\ composite relative to the \\hst\\ composite due to the Baldwin effect, seen also in the comparison between the \\hst\\ and SDSS composites (T02). (3) We find that the best-fit spectral index of the composite is $\\alpha=-0.56^{+0.38}_{-0.28}$. The conservative estimate of the total error in the spectral index includes the standard deviation in $\\alpha$ in 1000 bootstrap samples of the \\fuse\\ data set, uncertainties in the extinction correction applied to the \\fuse\\ spectra, and uncertainties in the column density distribution parameter of the intervening Ly$\\alpha$ forest. (4) The \\fuse\\ composite is harder than the EUV portion of the \\hst\\ composite spectrum of T02 who find $\\alpha=-1.76 \\pm 0.12$ for 332 spectra of 184 AGN with $z>0.33$. The Baldwin effect is generally attributed to the tendency for low-luminosity AGN tend to show harder ionizing continua (Zheng \\& Malkan 1993; Wang et al.\\ 1998; Dietrich et al.\\ 2002). The median luminosity of the AGN in the \\fuse\\ sample is $\\log L_{\\rm median}=41.2$, versus $\\log L_{\\rm median}=42.9$ for the \\hst\\ sample. One interpretation of these results is that both the enhanced high-ionization emission line strengths and the harder continuum shape of the \\fuse\\ composite spectrum are due to the larger fraction of low-luminosity AGN in the \\fuse\\ sample.  However, we note that splitting the \\fuse\\ sample itself into high- and low-luminosity/redshift subsamples as shown in Figure 2 results in only marginally different values for the EUV spectral index."
    },
    "0403/astro-ph0403348_arXiv.txt": {
        "abstract": "We report on  the detection of a molecular band centered at $\\sim$98~$\\mu$m ($\\sim$102~cm$^{-1}$), observed with the $\\textit {Infrared Space Observatory}$$^1$ in the young Planetary Nebula NGC~7027. The band structure and intensity can not be reproduced by atomic fine structure lines, recombination lines or by the rotational emission of abundant molecules. We discuss the possible contribution of the low--energy bending modes of pure carbon chains to the unidentified far--IR bands observed in  $C$--rich evolved objects. In particular, we speculate that the band emission could arise from the $\\nu_9$ and $\\nu_7$ bending modes of C$_6$ and C$_5$, for which wavenumbers of 90$\\pm$50 and  107$\\pm$5~cm$^{-1}$ have been estimated from photoelectron spectroscopy. ",
        "introduction": "Carbon chains molecules are of special interest for Astrophysics since Douglas (1977) proposed them as possible carriers of the diffuse interstellar bands (DIBs). In the past years, many polar carbon chains have been detected in the interstellar and circumstellar media (ISM and CSM) through their pure rotational spectrum at radio wavelengths. Major examples are the cyanopolyynes, HC$_{2n+1}$N with $n$=1--5 (Turner 1971, Avery et al. 1976, Kroto et al. 1978, Broten et al. 1978, Bell et al. 1997), and the hydrogenated carbon chain radicals such as C$_5$H, C$_6$H, C$_7$H  and C$_8$H (Cernicharo et al. 1986, Gu\\'elin et al. 1987, 1997, Cernicharo \\& Gu\\'elin 1996).   On the other hand, larger carbon molecular complexes such as the polyatomic aromatic hydrocarbons (PAHs) are believed to dominate the ubiquitous mid--IR emission seen in the  unidentified infrared bands (UIBs). However, only the single aromatic species, benzene, has been so far identified through its IR active modes (Cernicharo et al. 2001).  All this chemistry richness suggests that the growth mechanisms producing complex carbon molecules such as the PAHs or the fullerenes are highly efficient. Nevertheless, the set of `building blocks'' and the possible chemical reactions still have to be identified and clarified. Among the possible species, the pure carbon chains, (C$_n$; Van Orden \\& Saykally 1998), could be the `skeletons'' from which larger organic molecules can be formed (see Cernicharo, Goicoechea \\& Caux 2000, Cernicharo 2004). Due to the lack of permanent electric dipole, these species do not have rotational spectrum to be observed from radio telescopes. The only way to detect them in the dense ISM and CSM is through their asymmetrical stretching modes around $\\sim$5~$\\mu$m (2000~cm$^{-1}$) and/or through their low--energy bending modes around $\\sim$100~$\\mu$m ($\\sim$100~cm$^{-1}$). The first technique  allowed the detection of C$_3$ (Hinkle et al. 1988) and C$_5$ (Bernath et al. 1989) in the circumstellar envelope of the bright IR evolved star IRC+10216. However, few sources (specially in the ISM) have enough flux at $\\sim$5~$\\mu$m to allow systematic studies of the C$_n$ stretching modes, and the less known far--IR bending modes are the only way to detect these species.  Before the launch of ISO (Kessler et al. 1996), we proposed to observe the bending modes of several polyatomic molecules in the far--IR. As a result, the $\\nu_2$ bending mode of C$_3$ has been observed in Sgr~B2 and IRC+10216 (Cernicharo et al.  2000). In addition, we have detected several Unidentified far-IR bands (UfIBs) in the ISO spectrum of many C--rich sources that may also be related with the bending modes of carbon chains species. In particular we have tentatively assigned a  UfIB at $\\sim$57.5~$\\mu$m ($\\sim$174~cm$^{-1}$), observed  in Sgr~B2 and in $C$--rich evolved stars including the planetary nebula (PNe) NGC~7027, to the $\\nu_5$ bending mode of C$_4$ (Cernicharo, Goicoechea \\& Benilan 2002).  NGC~7027 is one of the most studied evolved objects in the Galaxy. It is a young PNe, $\\sim$1000~years since it left the AGB stage, characterized by a relatively compact ionized region (Volk \\& Kwok 1997) driven by the radiation that arises from a central source at $\\sim$200,000~K (Latter et al. 2000). In addition, a rich $C$--chemistry is taking place in the larger molecular envelope  ($\\sim$40$''$ in CO; Masson et al. 1985) that surrounds the inner photon dominated regions. In fact, NGC~7027 was the first object to show UIB emission latter attributed to the PAH emission (Gillett et al. 1973) and the only source where pure rotational lines of CH$^+$ have been detected (Cernicharo et al. 1997).  Therefore, the detection of unidentified features in the IR spectrum of NGC~7027 has historically contributed to a better understanding of the interstellar carbon complexity. In this work, we report the detection of a new UfIB at $\\sim$98~$\\mu$m ($\\sim$102~cm$^{-1}$) in the far--IR spectrum of NGC~7027 and speculate about it possible origin. ",
        "conclusions": "Both PAHs and pure carbon chains possesses low--energy bending modes in the far--IR. Contrary to the UIB emission in the mid--IR, the far--IR skeletal modes of PAHs (vibrations associated with the bending of the skeletal structure) depends of the exact nature of the species. Hence, the rovibrational structure of the PAH skeletal mode has to be calculated for each specific PAH. As a example, the lowest energy active mode of coronene (C$_{24}$H$_{12}$) lies at $\\sim$127~cm$^{-1}$ (Joblin et al. 2002), while for ovalene (C$_{32}$H$_{14}$), two far--IR skeletal modes at $\\sim$120~cm$^{-1}$ and $\\sim$66~cm$^{-1}$ have been theoretically investigated (Mulas et al. 2003). The far-IR emission of a few PAHs has also been observed in gas phase experiments (Zhang et al. 1996). However, if one of the observed UfIBs belongs to a specific PAHs, other bands  (generally stronger) should be observed in the IR spectrum. This is not the case of the new UfIB presented in this work and we have considered the low--energy modes of pure carbon chains as possible carriers of the observed UfIB.  The detection of several rovibrational lines of the $\\nu_2$ bending mode of C$_3$ ($\\sim$63~cm$^{-1}$) in Sgr~B2 and IRC+10216 evidenced  high abundances for this species and opened the possibility to detect more  carbon chains in the ISM and CSM (Cernicharo et al. 2000). At the low resolution of the LWS/grating, the more intense C$_3$ lines [Q(2,4,6)] are blended with the strong [\\CII]158~$\\mu$m line and no assignation could be made in NGC~7027. However, the tentative detection of the $\\nu_5$ bending mode of C$_4$ in many objects including NGC~7027 ($\\sim$174~cm$^{-1}$; Cernicharo et al. 2002), if confirmed, will imply a C$_3$/C$_4$$<$10 abundance ratio, which will suggest that the abundance of C$_n$ does not decrease  drastically when $n$ increases.  In spite of their importance, the C$_n$  bending modes are difficult to characterize spectroscopically in the laboratory or by \\textit{ab initio} computations. As expected, the main uncertainties difficulting their astronomical detection are the band origin and the IR intensity. Moreover, the induced dipole moments and other spectroscopic constants are often  known for the fundamental transition but less or nothing is known about overtone transitions within the bending mode or transitions to other excited states including the stretching modes.  The ground electronic state is  different for the odd-- and even-- numbered cumulenic chains. This property determines the observed band shape of the bending mode. Odd--numbered linear C$_n$'s have a singlet $^{1}\\Sigma^{+}_{g}$ ground electronic state which results in a perpendicular $^{1}\\Pi_{vib}-{^1}\\Sigma_{vib}$ vibronic spectrum with a strong --$Q$ branch and weak --$P$ and --$R$ branches. On the other hand, even--numbered linear C$_n$'s have a triplet $^{3}\\Sigma^{-}_{g}$ ground electronic state. Hence, each rotational line  is split in three components due to  spin couplings. This splitting increase with the number of atoms but is  particularly small for C$_4$ and C$_6$ (Giesen et al. 2001). The vibronic transition is now $^{3}\\Pi_{vib}-{^3}\\Sigma_{vib}$, and despite the null component of the electronic orbital momentum in the ground state, the spin--orbit constant (A$_{SO}$) could be large in  even-- C$_n$'s. In such a case, the resulting band--shape  will be notoriously affected by the value of A$_{SO}$.  \\subsection{Tentative detection of C$_6$ and C$_5$}  With the exception of the $\\nu_2$ mode of C$_3$, the band origins of other C$_n$ bending modes are not accurately constrained. Hence, their assignation as carriers of the UfIBs is not obvious. Although the low--lying bending modes of  C$_5$ and C$_6$ have not been directly observed in gas phase, different studies suggest wavenumbers around $\\sim$100~cm$^{-1}$ (see below). The observed new UfIB around $\\sim$98~$\\mu$m is composed by several peaks at $\\sim$91.8, 92.8,  93.9, 95.2, 97.6, 98.9, 100.5 and 101.7~$\\mu$m. In the following, we will consider the possibility that \\textit{any of these peaks is  related with the $\\nu_9$ and $\\nu_7$ bending modes of C$_6$ and C$_5$ respectively}.   C$_5$ was first detected in the gas phase by Heath et al. (1989) who observed the $\\nu_3$ stretching mode at $\\sim$2169~cm$^{-1}$ ($\\sim$4.6~$\\mu$m). The same mode was observed in IRC+10216 by Bernath et al. (1989). Actually,  C$_5$ is the largest C$_n$ detected in the CSM. On the other hand, C$_6$ was first identified  in the laboratory by its electronic spin resonance spectrum in an Ar matrix (Van Zee et al. 1987) and later observed in gas phase through its $\\nu_4$ stretching mode at $\\sim$1960~cm$^{-1}$ ($\\sim$5.1~$\\mu$m) by Hwang et al. (1993). Moazzen--Ahmadi et al. (1989) observed the ($\\nu_3$+$\\nu_7$)--$\\nu_7$ and ($\\nu_3$+2$\\nu_7$)--2$\\nu_7$ hot bands arising from the $\\nu_7$ bending mode of C$_5$. From the $l$--doubling constant $q_7$, they estimated a frequency of $\\nu_7$=118$\\pm$3~cm$^{-1}$. More recently, C$_5$ has been observed in photoelectron spectra. In particular, Arnold et al. (1991) measured the 2$\\nu_7$ transition and estimated $\\nu_7$=101$\\pm$45~cm$^{-1}$, while Kitsopoulos et al. (1991) obtained $\\nu_7$=107$\\pm$5~cm$^{-1}$. \\textit{Ab initio} calculations predict an infrared intensity around  36~km~mol$^{-1}$ [$A$($\\nu_7$=1-0)$\\simeq$0.06~s$^{-1}$; Hutter \\& L\\\"uthi  1994; Martin et al. 1995]. Fig.~1 shows the expected band--shape for a $^{1}\\Pi_{vib}-{^1}\\Sigma_{vib}$ transition with this intensity and  the molecular constants for the $\\nu_7$ mode (Moazzen--Ahmadi et al.) centered at 104.8~cm$^{-1}$ (95.4~$\\mu$m). As said before, any of the band peaks observed in NGC~7027 around $\\sim$98~$\\mu$m ($\\sim$102~cm$^{-1}$) could well be responsible of the $\\nu_7$ mode. However, only the best fit to the band including the C$_6$ model is shown in  Fig.~1. If the carrier of the 95.5~$\\mu$m feature is finally C$_5$ at $\\nu_7$=104.8~cm$^{-1}$, the estimated column density assuming an excitation temperature of 100~K (typical of a PNe envelope) is $N$(C$_5$)=1.8$\\times$10$^{14}$~cm$^{-2}$. Assuming 1.5--4.0  magnitudes of visual extinction in the neutral envelope (Hasegawa et al. 2000), the typical abundance of C$_5$ in NGC~7027 would be (0.5--1.0)$\\times$10$^{-7}$, similar to that derived by Bernath et al. (1989) in IRC+10216.  Fig.~2 shows that, for the same column densities required to reproduce the bending modes, the $\\nu_4$ and $\\nu_3$ stretching modes of C$_6$ and C$_5$ are not detected in NGC~7027. Different excitation conditions of the stretching and bending modes and geometrical effects may explain this conjecture. Although carbon clusters have a moderate size, we could also expect an efficient pumping of the stretching and bending modes through UV photons. Mid--IR photons could also contribute to the excitation of the high--$\\nu_{bending}$ excited states but also of the stretching modes. Finally, the low--energy bending modes could be pumped alone by the absorption of far--IR dust photons (see the case of the C$_3$ excitation in  Cernicharo et al. 2000). However, the smaller volume of the photon--dominated regions relative to the far--IR dusty envelope in NGC~7027, and the larger mid--IR flux of IRC+10216 (where the C$_3$ and C$_5$ stretching modes were detected) compared to NGC~7027 (a factor $\\sim$10$^3$ at 5~$\\mu$m, but only a factor $\\sim$4 at 100~$\\mu$m), can favor the observation of the C$_n$ bending modes within the large ISO beam. Under these conditions, the detection of the bending modes and the non--detection of the stretching modes in NGC~7027 can be plausible.\\\\      The $\\nu_9$ low--energy bending mode of  C$_6$ is even less known. \\textit{Ab inito} calculations predict a $\\nu_9$$\\sim$108~cm$^{-1}$ frequency and an infrared intensity of $\\sim$25~km~mol$^{-1}$ [$A$($\\nu_9=1-0$)$\\simeq$0.035~s$^{-1}$; Martin et al. 1990; 1995), while photoelectron spectroscopy  experiments estimate $\\nu_9$=97$\\pm$45~cm$^{-1}$ (Arnold et al. 1991) and $\\nu_9$=90$\\pm$50~cm$^{-1}$  (Xu et al. 1997). Theoretical predictions have to be taken into account with caution since they rely in the harmonic approximation which can be in error for large amplitude benders such as radicals with low-lying electronic states. In particular, the infrared intensity has to be considered as a lower limit (and column densities as upper limits). \\textit{Ab initio}  calculations do predict a low-lying electronic state for open--shell species such as C$_4$ or C$_6$. The term value for the $^{1}\\Delta_g$ lowest energy excited state of C$_6$ has been recently determined by Xu et al. only at  $\\sim$1400~cm$^{-1}$  from the ground. Hence, A$_{SO}$  can effectively be large even in the ground state ($\\Lambda$=0). As an example, the A$_{SO}$ constant of C$_4$H ($^{2}\\Sigma$ ground state) is very large, $\\simeq$3~cm$^{-1}$ (Yamamoto et al. 1987), because its $^{2}\\Pi$ electronic excited state is only $\\simeq$468~cm$^{-1}$ above the ground. In the case of even--numbered C$_n$ chains, A$_{SO}$ can easily be larger because of the higher  spin multiplicity of the ground state and the larger $\\Lambda$ value of the lowest--lying electronic excited state. Fig.~1 also shows the expected band--shape for a $^{3}\\Pi_{vib}-{^3}\\Sigma_{vib}$ transition with A$_{SO}$=3.5~cm$^{-1}$ (we estimated A$_{SO}$$\\sim$4~cm$^{-1}$ for C$_4$), C$_6$  molecular constants  from  Hwang et al. (1993) and van Zee et al. (1987;1988), and  the $\\nu_9$ band origin at 98.3~cm$^{-1}$ (101.7~$\\mu$m). Assuming the same excitation temperature than for C$_5$ we derive $N$(C$_6$)=0.8$\\times$10$^{14}$~cm$^{-2}$. Hence, the C$_6$ abundance will be  a factor 2 smaller than that of C$_5$.  Much more  work has to be done to fully understand and characterize the low--energy vibrations of C$_n$ chains in order to assign the UfIBs observed by ISO. The  cyclic C$_n$ isomers could also produce spectral features in the far--IR, however, their active modes are even less known. Waiting for such progresses, space observations offer the unique opportunity to obtain spectra in the far--IR domain where the bending modes appear and to motivate further laboratory and theoretical studies. The presence of small C$_n$ chains  in the space as well as their high reactivity, suggest that these species can be involved in the formation of more complex organic molecules. In fact, for cumulenic clusters with $n$=10 to 20, linear structures are thought to close into rings, while for  $n$$>$30, these species are thought to be more stable in aromatic and fullerene--like structures (O' Brien et al. 1987).  Future space heterodyne telescopes, such as the \\textit{Herschel Space Observatory}, with much better sensitivity and spectral resolution in the far--IR, should allow the detection of longer C$_n$ chains through their low--energy bending modes. This will be the fingerprint needed to understand  the formation and the nature of the UIBs and UfIBs carriers.  \\clearpage"
    },
    "0403/astro-ph0403104_arXiv.txt": {
        "abstract": "We use a physically motivated semi--analytic model, based on the mass function of dark matter halos, to predict the number of radio--loud quasars as a function of redshift and luminosity.  Simple models in which the central BH mass scales with the velocity dispersion of its host halo as $M_{\\rm bh}\\propto \\sigma_{\\rm halo}^5$ have been previously found to be consistent with a number of observations, including the optical and X--ray quasar luminosity functions.  We find that similar models, when augmented with an empirical prescription for radio emission, overpredict the number of faint ($\\sim 10\\mu$Jy) radio sources by 1--2 orders of magnitude. This translates into a more stringent constraint on the low--mass end of the quasar black hole mass function than is available from the Hubble and Chandra Deep Fields. We interpret this discrepancy as evidence that black holes with masses $\\lsim 10^7~{\\rm M_\\odot}$ are either rare or are not as radio-loud as their more massive counterparts.  Models that exclude BHs with masses below $10^7~{\\rm M_\\odot}$ are in agreement with the deepest existing radio observations, but still produce a significant tail of high--redshift objects.  In the 1-10GHz bands, at the sensitivity of $\\sim 10\\mu$Jy, we find surface densities of $\\sim 100$, $\\sim 10$, and $\\sim 0.3$ deg$^{-2}$ for sources located at $z>6$, $10$, and $15$, respectively.  The discovery of these sources with instruments such as the {\\em Allen Telescope Array (ATA)}, {\\em Extended Very Large Array (EVLA)}, and the {\\em Square Kilometer Array (SKA)} would open a new window for the study of supermassive BHs at high redshift.  We also find surface densities of $\\sim 0.1$ deg$^{-2}$ at $z > 6$ for mJy sources that can be used to study 21 cm absorption from the epoch of reionization.  We suggest that, although not yet optically identified, the FIRST survey may have already detected $\\sim 10^3-10^4$ such sources. \\\\ \\\\ ",
        "introduction": "\\label{sec:intro}  The past few years have seen significant progress in probing the ultra--high redshift universe, with both galaxies (e.g. Spinrad et al. 1998; Hu et al. 2002; Rhoads et al. 2003; Kodaira et al. 2003) and quasars (Fan et al. 2000, 2001, 2003) being discovered in increasing numbers around and beyond redshift $z=6$ (see recent reviews by Spinrad 2003 and Taniguchi 2003).  In hierarchical structure formation scenarios in cold dark matter (CDM) cosmologies, the first baryonic objects appear at still higher redshifts: at $z\\approx 20-30$, when the first high--$\\sigma$ peaks collapse near the Jeans scale of $\\sim 10^5~{\\rm M_\\odot}$ (Haiman, Thoul \\& Loeb 1996; see Barkana \\& Loeb 2001 for a recent review).  Radiative cooling is efficient in the dense gas that has collapsed on these scales, and in principle, it can facilitate efficient formation of stars and black holes (BHs). Indeed, significant activity must have taken place at high redshifts, in order to reionize the intergalactic medium (IGM) by $z\\sim 15$ (Spergel et al. 2003).  The deepest detections of galaxies and quasars to date have been obtained at optical or near infrared (NIR) wavelengths, where the objects were identified in broad--band filters by their continuum, or in narrow--band imaging observations by their Lyman--$\\alpha$ emission.  The expected number of faint sources in future, deep NIR observations have been studied extensively in the context of hierarchical structure formation, using simple semi--analytic models. Haiman \\& Loeb (1997; 1998) showed that if halos collapsing at high redshifts have reasonable star or BH formation efficiencies, they can be detected in the NIR continuum in great numbers, with surface densities possibly reaching $\\sim1000$ sources per arcmin$^{-2}$ out to redshifts $z \\gsim 10$.  Similar models predict that a few $z>6$ quasars per $\\sim 100$ arcmin$^{-2}$ could be visible in soft X--ray bands at the flux limits already accessible to deep ${\\it Chandra}$ and ${\\it XMM}$ observations (Haiman \\& Loeb 1999; Wyithe \\& Loeb 2003).  Predictions analogous to those above, based on physically motivated structure formation models, are currently lacking in the radio bands. Observations of $z \\lsim 6$ quasars have established that a significant fraction ($\\sim 10\\%$) of these objects are bright in the radio.  Although some of the detailed physics responsible for this emission remains elusive, it is known to be a direct consequence of outflows generated by accretion onto a central massive BH.  In this scenario, one would expect that the population of radio loud quasars extends to $z\\gg 6$, to the epoch when the first BHs appeared and started to accrete.  In fact, if the radio--loudness distribution does not evolve strongly at high redshift (Ivezic et al. 2002; Petric et al. 2003), and if high redshift supermassive BHs (SMBHs) radiate close to their Eddington limit, then black holes with masses as small as $M\\sim 10^6{\\rm M_\\odot}$ should have radio flux densities that are already being reached by the deepest existing observations ($\\sim 10\\mu$Jy at $\\sim$ GHz frequencies).  {\\it The purpose of this paper is to confront simple models of the radio-loud quasar population with current observations and to put forward predictions for the counts at redshift and flux thresholds that are beyond the current observational limits.} While such extrapolations are necessarily uncertain, the detection of these objects would provide important constraints on the formation and growth of the first SMBHs (see Haiman \\& Quataert 2004 for a recent review).  Our predictions are especially timely in light of recent data on the reionization history of the intergalactic medium (IGM).  On the one hand, SDSS quasar spectra imply that we may have reached the neutral epoch at $z\\sim 7$ (e.g. Fan et al. 2002).  On the other hand, observations of the cosmic microwave background (CMB) anisotropies suggest that the universe was significantly ionized as early as $z\\sim 15$ (Kogut et al. 2003; Spergel et al. 2003).  This behavior is a challenge to reionization models. Pinning down the value of the neutral fraction just beyond $z\\sim 6$ would be of great value in elucidating these models (see, e.g. Haiman 2003 for a recent review). One promising probe of neutral hydrogen at high redshift is to study redshifted 21cm absorption and emission features.  Background density fluctuations from the ``21cm forest'', both in absorption (Carilli et al. 2002) and in emission (Madau et al. 1997); absorption (Furlanetto \\& Loeb 2002) or emission (Iliev et al. 2002) from neutral gas in discrete minihalos; and a sharp step--feature analogous to the Gunn-Peterson trough (Shaver et al. 1999) could all, in principle, be detected against a bright enough background source.  Such observations would provide a powerful probe of the amount and distribution of neutral hydrogen in the high--redshift IGM.  These studies will obviously depend critically on the number of available radio sources (although 21 cm features can also be studied against the CMB; e.g., Tozzi et al. 2000).  The rest of this paper is organized as follows.  In \\S~\\ref{sec:quasars} we summarize the phenomenology of radio--loud quasars, and in \\S~\\ref{sec:model}, we briefly discuss the model we adopt to describe their abundance and evolution.  In \\S~\\ref{sec:BHmass}, we compare this model to existing observations and in \\S~\\ref{sec:predict} we present our predictions for even higher redshift.  In \\S~\\ref{sec:discuss} we conclude with a discussion of our results and their implications.  Throughout this paper, we adopt the background cosmological parameters as measured by the {\\it WMAP} experiment, $\\Omega_m=0.27$, $\\Omega_{\\Lambda}=0.73$, $\\Omega_b=0.044$, $h=0.71$, $\\sigma_{8h^{-1}}=0.9$ and $n_s=1$ (Spergel et al. 2003). ",
        "conclusions": "\\label{sec:discuss}  We have presented a simple physically motivated model for the SMBH population and its evolution that fits the optical/IR and X-ray quasar luminosity functions out to $z \\approx 5$, and the luminosity function and number counts of bright radio sources at high redshift.  This model significantly overpredicts the counts at the $\\approx 10 \\ \\mu$Jy level in deep radio observations.  This discrepancy can be reconciled by postulating the existence of a lower limit to the SMBH mass, below which SMBHs are either rare or do not produce as much radio emission as their more massive counterparts.  We find that this lower limit is $\\approx 10^7~{\\rm M_\\odot}$, a constraint that is approximately an order of magnitude more stringent than that available from either the {\\it Hubble} or {\\it Chandra} deep fields (e.g., Haiman et al. 1999).  This constraint is especially interesting since there are several SMBHs in local galaxies with dynamically determined masses $\\lsim 10^7 M_\\odot$ (e.g. the central BHs in M32 -- van der Marel et al. 1998 -- and in the Milky Way -- Sch\\\"odel et al. 2002; Ghez et al. 2003).  In addition, Filippenko \\& Ho (2003) argued for a $\\sim 10^5 M_\\odot$ BH for the Seyfert galaxy in the late type (bulgeless) spiral NGC 4395, and Barth et al. (2004) reached a similar conclusion for the dwarf Seyfert 1 Galaxy POX 52.  If indeed SMBHs with masses below $\\approx 10^7 M_\\odot$ are rare, this could be because it is difficult to form SMBHs in shallow potential wells (e.g., Haehnelt et al. 1998; Haiman et al. 1999). Alternatively, during the coalescence of SMBH binaries, the remnant BH receives a ``kick'' velocity of up to several hundreds of km s$^{-1}$ due to the net linear momentum carried away by gravitational waves (e.g., Favata et al. 2004).  This kick may be sufficient to unbind lower-mass BHs from their host galaxies, leading to a dearth of low-mass BHs in galactic nuclei (e.g., Madau \\& Quataert 2004; Merritt et al. 2004).  An alternative explanation for the discrepancy between our models and the number counts of faint radio sources is that their bright accretion phase (near Eddington) is significantly shorter than that of their high--mass counterparts (note, however, that the Salpeter time characterizing the growth of SMBHs is independent of BH mass), or that lower mass BHs with $M \\lsim 10^7 M_\\odot$ are intrinsically less radio loud (we find that a reduction by a factor of $\\approx 15$, or changing $\\bar R = 2.8$ to $\\bar R \\lsim 1.6$, is needed). There are indeed some suggestions in the literature that very massive BHs are preferentially radio loud (e.g., Laor 2000; Lacy et al. 2001). However, other analyses suggest instead that the strongest correlation is between radio loudness and Eddington ratio (e.g., Ho 2002), consistent with the assumptions used here.  Our results also show that even in the presence of a low--mass cutoff in the distribution of radio-emitting SMBHs (which we assume is independent of redshift), a significant number of sources at redshifts as high as $z\\sim 15$ could be detectable.  In the 1-10GHz bands, at the sensitivity of $\\sim 10\\mu$Jy, we find surface densities of $\\sim 100$, $\\sim 10$, and $\\sim 0.3$ deg$^{-2}$ for sources located at $z>6$, $10$, and $15$, respectively (Fig.~\\ref{fig:rlcounts2}).  These predictions for the cumulative counts are insensitive to our choice of average radio loudness $\\bar R$.  The reason is that a typical $10^7 M_\\odot$ BH radiating at Eddington at $z = 6$ produces $\\approx 100 \\ \\mu$Jy if $R \\approx 2.8$ and $10 \\ \\mu$Jy if $R \\approx 1.8$.  Thus even if we decrease the average radio flux by an order of magnitude, i.e., from $\\bar R \\approx 2.8$ to $\\bar R \\approx 1.8$, the total surface density of the faintest sources shown in Figure~\\ref{fig:rlcounts2} does not change significantly (i.e., although the predicted fluxes for most of the sources decreases from $100 \\ \\mu$Jy to $10 \\ \\mu$Jy, the total number of sources with $F_\\nu \\gsim 10 \\ \\mu$Jy is relatively unchanged).  Figure~\\ref{fig:alpha} shows that, except at very high $z$, these predictions are also insensitive to the choice of the spectral index $\\alpha$.  \\begin{figure}[t] \\plotone{fig6.eps} \\caption{This figure shows the sensitivity of our predictions for the $10\\mu$Jy and 1mJy counts to the assumed spectral slope $\\alpha$ of the radio spectrum $F_\\nu\\propto\\nu^{-\\alpha}$.  The counts for $\\alpha=0$ are independent of frequency.  When $\\alpha=0.5$ is assumed, however, the counts increase with decreasing frequency, as shown by a comparison between the predictions at 100 MHz, 1GHz, and 10 GHz (short--dashed, dotted, and long--dashed curves, respectively).} \\label{fig:alpha} \\end{figure}  A particularly interesting question is the redshift evolution of bright radio sources, which can be used to study redshifted 21cm absorption features and hence the reionization history of the universe.  Carilli et al. (2002) find that a $\\sim 6$mJy source (at a few 100 MHz) is needed to achieve the S/N necessary for 21 cm absorption studies.  For a flat spectrum ($\\alpha = 0$), our results (Figure~\\ref{fig:rlcounts2}) imply that $\\sim 2.5$ such sources should be available in a 10 deg$^2$ field at $z=6-7$, with $\\approx 2,000$ sources available over the full sky in the redshift range $8<z<12$. Figure~\\ref{fig:alpha} shows that the predicted counts are larger by a factor of few if we instead assume a steeper spectrum with $\\alpha = 0.5$.  It is also important note that since these bright mJy sources are produced by massive BHs, their counts are independent of the low-mass cutoff considered above.  They are, however, sensitive to an evolution in the average radio--loudness $\\bar R$ (in contrast to the counts of the faint sources). For example, decreasing $\\bar R$ from ${\\bar R}=2.8$ to ${\\bar R}=2.3$ beyond $z>6$ would decrease the number of 6mJy sources at $6<z<7$ and $8<z<12$ by a factor of $4.6$, and $6.3$, respectively.  Another important point to note is that our fiducial results, obtained assuming $\\bar R=2.8$, predict that the FIRST survey may have already detected $\\sim 10^3-10^4$ quasars at $\\sim 1$mJy from redshift $z\\gsim 7$.  The identification of these quasars is a challenge, but should, in principle, be feasible with deep optical/IR observations.  The extrapolation of the radio-source population to high redshifts is necessarily uncertain; nevertheless, the results presented here should serve as useful order of magnitude estimates. We also note that our assumption of $M_{\\rm bh}\\propto\\sigma^5$ is somewhat conservative; a shallower relation would imply that the typical BHs reside in lower--mass halos (with shorter lifetimes; see Haiman \\& Loeb 1998 and Haehnelt et al. 1998); their abundance would then decrease less rapidly at high redshifts.  An important issue in designing future surveys is the relative merits of area vs. depth. The DM halo mass function has an approximate power--law shape at low masses $dN/dM \\propto M^{-2}$, and the velocity dispersion scales as $\\sigma\\propto M^{1/3}$.  As a result, if $M_{\\rm bh}$ scales as $\\sigma^\\alpha$, and if the flux is proportional to $M_{\\rm bh}$, then the number of sources will go as $F dn/dF \\propto F^{-3/\\alpha}$.  The number of detections scales linearly with the solid angle $\\Delta\\Omega$, and with the observation time as $t^{1.5/\\alpha}$; area is therefore more important than depth for the empirically determined slopes of $\\alpha \\approx 4$ (Gebhardt et al. 2000; Merritt et al. 2001; Tremaine et al. 2002). In principle, at bright flux limits, corresponding to BH masses where the halo mass function turns over and drops exponentially, decreasing the flux threshold would result in a larger yield of sources.  However, Figure~\\ref{fig:rlcounts2} reveals that even at the brightest fluxes shown, the source counts increase linearly with the flux threshold, and it would therefore be more advantageous to cover a larger area.  To conclude, we briefly discuss how to identify the high redshift radio sources.  At low flux densities ($\\lsim 30 \\ \\mu$Jy), starbursts dominate over quasars in deep radio observations (e.g., Richards et al. 1998).  There is also a contribution from moderate redshift low--luminosity AGN.  AGN can be distinguished from starbursts by their flatter spectra and variability.  Isolating the high redshift sources, however, will require optical/IR followup.  Our typical $10 \\ \\mu$Jy source at $z \\approx 6$ and $z \\approx 10$ is powered by a $\\approx 10^7 M_\\odot$ BH (without the cutoff described above it would be $\\lsim 10^6 M_\\odot$, but we would then overpredict the number of such sources).  With the Elvis et al. (1994) spectral template, normalized to a bolometric Eddington luminosity, these sources would have a flux density of $\\sim 0.3 \\mu$Jy at observed wavelengths of $\\sim 1-5\\mu$m, or an AB magnitude of $\\sim 25.5$ (see, e.g., figure 1 in Haiman \\& Loeb 1998).  They should be detectable in moderate integrations with the {\\it Hubble} or {\\it Spitzer Space Telescopes}.\\footnote{The {\\it Spitzer Space Telescope} has a point--source sensitivity of $\\sim 3\\mu$Jy in a 10-second exposure at 3.6$\\mu$m, see http://ssc.spitzer.caltech.edu/irac/sens.html} The very bright $\\sim$mJy sources at $z \\approx 6-10$, relevant for 21 cm absorption studies, should have flux densities of $\\sim 3\\mu$Jy (or AB magnitudes of $\\sim 23$ at $\\sim 1-5\\mu$m), and should be detectable in short exposures with {\\it HST} or {\\it Spitzer}, and potentially from the ground as well. The discovery and confirmation of even a few radio sources at $z>10$ by instruments such as the {\\em Allen Telescope Array (ATA)}, {\\em Extended Very Large Array (EVLA)}, and the {\\em Square Kilometer Array (SKA)} would open a new window for the study of supermassive BHs at high redshift, and of the pre--reionization universe.  Of even more immediate interest is the prediction from our models that, although not yet optically identified, the FIRST survey may have already detected $\\sim 10^3-10^4$ distant $z>7$ quasars.  Deep surveys, such as NOAO Deep Wide-Field Survey (NDWFS), VIRMOS and DEEP may cover the area necessary to identify a handful of these high-$z$ FIRST sources."
    },
    "0403/astro-ph0403618_arXiv.txt": {
        "abstract": "We review the continuum emission of Active Galactic Nuclei (AGN) over the entire electromagnetic spectrum. After a brief historical introduction, we describe the main spectral properties of unobscured AGN, discussing the selection biases which prevent us from having a complete view of the AGN population in the universe, and trying to build an updated spectral energy distribution of optically selected quasars. In the second part of the review, we describe the spectral properties of obscured AGN. Finally, we discuss the main observational methods in the different wavelength bands for disentangling AGN and stellar emission, and the ability of these methods to find new (mainly obscured) AGN, a significant fraction of which are probably still missing in current surveys. ",
        "introduction": "Active Galactic Nuclei (AGN) shine over $\\sim20$ decades of the electromagnetic spectrum, from the radio to the gamma rays. In almost all of this huge energy range, AGN are the brightest sources in the sky, except for the relatively narrow ($\\sim$3 decades) range from the infrared to the ultraviolet (UV).  The ``modern history'' of quasars is closely linked with the development of non-optical astronomy. In the early 60s, hard won precise positions of radio sources enabled the large optical telescopes of the day to take spectra of whatever optical object lay at that position. Some showed the starlight of normal galaxies, but often at large distances. In 1963 Marteen Schmidt made the ``official'' discovery of the first quasar, 3C~273, realizing that the point-like counterpart of a powerful radio source was at a redshift $z=0.16$, implying an enormous luminosity. In the early sixties, the availability of new space-borne instruments also opened the new field of astronomical observations in the X-rays, which soon revealed a sky dominated by AGN.  It is now clear that the initial population of quasars discovered in radio surveys is only a small fraction of a 10 times more numerous class of quasars, most of which are ``radio quiet''. These sources are, in turn, a small fraction of the total AGN population, which is dominated by obscured sources that can only be detected through their hard X-ray emission or through their reprocessed radiation in the infrared.  In this chapter, we describe the main continuum properties of AGN across the whole electromagnetic spectrum, according to our present, still incomplete, knowledge. We do not treat the emission of the small minority of quasars known as ``blazars'', whose emission appears significantly altered by relativistic beaming effects. We focus on the {\\em observed\\/} properties of quasars from a phenomenological point of view. More emphasis on the physical processes responsible for the observed emission of quasars can be found in Armitage (this volume).  We start in \\S\\ref{risalitisectype1} with unobscured (type I) AGN. We briefly discuss the main limitations in defining a representative sample of AGN. We then focus on a sample of bright, optically selected quasars (the Bright Quasar Survey, Schmidt \\& Green 1983) in order to build a mean Spectral Energy Distribution (SED) that takes into account all of the most recent observational results. This can be considered an update of the 10 year old compilation of Elvis et al.\\ (1994; hereafter, E94). We conclude \\S\\ref{risalitisectype1} with a brief overview of the possible dependence of the average spectral properties of type I AGN on redshift and/or luminosity and with a brief discussion on the often underestimated issue of the intrinsic dispersion of quasar spectra with respect to their average shapes.  Section~\\ref{risalitisecseds} is devoted to obscured (type II) AGN. We first review the main spectral properties for each energy band. We then discuss the relation between dust and gas absorption, describing some recently discovered sources for which the ratio between gas and dust, and the dust properties, are strongly different from those in our Galaxy. Finally, in \\S\\ref{risalitisectype2} we discuss the current methods that can be used to discover obscured AGN in the different energy bands. There are several indications that a significant fraction of these sources are still missing in current surveys. In particular, we focus on the techniques that can help to disentangle the AGN and stellar emission in the population of luminous sources that dominate the infrared and submillimeter sky. We conclude with a brief summary in \\S\\ref{risalitisecsummary}.  We finish this introduction with a note on terminology: historically, the term ``quasar'' has been used widely to refer to AGN of high luminosity whose emission completely overwhelmed that of the host galaxy. (Thus, these sources appeared as point sources in optical observations.) Low-luminosity AGN were (and still are) referred to as ``Seyfert Galaxies''. However, evidence for any basic physical difference between these types of active objects has diminished through the years, essentially to the vanishing point. For this reason, in this chapter, we will use the term ``quasar'' as synonymous with type 1 AGN. ",
        "conclusions": "\\label{risalitisecsummary}  We have reviewed the main spectral properties of quasars, with an emphasis on continuum emission and a brief treatment of the main emission and absorption features.  In the first part, we built an updated SED of quasars, mainly based on continuum emission of local ($z<0.4$) PG quasars. The main improvements with respect to the previous reference work in the field, the atlas of Elvis et al.\\ (1994), are: 1) the treatment of selection effects in the X-rays, which allowed us to estimate the correct optical-to-X-ray flux ratio for local quasars, 2) the extension of the SED in the UV at wavelengths shortward of the Lyman break, using {\\em HST\\/} observations of quasars, 3) the inclusion of the new IR data obtained with {\\em ISO\\/}, and 4) the use of a larger sample in the computation of the average spectral properties of quasars.  In the second part, we reviewed the main spectral properties of obscured quasars. We also discussed the relation between dust and gas absorption, briefly describing the cases of X-ray obscured objects with no or little optical/near-IR absorption.  In the third and final part, we summarized the methods for finding obscured AGN in the universe, focusing on the problem of disentangling the active nucleus component in infrared luminous sources and disucssing the still open problems. Despite the improvements of the last few years, what we know is most likely only a fraction of the obscured quasars present in the universe."
    },
    "0403/astro-ph0403332_arXiv.txt": {
        "abstract": "The unique high--resolution wide--field imaging capabilities of HST with ACS have allowed the characterization of galaxies at redshift 6, less than 1 Gyr from recombination.  The dropout technique, applied to deep ACS $i$, $z$ images in the RCDS 1252--2927, GOODS and UDF--Parallel fields has yielded large samples of these objects, allowing a detailed determination of their properties (e.g., size, color) and meaningful comparisons against lower redshift dropout samples.  The use of cloning techniques has enabled us to control for many of the strong selection biases that affect the study of high redshift populations.  A clear trend of size with redshift has been identified, and its impact on the luminosity density and star formation rate can be estimated.  There is a significant, though modest, decrease in the star formation rate from redshifts $z\\sim 2.5$ out through $z\\sim 6$. The latest data also allow for the first robust determination of the luminosity function at $z\\sim 6$. ",
        "introduction": "The advent of the HST Advanced Camera, the ACS (Ford et al 2003) has greatly increased our ability to ``watch galaxies form''.  The sensitivity, resolution and excellent filter set have provided us with images from which large samples of high redshift galaxies can be derived.  Of particular interest are those galaxies with red enough {\\it i-z} colors to qualify as {\\it i}-dropouts -- galaxies at redshifts $z \\sim 6$.  Such objects have been the focus of a number of papers over the last year (e.g., Bouwens et al 2003b, Stanway et al 2003, Yan et al 2003, Dickinson et al 2004). Spectroscopic confirmation is beginning to appear (e.g., Bunker et al 2003, Dickinson et al 2004) but is challenging as Weymann et al (1998) demonstrated with their $z=5.6$ object, which took over 6 hours on Keck.  The current frontier for high redshift objects is at $z\\sim 6$ (the ACS UDF and NICMOS UDF--IR images together will likely extend the dropout sources to redshifts 7 and beyond, but the samples will be small).  Rapid changes in the properties of high redshift galaxies must occur beyond $z\\sim6$ and so careful characterization of objects even those separated by small intervals of time, is a well-justified goal.  Given this, there is great value in having large samples of $z\\sim3-5$ objects to contrast with $z\\sim6$ galaxies.  Though only $0.2-1.0$ Gyr later in cosmic history, $z\\sim3-5$ object samples are larger and much better characterized, providing key information on evolutionary changes in high redshft galaxies.     \\begin{figure} \\begin{center} \\begin{minipage}[c]{0.5\\linewidth} \\includegraphics[width=1.0\\textwidth]{gdiF1.eps} \\end{minipage}\\hfill \\begin{minipage}[c]{0.45\\linewidth} \\caption[]{Selection of {\\it i}-dropouts in the  {\\it i-z} {\\it z-J} two--color plane.  This example is for the HST ACS from the RCDS 1252--2927 field (Rosati et al 1998). The selection limits (particularly the $(i-z)>1.5$ cut -- see Bouwens et al 2003b) returns $z \\sim 6$ galaxies with little contamination (an estimated 11\\% contamination rate).} \\end{minipage} \\end{center} \\end{figure} ",
        "conclusions": "There has been a remarkable growth in the number of objects known at high redshifts ($z\\sim 6$) since the HST Advanced Camera came into operation after servicing mission SM3B. Not only are large numbers of sources being detected at high redshift, but the development of new techniques for detecting, characterizing and comparing high redshift objects from photometric datasets has led to many quantitative results on the nature and evolution of galaxies in the first $\\sim1-3$ Gyrs.  \\smallskip  {\\bf Acknowledgments} We would like to thank the organizers for an excellent meeting in a great place. We acknowledge the remarkable advances that have come about because of HST and its amazing imagers, and regret the decision to cancel SM4 that will lead to the premature death of HST. We owe a lot to our team members on the ACS GTO team and the UDF--IR team, and particularly the PIs, Holland Ford and Rodger Thompson.  Support from NASA grant NAG5--7697 and  NASA/STScI grant HST--GO--09803.05--A is gratefully acknowledged.   \\def\\apj{ApJ.} \\def\\apjl{ApJ.} \\def\\aap{A\\&A} \\def\\mnras{MNRAS}"
    },
    "0403/astro-ph0403568_arXiv.txt": {
        "abstract": "We discuss a general mechanism which allows to explain naturally both radio and high energy emission by pulsars. We also discuss the plasma distribution in the region surrounding the pulsar, the pulsar wind and the formation of jet along the magnetic axis. We suggest a plausible mechanism to generate pulsar proper motion velocities. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403042_arXiv.txt": {
        "abstract": "We report the discovery of a dwarf (M$_B = -13.9$) post-starburst galaxy coincident in recession velocity (within uncertainties) with the highest column density absorber (N$_{\\rm H~I} = 10^{15.85}$ cm$^{-2}$ at $cz = 1586$ \\kmsno) in the 3C~273 sightline. This galaxy is by far the closest galaxy to this absorber, projected just 71h$^{-1}_{70}$ kpc on the sky from the sightline. The mean properties of the stellar populations in this galaxy are consistent with a massive starburst $\\approx$3.5 Gyrs ago, whose attendant supernovae, we argue, could have driven sufficient gas from this galaxy to explain the nearby absorber. Beyond the proximity on the sky and in recession velocity, the further evidence in favor of this conclusion includes both a match in the metallicities of absorber and galaxy, and the fact that the absorber has an overabundance of Si/C, suggesting recent type II supernova enrichment. Thus, this galaxy and its ejecta are the expected intermediate stage in the fading dwarf evolutionary sequence envisioned by Babul \\& Rees to explain the abundance of faint blue galaxies at intermediate redshifts. While this one instance of a QSO metal line absorber and a nearby dwarf galaxy is not proof of a general trend, a similar dwarf galaxy would be too faint to be observed by galaxy surveys around more distant metal line absorbers. Thus, we cannot exclude the possibility that dwarf galaxies are primarily responsible for weak (N$_{\\rm H~I} = 10^{14-17}$ cm$^{-2}$) metal line absorption systems in general. If a large fraction of the dwarf galaxies expected to exist at high redshift had a similar history (i.e., they had a massive starburst which removed all or most of their gas), these galaxies could account for at least several hundred high-$z$ metal-line absorbers along the line-of-sight to a high-$z$ QSO. The volume filling factor for this gas, however, would be $< 1\\%$. ",
        "introduction": "The relationship between \\lya absorbers and galaxies has been debated in the literature for many years \\citep[see e.g.,][]{bahcall96,arons72,weymann81,bergeron91,lanzetta95,steidel98,dave99} and is the source of on-going controversy \\citep[see recent volume ``Extragalactic Gas at Low Redshift'';][]{mulchaey02}. The two-point correlation function (TPCF) of absorbers is the only available datum on this relationship at $z\\geq$1.6, because at those redshifts even bright galaxies are hard to detect and study \\citep[although see][]{adelberger03}. The latest TPCF results continue to draw a distinction between the \\lya-only (i.e., apparently metal-free) absorbers and those which also contain metals \\citep[e.g., \\ion{C}{4} 1548\\AA;][]{womble96}. While the metal-line systems show a significant excess power in the TPCF at low velocity differences, the \\lya-only absorbers show very little evidence for clustering in velocity \\citep[i.e., only a 3-$4\\sigma$ excess at $\\Delta cz \\leq 200$ \\kmsno;][]{rauch92,hu95,cristiani97}. A similar result is found at $z\\sim0$ for low column density \\lya absorbers by \\citet[, Papers I-IV hereafter]{penton00a,penton00b,penton02,penton04}. This result is interpreted to mean that metal-line systems are related to galaxies (e.g., bound galactic disks, halos or outflowing winds) while the \\lya-only absorbers are more uniformly distributed in space and may not be related to galaxies at all.  At high-$z$, absorption lines of \\ion{C}{4} and/or \\ion{O}{6} are visible in individual absorption systems down to \\ion{H}{1} column densities of N$_{\\rm H~I} \\approx 10^{14.5}$ cm$^{-2}$ \\citep{songaila96,telfer02}. And, based upon analysis techniques that use many, very weak absorption systems in concert \\citep[e.g., ][]{cowie98,ellison00,schaye00,telfer02,simcoe04}, metals may be present in a substantial fraction of the still lower column density systems. These TPCF and metal abundance results suggest, but do not prove, that even the uniformly-distributed, lower column density \\lya absorbers could contain metals at a level of 10$^{-2}$ to 10$^{-2.5}$ Solar. If this inference is correct, a very early epoch of star formation may have ``enriched'' the intergalactic medium (IGM) more or less uniformly (but see \\citealp{simcoe04} for evidence against this conclusion).  At the highest hydrogen column densities ($\\geq 10^{17.3}$ cm$^{-2}$) the association between galaxies and absorbers is much clearer. Imaging and spectroscopy of strong \\ion{Mg}{2} metal-line systems at moderate $z$ \\citep{bergeron91,steidel95,steidel98} have found strong evidence for the association of metal-bearing absorbers with luminous galaxies brighter than $\\sim 0.1$L$^*$. In Steidel's (1995, 1998) survey of 55 \\ion{Mg}{2} absorbers, 53 have galaxies brighter than $\\sim 0.1$L$^*$ within 50h$_{70}^{-1}$ kpc in projection. Since the high equivalent-width \\ion{Mg}{2} systems are invariably optically-thick at the \\ion{H}{1} Lyman limit, these systems must include both Lyman limit and damped \\lya absorption systems. The two absorbers in the \\citet{steidel95} sample lacking ``associated'' bright galaxies are damped systems \\citep{steidel97} and it is now known that some low-$z$ damped \\lya absorbers are associated with dwarf \\citep{nestor02} and even low surface brightness (LSB) galaxies \\citep{bowen01}. A new blind \\ion{H}{1} 21cm survey \\citep{rosenberg02} confirms that a very broad range of \\ion{H}{1} masses and thus optical and near-IR luminosity galaxies contribute to the damped \\lya absorber population \\citep{rosenberg03a}. These observations are strongly suggestive that strong \\ion{Mg}{2}/Lyman limit absorbers are due to metal-rich gas in bound halos of brighter galaxies ($\\geq 0.1$L$^*$), while damped absorbers are due to metal-rich gas in disks of galaxies of all descriptions and luminosities.  Weaker (N$_{\\rm H~I}\\leq 10^{17}$ cm$^{-2}$) low-ionization, (\\ion{Mg}{2}, \\ion{Fe}{2}) and high-ionization (\\ion{C}{4}, \\ion{Si}{4}, and \\ion{O}{6}) metal-line absorbers are more numerous and poorly understood \\citep[see e.g.,][]{lanzetta95,charlton02,chen00}. One might speculate that these weaker and higher-ionization absorbers are the more diffuse, outer regions of bound halos or, perhaps, the halos or winds associated with less luminous galaxies than found by the \\citet{steidel98} survey \\citep[see e.g.,][]{lanzetta95}. Particularly enigmatic are the low ionization, weak \\ion{Mg}{2} systems \\citep{rigby02}. While there is little imaging and spectroscopy of these systems that might reveal ``associated'' galaxies, a photoionization model leads to surprisingly high densities and thus small sizes \\citep[$\\leq 1$ kpc;][]{charlton02,tripp02,rosenberg03b}, at least along the line-of-sight. Suggestions concerning the enigmatic nature of these very numerous, low-ionization absorbers include extragalactic starforming regions, cool clumps in bound galaxy halos similar to Galactic High Velocity Clouds (HVCs), and cool clumps in unbound galaxy winds. Regardless, if the extremely small inferred cloud depths are indicative of the overall sizes of these weak \\ion{Mg}{2} systems, they would be extremely numerous (greatly exceeding the number of bright galaxies) and difficult to understand in any context \\citep{charlton02}.  {\\it Hubble Space Telescope} (HST) ultraviolet spectroscopy stands to make an essential contribution to our understanding of the relationship between galaxies and the variety of absorbers described above, because some of the HST-discovered absorbers are at sufficiently low redshift that their surroundings can be scrutinized for galaxies in great detail. Statistical evidence for the relationship between low-$z$ absorbers and galaxies is based upon two major HST spectroscopic surveys: (1) the HST QSO Absorption Line Key Project \\citep{bahcall93,bahcall96,weymann98,jannuzi98} used the Faint Object Spectrograph (FOS) to study high column density (log N$_{\\rm H~I} \\geq 14$ cm$^{-2}$) absorbers and (2) Papers~I-IV used the Goddard High Resolution Spectrograph (GHRS) and the Space Telescope Imaging Spectrograph (STIS) to study lower column densities (log N$_{\\rm H~I} = 12.5$ to 16 cm$^{-2}$) but over a shorter pathlength ($\\Delta z \\sim 1$ compared to $\\Delta z \\sim 3.5$ at $z \\leq 0.3$ for the Key Project spectroscopy). The ground-based galaxy survey work of \\citet{lanzetta95} based upon the FOS Key Project absorber sample finds that $\\sim 30\\%$ of these high column density absorbers can be identified with bright galaxies at impact parameters $\\leq 160$h$_{70}^{-1}$ kpc. \\citet{lanzetta95} and \\citet{chen98} extrapolate this result to conclude that all high equivalent width \\lya absorbers (N$_{\\rm H~I}\\geq 10^{14.5}$ cm$^{-2}$) can be identified as the very extended halos of galaxies, although most would be too faint to be detectable at the distance of the Key Project absorbers. On the other hand, investigations of the distribution of galaxies near lower column density (N$_{\\rm H~I} \\leq 10^{14.5}$ cm$^{-2}$) \\lya absorbers, first by \\citet[M93 hereafter]{morris93}, then by Paper~III \\citep[see also][]{tripp98,impey99} found little evidence for a close association between \\lya absorbers and galaxies, although there was some evidence for an association between most ($\\sim 78\\%$) \\lya absorbers and larger-scale galaxy filamentary structures; the remaining 22\\% of these absorbers were found in galaxy ``voids'' \\citep[Paper~III and][]{mclin02a}. These latter results have suggested to some investigators (including the current authors), that the proximity of high column density absorbers found by \\citet{lanzetta95} is not a cause-effect relationship between the absorber and its nearest galaxy, but rather is caused by both galaxies and absorbers being parts of the same large-scale structure filament. This conclusion is supported both by the theoretical N-body simulations \\citep[e.g.,][]{dave99} and by the nearest neighbor analysis of Paper~III.  In an effort to utilize the most nearby \\lya absorbers to study the absorber/galaxy relationship, we have initiated a deep pencil-beam galaxy survey around sightlines where very nearby ($cz \\leq 10,000$ \\kmsno) \\lya absorbers have been found in the Penton et~al. survey. This sample is ideal for studying the galaxies surrounding low column density \\lya -only absorbers and the weak metal-line absorbers because even very faint galaxies can be inventoried. For example, a first result from this survey \\citep{mclin02a} finds no galaxy close to several \\lya absorbers at $cz \\leq 3000$ \\kms in galaxy voids, in some cases with absolute mag limits as faint as M$_B = -12.5$. In this Paper we report on a second specific result from our deep optical galaxy survey involving one of the nearest and strongest \\lya absorbers in the Penton et~al. study: a low ionization metal bearing absorber at Virgo distance ($cz$=1586 \\kms) along the 3C~273 sightline \\citep[T2002 hereafter]{bahcall91,morris91,weymann95,sembach01,tripp02}. While our overall survey results will be presented elsewhere (Stocke et al. 2004), here we present the discovery of the closest galaxy-absorber association found by our survey, a dwarf post-starburst galaxy 71h$_{70}^{-1}$ kpc on the sky and $\\Delta cz = 50 \\pm 50$ \\kms away from a high column density (log N$_{\\rm H~I} = 15.85$ cm$^{-2}$) \\lya absorber. This one absorption system has more \\ion{H}{1} column than all of the other \\lya absorptions along this sightline, so it is at the very high end of the column density distribution of absorbers studied by our HST and galaxy surveys (Papers~I-IV).  This absorber is very similar to the ``weak \\ion{Mg}{2}'' absorbers studied by \\citet{charlton02}, as described above. T2002 use photoionization models to infer a cloud thickness of 70 pc, quite comparable to the small line-of-sight extents inferred by \\citet{rigby02} for weak \\ion{Mg}{2} systems. Thus, this absorber and its possibly related galaxy allow us to study weak metal-line absorption systems in close detail. In the next Section we describe the previous work on this sightline and the nearby sightline of \\RXJ\\ with HST as well as ground-based optical and radio \\ion{H}{1} 21cm galaxy survey work in this region. Section 3 describes the discovery and properties of this dwarf galaxy. Section 4 presents a discussion of these results in the context of other work and Section 5 summarizes our conclusions. As with the distances on the sky that we quote above, throughout this paper we will quote distances obtained from redshifts assuming a pure Hubble velocity field in this direction with $H_0$ = 70h$_{70}$ \\kms Mpc$^{-1}$. Using the Local Supercluster velocity field model of \\citet{tonry00} reduces distance values in this direction by 10\\%. ",
        "conclusions": "In this Paper we have reported the discovery of a dwarf, post-starburst galaxy 71h$^{-1}_{70}$ kpc away on the sky and at the same recession velocity to within the combined errors as the highest column density \\lya absorber in the 3C~273 UV spectrum obtained with HST and FOS, GHRS and STIS. A detailed study of this absorber using a STIS E140M spectrum \\citep{heap02} in conjunction with FUSE observations of higher order Lyman lines \\citep{sembach01} has been made by T2002. Based upon standard photoionization models, T2002 found a metallicity for this N$_{\\rm H~I}$ = 10$^{15.85}$ cm$^{-2}$ absorber of 6\\% Solar, an absorber thickness along the line-of-sight of 70 parsecs and an overabundance of Si to C of 0.2 dex relative to Solar. Using Lick absorption line indices measured from a ground-based spectrum of the nearby dwarf (M$_B = -14$) obtained at the Apache Point Observatory's 3.5m telescope, we estimate that the galaxy's stellar population has a mean age of 3.5 $\\pm$ 1.5 Gyrs and a mean metallicity of 3-30\\% Solar. Strong Balmer absorption (e.g., 2 \\AA\\ equivalent width of H$\\delta$) and the absence of [O II] and H$\\alpha$ emission, as well as the absence of \\ion{H}{1} 21 cm emission (M$_{\\rm H~I} \\leq 5 \\times 10^6$ M$_{\\Sun}$) are strongly suggestive, but cannot prove, that this galaxy experienced a massive starburst $\\sim$3.5 Gyrs ago, whose subsequent supernovae blew out all of its remaining gas. Given this interpretation, the evidence for a direct connection between this small galaxy and the nearby \\lya absorber is quite convincing: approximate matches in recession velocity and metallicity, and a Si/C ratio in the absorber suggesting recent supernova type II enrichment. Despite an extremely thorough search for faint galaxies near the 3C~273 sightline \\citep[M93;][and the present work]{rauch96,vangorkom93}, only this one dwarf has been found close to this absorber. The next nearest galaxies are two other dwarf galaxies $\\approx 250$h$_{70}^{-1}$ kpc away, so that the only viable alternative picture for the 3C~273 absorber is gas produced by at least several nearby galaxies, including this dwarf.  Optical (B,R) images of this dwarf obtained at the APO 3.5m show that, despite the galaxy's elliptical appearance, its surface brightness profile is much better fit by an exponential disk with scale length 0.4h$^{-1}_{70}$ kpc and no optical peculiarities in the B-band image that would indicate the presence of dust. If a massive starburst did occur in this galaxy $\\sim 3.5$ Gyrs ago, the type II supernovae would have had sufficient energy to accelerate several $\\times 10^8$ M$_{\\Sun}$ of \\ion{H}{1} to the speeds necessary to reach 70h$^{-1}_{70}$ kpc radius by now. Thus, although we cannot prove beyond all doubt that the atoms producing the absorber came from this dwarf galaxy due to a massive starburst superwind, this seems to be the most logical conclusion. Also, since spectral synthesis models of the evolving stellar population of starbursts \\citep{bruzual93} suggest that this galaxy was $\\sim 10$ times brighter in B-band during its starburst phase and will fade ten times more in the next few Gyrs, this galaxy is an example of the intermediate stage in the evolutionary process that \\citet{babul92} suggests begins as a blue, starbursting dwarf galaxy and ends as a dormant, dwarf spheroidal.  The primary uncertainty in this proposed scenario is that this galaxy experienced a single, massive starburst $\\sim$3.5 Gyrs ago, which accounts for its current luminosity and stellar properties. A starburst of this magnitude in a dwarf galaxy \\citep[requiring 1 M$_{\\Sun}$ yr$^{-1}$ for 10$^8$ yrs;][]{martin03} is unexpectedly large. Further, the indirect evidence from observations of the stellar populations of Local Group dwarf spheroidals is that low luminosity galaxies typically experience several epochs of star formation \\citep{mateo98}. However, even in the case that the starburst which took place in this galaxy accounts for only $\\sim 30\\%$ of all of its stars and occurred as recently as 2 Gyrs ago, a supernovae-driven wind can still account for the observed properties and location of the 1586 \\kms absorber towards 3C~273. A near-UV image of this dwarf galaxy could determine whether the massive, single starburst model is correct. \\citet{bruzual93} models of instanteous starbursts have little mid-UV flux more than 1 Gyr after the starburst, fading from 6 to 20 times less flux than at B-band between 1 and 4 Gyrs after the starburst. However, stellar synthesis models with several star formation epochs or continuous star formation have sufficient blue horizontal branch stars so that the mid-UV flux never gets this low compared to B-band \\citep[see e.g., \\citet{brown97} on the ``UV-upturn'' in ellipticals, which has an onset of $\\geq 7$ Gyrs after a starburst;][]{tantalo96}.  While the association of this dwarf galaxy with a \\lya + metal-line absorption system is interesting as an individual example of a ``Cheshire Cat'' galaxy \\citep{salpeter93,charlton95}, it also may be indicative of a large population of starbursting dwarf galaxies, whose supernovae can drive gas away from these dwarfs due to their shallow gravitational potential wells \\citep{maclow99,ferrera00}. We are currently searching for a faint galaxy near the \\RXJ\\ sightline that could be responsible for a very similar pair of metal-line absorbers at 1666 \\kms, $\\sim 350$h$^{-1}_{70}$ kpc away on the sky from the 1586 \\kms absorber in 3C~273. If a dwarf galaxy is present near \\RXJ\\, this will provide further proof in favor of the dwarf starburst superwind model for these absorbers. All other known metal line systems are too far away ($cz \\geq 3000$) to allow dwarfs similar to this one to be observed by current surveys for galaxies near QSO absorbers (i.e., at m$_B > 19$; M93). Therefore, metal-line absorbers would be identified with more luminous galaxies at greater impact parameters \\citep[e.g.,][]{chen01,adelberger03}, which happen to be in the same galaxy filament as a dwarf (or dwarves) closer to the sightline.  We view the discovery of this dwarf close to a \\lya + metal-line absorption system as strong evidence in favor of the \\citet{babul92} scenario used to explain the large population of ``faint, blue galaxies'' seen on deep optical images. The large numbers of faint, blue galaxies plus the theoretical expectations of standard heirarchical pictures of galaxy formation strongly suggest that there were many more of these dwarf galaxies in the past than now. If most of these dwarfs went through a violent starburst phase, dwarf starburst superwinds could be responsible for ${\\rm dN/dz} \\gtrsim 500$ \\lya + metal-line absorption systems in the spectra of high-$z$ QSOs. This estimate is based upon the assumption that all dwarf galaxies produce a 75h$^{-1}_{70}$ kpc radius shell of metal-enriched gas around them before they  merge to form larger galaxies. Evidence from other work \\citep[e.g., ][]{bergeron91,steidel95,steidel98} is that larger galaxies (${\\rm L} > 0.1$L$^*$) form bound halos of denser gas and may not contribute significantly to the spread of metals into the IGM. But, even if all dwarfs produce massive starburst superwinds, $< 1\\%$ of the IGM would become enriched in metals, leaving most of the \\lya forest absorbers metal-free at the lowest column densities. This conclusion can be tested by determining if \\lya absorbers in galaxy voids \\citep[e.g., ][]{mclin02a} contain metals at a level $\\geq 0.3\\%$ Solar.  We thank Erica Ellingson and the APO 3.5m observing specialists for assistance in obtaining the dwarf galaxy spectrum. We thank Carnegie Observatories for their support of the wide field CCD multi-object spectroscopy system and the WIYN queue observers for their excellent observing support. We thank Bianca Poggianti for the use of her figure from a previous publication, Yeong-Shang Loh for accessing the public Sloan Digital Sky Survey database to provide information for us on the galaxies near the 3C~273 sightline and Simon Morris for helpful discussions of his galaxy survey work near 3C~273. JTS, KMM, BAK and JLR acknowledge support from NASA HST General Observer grants GO-06593.01-A, GO-08182.01-A, GO-09506.01-A and archival research grant AR-09221.01-A. JLR also acknowledges support from an NSF Astronomy \\& Astrophysics Postdoctoral Fellowship, AST-0302049.    \\clearpage"
    },
    "0403/astro-ph0403274_arXiv.txt": {
        "abstract": "The Tycho supernova remnant (SNR), as one of the few historical SNRs, has been widely studied in various wavebands. Observations show evidence that Tycho is expanding in a medium with density gradient and possibly  interacting with dense ambient medium toward the northeast direction. From the FCRAO CO survey of the outer Galaxy, we have identified a patch of molecular clouds in this area and have conducted a follow-up observation with Nobeyama 45m radio telescope. The high-resolution (16\\arcsec) Nobeyama data shows that a large molecular cloud surrounds the SNR along the northeastern boundary. We suggest that the Tycho SNR and the molecular cloud are located in the Perseus arm and that the dense medium interacting with the SNR is possibly the molecular cloud. We also discuss the possible connection between the molecular cloud and the Balmer-dominated optical filaments, and suggest that the preshock gas may be accelerated within the cosmic ray and/or fast neutral precursor. ",
        "introduction": "The Tycho supernova remnant (SNR) is known as a remnant of a Type Ia supernova that occurred in the year 1572 \\citep{1945ApJ...102..309B}. As one of the few historical SNRs, the remnant has been widely studied in various wavebands. In radio continuum, it shows a shell morphology with clearly defined outer boundary and clumpy inner structure \\citep[e.g.,][]{1991AJ....101.2151D% }. Recent high resolution X-ray images also exhibit similar features \\citep{2001A&A...365L.218D,2002ApJ...581.1101H}. \\cite{1980ApJ...235..186C} % derived a shock velocity by modelling the optical spectrum from a fast nonradiative shock and compared it with an optical proper motion \\citep{1978ApJ...224..851K}% , giving an estimated distance of $2.3 \\pm 0.5$ \\kpc. Other independent estimates give similar distances \\citep{1986MNRAS.219..427A,1988MNRAS.230..331S}. It is, however, noted that there has been controversy in interpreting the \\hi absorption data, giving a different distance \\citep[e.g.,][]{1995A&A...299..193S}. This will be discussed in \\S\\ \\ref{distance}.  Observations show evidence that Tycho is interacting with ambient dense clouds toward its northeast (NE) direction. \\citet{1997ApJ...491..816R} has studied its expansion using a 10-year-lap VLA data and found that the NE area, which is farthest from the geometrical center of the SNR, is currently moving slowest, suggesting that the remnant is expanding into a higher density medium. \\citet{1999AJ....117.1827R} % has conducted a \\HI absorption study towards the remnant % and suggested that Tycho is currently interacting with a dense \\HI concentration in the NE ($\\vlsr = -51.5$\\ \\kms) which locally slows down the expansion of the shock front. There are other observations that also suggest its interaction with inhomogeneous environment \\citep{ 1998A&A...334.1060V, 2000ApJ...529..453K, 2001A&A...373..281D}.  From the FCRAO $^{12}$CO survey of the outer Galaxy \\citep{1998ApJS..115..241H}, we have identified molecular clouds % which could be associated with Tycho (Figure \\ref{fig:tycho_fcrao}). They are arc-shaped and the inner boundary roughly traces the NE rim of Tycho. The morphology of the molecular clouds suggests physical association with Tycho. We have conducted high-resolution $^{12}$CO J=$1\\!-\\!0$ line observations and present the results in this \\paper. ",
        "conclusions": "The study of Balmer-dominated filaments in SNRs has given us much information about the shock structure. The velocity of an unperturbed preshock medium is an invaluable observational ingredient for fully understanding underlying shock  structure. In Tycho, a velocity shift between an ambient molecular cloud, which is believed to be the unperturbed preshock gas, and the \\halpha\\ filaments has been observed. Although the preshock acceleration was searched for previously, no evidence has obtained before our observation \\citep[e.g.][]{1994PASP..106..780E}. Our observation serves as the first direct evidence for the acceleration of preshock gas by a cosmic ray (or fast neutral) precursor. We, however, note that our conclusion is based on the single velocity measurement and more observations are needed to test it.   There has been not much study on the structure of collisionless shocks propagating into  molecular clouds. But most of the volume of a molecular cloud may be occupied by low density ($\\lesssim 10\\ \\mathrm{cm}^{-3}$) interclump medium, and the structure would not much differ from that in atomic medium. The molecules are supposed to be dissociated within a precursor either by radiation from postshock region or hot electrons. The dissociation rate of \\hh\\ is only slightly larger than ionization rate of the hydrogen atom \\citep{1998ApJ...496.1044Y, 2000A&AS..146..157L}, and complete dissociation of \\hh\\ may not happen since they occupy more dense region. The velocity of the transmitted shock into a molecular clump can be sufficiently low to be radiative, but currently there is no sign of radiative shock in Tycho, presumably because the shock is propagating into the low-density outskirts of the molecular cloud. Future study should address these issues together with the nature of cosmic ray acceleration for a high-velocity ($\\gtrsim 1000\\ \\kmss$) shock propagating into a molecular cloud."
    },
    "0403/astro-ph0403326_arXiv.txt": {
        "abstract": "We propose a new method of estimation of the black hole masses in AGN based on the normalized excess variance, \\rms. We derive a relation between \\rms, the length of the observation, \\T, the light curve bin size, $\\Delta t$, and the black hole mass, assuming that (i) the power spectrum above the high frequency break, \\nb, has a slope of $-2$, (ii) the high frequency break scales with black hole mass, (iii) the power spectrum amplitude (in $frequency \\times power$ space) is universal and (iv) \\rms\\ is calculated from observations of length $T < 1/$\\nb. Values of black hole masses in AGN obtained with this method are consistent with estimates based on other techniques such as reverberation mapping or the \\bhm-stellar velocity dispersion relation. The method is formally equivalent to methods based on power spectrum scaling with mass but the use of \\rms\\ has the big advantage of being applicable to relatively low quality data. ",
        "introduction": "The X-ray emission of active galactic nuclei (AGN) displays variations over a wide range of time scales.  The first convincing demonstration of this phenomenon came with the {\\it EXOSAT} `long looks' (Lawrence \\etal\\ 1987; McHardy \\& Czerny 1987). The data showed no characteristic time-scales, and the power spectral density function (PSD) showed a `red-noise' shape. At the same time, it was also noticed that more luminous sources show slower variations. Various methods have been used in the past in order to determine the X-ray variability amplitude of AGN: determination of the `two-folding' time-scale (i.e. the time-scale for the emitted flux to change by a factor of two), calculation of the PSD amplitude at a given frequency, and estimation of the so called `normalized excess variance' (\\rms, i.e.  the variance of a light curve normalized by its mean squared after correcting for the experimental noise). In all cases, these quantities appear to anti-correlate with the source luminosity (Barr \\& Mushotzky 1986; Lawrence \\& Papadakis 1993; Green, McHardy \\& Lehto 1993; Nandra \\etal\\ 1997; Turner \\etal\\ 1999; Leighly 1999; Markowitz \\& Edelson 2001).  One possible explanation for the observed anti-correlation between the X-ray variability amplitude and X-ray luminosity, \\lx, is that more luminous AGN have larger black hole masses (\\bhm) as well.  In this case, as \\bhm\\ increases, so does the size of the X-ray source, and a change of the source luminosity by a constant fraction takes relatively longer; equivalently, variability amplitude measured at a fixed timescale should decrease with mass. The significant progress in the measurement of the black hole (BH) mass in the centers of nearby galaxies has made it possible to actually directly test the dependence of the X-ray variability amplitude on \\bhm\\ in AGN. The data show that \\rms\\ measured at a certain timescale is indeed anti-correlated with \\bhm\\ (Lu \\& Yu 2001; Bian \\& Zao 2003; Papadakis 2004).  If there is an intrinsic correlation between X-ray variability and \\bhm, then X-ray variability measurements could be used in order to measure the central black hole mass in these objects.  This possibility has already been studied in the past. Hayashida \\etal\\ (1998) and Czerny \\etal\\ (2001) have used the PSD normalized by the square of the mean flux as a measure of the X-ray variability of a source. By calculating the ratio of the frequencies at which the PSD$\\times frequency$ has a certain value in AGN and Cyg X-1, they were able to estimate the \\bhm\\ in AGN.  Recently, long, well-sampled \\xte\\ light curves have been used in order to accurately estimate the X-ray PSD of AGN. One of the main results from these studies has been the unambiguous detection of a characteristic `break frequency', \\nb, at which the power spectrum changes its slope from a value of $-2$ to $-1$ above and below \\nb, respectively (Uttley \\etal\\ 2002; Markowitz \\etal\\ 2003). This break frequency is analogous to the high frequency break in galactic sources and should not be confused with the low frequency break, where the spectrum changes slope from $-1$ above to $0$ below.  This high frequency break appears to correlate well with \\bhm, in the sense that $1/\\nu_{\\rm bf} \\propto {\\rm M_{BH}}$ (Markowitz \\etal\\ 2003). Thus another way to estimate \\bhm\\ is through the determination of \\nb\\ in the X-ray PSD of a source. However, this method requires rather long, high quality data.  We propose a new method to estimate \\bhm\\ in AGN, which may be used even for relatively low quality and/or short data coverage. The method is based on the use of \\rms\\ as a measure of their X-ray variability amplitude, and takes Cyg X-1 as a reference point. Using recent, archival \\xte\\ and \\asca\\ light curves of 21 AGN, we estimate their \\bhm, and we compare our results with \\bhm\\ estimates which are based on other methods.  We find that the accuracy of the \\bhm\\ estimates using our new method is roughly similar to the accuracy of the estimates with well established methods like reverberation mapping. Furthermore, the \\rms\\ estimation of an X-ray light curve is rather easy and straightforward. Consequently, as long as typical AGN observations performed by present day satellites (like {\\it XMM-Newton} and {\\it CHANDRA}) detect significant variations, they can be used to estimate \\rms\\ and hence \\bhm, using the method which we present in this work. It is possible then that the use of our method can provide \\bhm\\ estimates for a large number of AGN in the near future. ",
        "conclusions": "\\begin{itemize} \\item We propose a new method based on measurements of \\rms\\ of X-ray variability of AGN to \\bhm\\ estimates. The estimates of black hole masses in AGN obtained from our method (\\bhms) are consistent with the estimates based on other techniques such as the reverberation mapping technique or the \\bhm-stellar velocity dispersion relation. \\item The constant $C$ in Eq.~(\\ref{bhm}) is equal to 0.96 for `classical' AGN and 19.2 (i.e. 20 times larger) for NLS1 \\item We suggest that IC5063, classified as S2 galaxy, is in fact a hidden NLS1. \\end{itemize}"
    },
    "0403/astro-ph0403581.txt": {
        "abstract": "{ We have performed 23 GHz VLA observations of 7 compact, luminous infrared galaxies, selected to have evidence of starburst activity. New and published multi--frequency data are combined to obtain the spectral energy distributions of all 7 galaxies from the near--infrared to the radio (at 1.4 GHz). These SEDs are compared with new models, for dust enshrouded galaxies, which account for both starburst and AGN components.  In all 7 galaxies the starburst provides the dominant contribution to the infrared luminosity; in 4 sources no contribution from an AGN is required. Although AGN may contribute up to 50 percent of the total far--infrared emission, the starbursts always dominate in the radio. The SEDs of most of our sources are best fit with a very high optical depth of $\\sim 50$ at $1\\;\\rm \\mu m$.  The scatter in the far--infrared/radio correlation, found among luminous IRAS sources, is due mainly to the different evolutionary status of their starburst components. The short time--scale of the star formation process amplifies the delay between the far--infrared and radio emission. This becomes more evident at low radio frequencies (below about 1 GHz) where synchrotron radiation is the dominant process. In the far--infrared (at wavelengths shorter than 100 $\\mu$m) an additional source of scatter is provided by AGN, where present. AGN may be detected in the near--infrared by the absence of the knee, typical of stellar photospheres. However, near--infrared data alone cannot constrain the level at which AGN contribute because the interpretation of their observed properties, in this wave--band, depends strongly on model parameters. ",
        "introduction": "One of the most intriguing aspects of IRAS galaxies (Soifer et al., 1984, Sanders \\& Mirabel 1996), characterized by their extremely cool emission (peaking between 60 and $100\\; \\rm \\mu m$), is the origin of their high luminosities, often exceeding $10^{12}\\; \\rm L_\\odot$. In such galaxies it is clear that we observe radiation from dust at temperatures of a few 10s of Kelvin. However, the question of whether the heat source is baryonic nuclear processing in young stars, or accretion of baryons onto a central black hole is still debated. Recent evidence (e.g. Ferrarese 2002) that all spheroids harbour a central black hole, suggested by the existence of a tight correlation between the masses of black holes and the spheroidal components, implies that there may be a relationship between these two processes. The evidence that almost all ultra luminous IRAS sources (ULIRGs) exhibit signatures of on-going or recent interaction, suggests that the triggering mechanism is dynamical interaction.  Several questions remain: \\begin{enumerate} \\item What are the relative fractions of mass going into star formation and accretion respectively? \\item How tight is the link between star formation and nuclear activity? And is there any causal link between the two processes? \\item How important is the role of mutual feedback in the termination of these processes? \\end{enumerate} The answers to these questions %are even more important take on greater importance if, as deduced by deep ISO and sub--millimetric surveys (e.g. Franceschini et al., 2001, 2003; Elbaz et al., 2002; Smail et al., 2000; Barger et al., 2000), ULIRGs are the local counterparts of a larger population of luminous galaxies at redshifts z$\\geq 0.5$.  Recent spectroscopic determinations of the redshift of SCUBA sources (Smail et al., 2003) confirm that these objects have extreme luminosities. If these luminosities were due to starburst episodes (accounting for possible magnification by gravitational lensing), they would correspond to star formation rates (SFRs) in excess of several thousand M$_\\odot$/yr. If these extreme SFRs were maintained over a dynamical timescale, a spheroidal component of a typical L$_*$ galaxy could be assembled.  Both star formation and nuclear activity can be significant sources of energy feedback into the interstellar and intergalactic media -- galactic winds and re-ionization respectively (e.g. Granato et al., 2001; Granato et al., 2003). The latter mechanism is one of the most important, if least understood, ingredients in current models of galaxy formation (Cavaliere, Lapi \\& Menci, 2002).  Bressan, Silva \\& Granato (2002) showed that comparison of the radio and far--infrared (FIR) emission may provide an accurate diagnostic for the study of the star formation history during the strong starburst phase characteristic of compact ULIRGs.  By means of new starburst models they were able to show that deviations from the FIR/radio correlation could be used to derive the evolutionary status of the starburst.  These deviations are expected in bursts of short duration because at early times (a few $10^6\\;\\rm y$) even the most massive stars formed in the burst will not have ended their lives as supernovae and no excess synchrotron emission should result from the starburst activity. Such starbursts would be relatively strong FIR emitters. Another important signature of a young starburst is its radio slope; as young starbursts are not yet dominated by synchrotron radiation from supernovae their radio slope  should be flatter than $S_{\\nu}\\propto \\nu^{-0.7}$, due to the large free-free component from ionized gas.  Bressan, Silva \\& Granato's (2002) ability to derive the radio slope of objects in Condon's (1991) sample of compact ULIRGs was limited by their use of existing radiometry at $1.4\\;\\rm GHz$ (which may be affected by free-free absorption). Moreover, they did not attempt to perform a complete fit to the complete SED of the galaxies, nor did they  account for the effects of the AGN in the infrared.  The new $23\\; \\rm GHz$ fluxes presented here allow for both the determination of the radio spectrum, unaffected by free-free absorption, and place a useful constraint on the level of free-free emission. The combination of radio data with IRAS fluxes at 60 and $100\\;\\rm \\mu m$ provides a strong constraint on the age of starburst activity. FIR data below $60\\;\\rm \\mu m$ are sensitive to the contribution of AGN, which may then be estimated as the excess above the starburst component.  By studying local ULIRGs we can probe all of the important physical mechanisms involved in the formation of spheroidal systems: dynamical interaction; the triggering of star formation; QSO activation and, perhaps most important, feedback. In this paper we present new $23\\;\\rm GHz$ observations of 7 compact ULIRGs. With these new data and existing archival photometry/radiometry, we study their star formation history and possible AGN contributions.   Source selection is briefly discussed in Section \\ref{sample}, corresponding observations are presented in Section \\ref{observations}. In Section \\ref{datared} we describe the procedure used to obtain the flux measurements of the observed sources while  in Section \\ref{results} we present the data.  In Section \\ref{sed} we present starburst and AGN models adopted to fit the SEDs of the observed galaxies. Discussion of the results and  conclusions are given in Sections \\ref{discussion} and \\ref{conclusions}, respectively.    \\begin{table*} \\caption{Previously published radio, FIR, mid-IR and sub--mm fluxes of the compact luminous infrared galaxy sample.} \\begin{tabular}{lccccrrrrcc} \\hline \\hline Source   & $\\rm{S}_{1.4}$ &  $\\rm{S}_{4.85}$ &$\\rm{S}_{8.44}$ &$q$&$\\rm{S}_{12\\mu}$&$\\rm{S}_{25\\mu}$&$\\rm{S}_{60\\mu}$&$\\rm{S}_{100\\mu}$ & $\\rm{S}_{850\\mu}$ & log L$_{\\rm FIR}$ \\\\ Name     & / mJy      & / mJy        &  / mJy         & -  & /mJy & /mJy & /mJy & /mJy & /mJy & L$_{\\odot}$ \\\\ & (1)        & (2)              &  (3)       & (4)  & (5) & (6) & (7) & (8) & (9) & (10)  \\\\ \\hline NGC0034        &$67.5\\pm2.5$   & \\ldots  & $15.2\\pm0.8$                &2.52 &360    &2380 &16080 &16970 &\\ldots    & 11.54\\\\ CGCG436-30     &$50.2\\pm1.6$   & 21.5    & $12.7\\pm0.6$                &2.46 &270    &1410 &10720 &9600  &$39 \\pm 8$& 11.77\\\\ IRAS01364-1042 &$15.8\\pm0.7$   & \\ldots  & $8.2\\pm0.4$                 &2.67 &$<150$ &430  &6530  &7000  &\\ldots    & 11.93\\\\ IRAS08572+3915 &$4.89\\pm0.19^*$& \\ldots  & $4.1\\pm0.2$                 &3.11 &340    &1840 &7660  &5060  &$17\\pm7$  & 12.24\\\\ NGC7469        & $180.5\\pm5.4$ & $70\\pm8$&$>15.0\\pm0.8^\\dagger$        &2.29 &1488   &5430 &25740 &32466 &$264\\pm30$& 11.68\\\\ IC5298         & $34.7\\pm1.4$  & 13.9    & $8.2\\pm0.4$                 &2.65 &288    &1690 &7880  &10480 &$77\\pm15$ & 11.61\\\\ UGC12812       & $70.7\\pm2.2$  & 31.1    &$21.5\\pm1.1^{\\dagger\\dagger}$&2.51 &492    &2470 &16710 &20130 &$132\\pm25$& 11.54\\\\ \\hline \\end{tabular} {\\footnotesize (1):  NVSS 1.4 GHz flux densities Condon et al. (1998); $^*$ FIRST 1.4 GHz flux densities Becker et al. (1995);  (2): Sopp \\& Alexander (1992);  (3): Condon et al.'s (1991) total 8.44 GHz source flux densities, $^\\dag$ main component flux density, $^{\\dag\\dag}$ indicates a relatively uncertain flux density due to surrounding diffuse emission. Errors in the 8.44 GHz fluxes were estimated from the observational parameters given in Condon et al. using the method described in Section \\ref{radiometrysection}  (4): logarithmic far infrared to 1.4 GHz luminosity ratio as defined by Condon et al. (1991);  (5), (6), (7) \\& (8): IRAS-BGS flux densities (Soifer et al., 1989); (9): SCUBA $850\\;\\rm \\mu m$ flux densities Dunne et al. (2000); (10) far infrared luminosity in the range $8-1000\\;\\rm \\mu m$. } \\label{sourceproperties} \\end{table*} ",
        "conclusions": "We have performed new high frequency VLA radio measurements ($23\\; \\rm GHz$)   of 7  compact ULIRGs chosen from Condon et al.'s (1991) sample.  We have combined the new observational data with existing data  at other radio frequencies (1.4 GHz, 5.4 GHz, 8.4 GHz); in the millimetre; from SCUBA $850\\;\\rm \\mu m$; infrared IRAS 100, 60 25, $12\\;\\rm \\mu m$ bands and in the NIR with J, H and K band  observations from 2MASS.  We fit the global SEDs  by means of a combination of new AGN and starburst models. We find:  \\begin{enumerate}  \\item The high frequency radio data have proven essential for the determination of the relative contribution of the star-burst and the AGN. Although AGN can be detected in several ways (such as via their hard X-ray emission, near-infrared colours or optical line emission), quantifying the energy contribution of the AGN to the bolometric luminosity of the source requires an evaluation of the starburst power. On one hand, the new radio observations at $23\\; \\rm GHz$ allow a clean definition (unaffected by free-free absorption) of the shape of radio emission over a wide range of frequencies. On the other, the flux at $23\\; \\rm GHz$ is mostly contributed ($>$50 percent) by free-free emission from ionized HII regions and is therefore directly related to ongoing star formation. The flux at lower frequencies is instead sensitive to the supernova rate and thus carries information about the previous history of the starburst.  \\item We have found that radio data, together with the $100\\;\\rm \\mu m$ flux define almost unequivocally the recent star formation history of the burst. This confirms the suggestion by Bressan, Silva \\& Granato 2002, that combining FIR and Radio data provides a fine resolution age determination, for these obscured sources.  \\item  While the starburst contribution is generally well constrained by the SED longward of $60\\;\\rm \\mu m$, the SED below $60\\;\\rm \\mu m$ is sensitive to the contribution from the hot dust of the AGN. Thus, the excess emission above the starburst contribution at 12, 25 and $60\\;\\rm \\mu m$ may be used to quantify the contribution from the AGN. A degree of uncertainty  remains regarding the strength of the AGN.  Low optical depth in molecular clouds  favours warm emission from the starburst, thus a smaller contibution from AGN is required to fit the mid-infrared data. However, as the emission from the starburst gets warmer, its contribution at $100\\;\\rm \\mu m$ decreases and to fit this data point a larger SFR is required; eventually becoming incompatible with the radio data. Thus, while still affected by uncertainties, the method we propose proves to be useful in providing an estimate of the relative contribution of AGN and starburst.  \\item The effect of the AGN is also visible from the total J, H and K fluxes but its quantification is impossible from these NIR colours alone because of the strong dependence of the AGN emission on the model parameters at these wavelengths.  \\item It is possible to fit the SEDs of most of the sources in our sample with no contribution from an AGN. However, better fits, especially for the near infrared colours, are obtained when the bolometric luminosity has a small AGN contribution. We note that Farrah et al. (2003), who fitted the SEDs of ULIRGs with template spectra, find that in general the starburst contributions dominante.  \\item To fit the mid-infrared data a very high optical depth at 1 micron is generally required. In the most extreme case, IRAS01364-1042, the SED cannot be fit unless $\\tau_{1\\;\\rm \\mu m}\\sim 180$. Although we cannot do an object by object comparison, values as high as at least 50 are consistent with gas column densities measured in the centre of ULIRGs. The fact that extinction measurements at shorter wavelengths give much lower values probably indicates that such observations penetrate only a surface layer of the ISM and do not probe the whole sight-line towards the central energy source.   \\end{enumerate}"
    },
    "0403/astro-ph0403676_arXiv.txt": {
        "abstract": "The primary objective of the study of unidentified EGRET sources with INTEGRAL is to locate with a few arcminute accuracy the hard X-ray / soft $\\gamma$-ray counterparts within the EGRET error circle or to determine upper limits on their emission in the INTEGRAL energy range. The INTEGRAL Galactic Plane Scan (GPS) and Galactic Centre Deep Exposure (GCDE) data have been analysed and cross--correlated with the EGRET $3^{\\rm rd}$ Catalogue to search for counterparts of EGRET sources. The IBIS detection of a source within the EGRET error circle of 3EG J$1736-2908$ is presented and its possible identification with the active galaxy GRS $1734-292$ is discussed. Finally, preliminary results on the cross-correlation between EGRET unidentified sources and the first data from the IBIS survey of the Galactic Centre are presented. ",
        "introduction": "The $3^{\\rm rd}$ Energetic Gamma--Ray Experiment Telescope (EGRET) Catalogue of high--energy (E$>$100 MeV) $\\gamma$--ray sources (Hartman et al. 1999)  refers to the observations carried out by the EGRET experiment onboard the Compton Gamma--Ray Observatory (CGRO) between April 1991 and October 1995. It contains 271 point sources, including: 1 solar flare, the Large Magellanic Cloud, 5 pulsars, the radio galaxy Centaurus A, and 66 high--confidence identifications of blazars. In addition, 27 lower--confidence blazar identifications are included in the catalog, which leaves 170 sources unidentified.  Since the publication of the catalog, a lot of effort has been dedicated to the identification of the 170 unidentified EGRET sources (Grenier 2003). While most of the extragalactic counterparts are blazars, two possible identifications of EGRET sources with radiogalaxies have been recently proposed (Combi et al. 2003, Mukherjee et al. 2003), in addition to the well known radiogalaxy Cen A. However, the multiwavelength counterpart search is hampered by the limited EGRET point spread function, which gives error radii of $0.5^\\circ-1^\\circ$.  One of the two main telescopes on board INTEGRAL (Winkler et al. 2003) is the IBIS imager (Ubertini et al. 2003).  IBIS combines for the first time broad band high energy coverage ($15$~keV -- $10$~MeV) with imaging capability.  IBIS has a large field of view ($19^\\circ\\times 19^\\circ$ at half response), and a fine angular resolution of $12'$, sampled in $5'$ pixels in the low energy ($0.015-1$~MeV) layer ISGRI (Lebrun et al. 2003), and in $10'$ pixels in the high energy ($0.175-10$~MeV) detector PICsIT (Di Cocco et al. 2003). ISGRI reaches a point source location accuracy better than $1'$ for a $30\\sigma$ detection (Gros et al. 2003). These unprecedented characteristics for a $\\gamma$--ray telescope clearly make IBIS one of the best instruments available to date to search for counterparts of unidentified EGRET sources. Within the framework of the Core Programme (CP) observations, a project has been set up to systematically search for associations between INTEGRAL sources and EGRET unidentified error sources. A first possible identification of an EGRET source has already been found (see Di Cocco et al. 2004, A\\&A submitted, for details). ",
        "conclusions": "We presented the first results of the early activity in the search for counterparts of the unidentified EGRET sources in the framework of the INTEGRAL Core Programme. One candidate counterpart has been found close to the Galactic Centre and the identification with the active galaxy GRS$1734-292$ has been proposed. The cross--correlation of the 2~Ms mosaic of the central radiant (Revnivstev et al. 2004) with the $3^{\\rm rd}$ EGRET catalog has not produced other possible counterparts. The research should now be extended to the others regions of the sky covered by the INTEGRAL core programme."
    },
    "0403/astro-ph0403395_arXiv.txt": {
        "abstract": "Effective temperatures and radii for 92 planet-hosting stars as determined from the InfraRed Flux Method (IRFM) are presented and compared with those given by other authors using different approaches. The IRFM temperatures we have derived are systematically lower than those determined from the spectroscopic condition of excitation equilibrium, the mean difference being as large as 110~K. They are, however, consistent with previous IRFM studies and with the colors derived from Kurucz and MARCS model atmospheres. Comparison with direct measurements of stellar diameters for 7 dwarf stars, which approximately cover the range of temperatures of the planet-hosting stars, suggest that the IRFM radii and temperatures are reliable in an absolute scale. A better understanding of the fundamental properties of the stars with planets will be achieved once this discrepancy between the IRFM and the spectroscopic temperature scales is resolved. ",
        "introduction": "Accurate determination of the fundamental stellar parameters for the planetary host stars is required to improve our knowledge of the physical properties of the extrasolar planets. The temperatures and luminosities of these stars are used to obtain their masses and radii, which are the stellar parameters required to determine the masses and radii of the planets from radial velocity and transit observations, respectively.  The temperatures of the planetary host stars have been spectroscopically determined from the excitation equilibrium of Fe I lines (e.g. Heiter \\& Luck 2003; Gonzalez et al. 2001; Laws et al. 2003; Santos, Israelian, \\& Mayor 2004). The most complete study of this kind is that of Santos et~al. (2004, SIM04), who determined temperatures of 98 of them. Using a different approach, Ribas et~al. (2003, R03) determined the temperatures of these stars from infrared photometry. Both results are similar and, in fact, consistent with most of the previous studies, with the exception of the Fe~I temperatures of Sadakane et~al. (2002) and Takeda et~al. (2002), which are about 100~K lower.  To date, it has not been verified whether the spectroscopic temperatures are consistent with those determined from stellar angular diameters. This is basically due to the lack of dwarf stars with measured angular diameters. However, recent interferometric measurements (Kervella et~al. 2003, 2004; Lane, Boden, \\& Kulkarni 2001; Pijpers et~al. 2003) and transit observations (Brown et~al. 2001) allow one to test this now. Here we show that our IRFM temperatures are consistent with direct measurements of angular diameters. Then we apply the IRFM to obtain the temperatures of 92 of the known planet-hosting stars and compare them with those obtained by other methods. ",
        "conclusions": "The temperatures obtained in this work for the stars with planets differ from those given by most of the other groups by about 110~K. In particular, there is a discrepancy with the excitation equilibrium temperatures of Santos et al. (2004) % and the method of infrared photometry of Ribas et~al. (2003).  On the other hand, the $\\teff$ vs $(V-K)$ relation from model atmospheres is closer to our $\\teff$ scale and the angular diameters measured by interferometry and transit observations for 7 dwarf stars are well scaled to those derived from the IRFM. The bolometric flux calibration adopted (AAM95) seems to be in excellent agreement with the measurements reported in the literature ($\\pm1$\\%), which strengthen the present results. The comparison with direct measurements suggest that the zero point of the IRFM $\\teff$ scale is within 50~K, considering twice the standard error. New measurements of angular diameters for stars in the range 4000~K$<\\teff<$6000~K are encouraged to better define it.  The 110~K difference makes the stars bigger according to the IRFM, which implies that the planetary radii obtained from transit observations would be larger if our $\\teff$ scale is adopted instead of the spectroscopic one.  %  Metallicity determinations for the stars with planets taking into account the IRFM $\\teff$ scale would be interesting. Note, however, that large samples of stars both with and without detected giant planets have been analyzed in a homogeneous way, so it is unlikely that this will affect the well stated metallicity enhancement of the planet-hosting stars, although small differences may be found in an absolute scale."
    },
    "0403/astro-ph0403440_arXiv.txt": {
        "abstract": "Cold fronts -- sharp discontinuities recently discovered by Chandra in many clusters of galaxies -- are believed to be due to a hot gas flow over a colder gravitationally bound gas cloud. We analyze the stability of the fronts with respect to Kelvin-Helmholtz instability and show that an intrinsic width of the interface of the order of a few per cent of the curvature radius strongly limits the growth of perturbation. For the best studied case of a front in the Cluster Abell 3667 we conclude that current observational data on the width and extent of the front can be explained even in the absence of dynamically important magnetic fields. ",
        "introduction": "Cold fronts were discovered as sharp features in the X-ray surface brightness distribution in Chandra observations of the clusters A2142 and A3667 (Markevitch et al., 2000, Vikhlinin, Markevitch \\& Murray, 2001a), see also Markevitch et al. (2002). Similar features have now been found in several other clusters (e.g. Sun et al., 2002, Kempner, Sarazin \\& Ricker, 2002). Unlike shocks, these features have lower gas temperature on the X-ray brighter side of the discontinuity. For that reason they are called ``cold fronts''. It is believed that some cold fronts are formed when a subcluster merges with another cluster and the ram pressure of gas flowing outside the subcluster gives the contact discontinuity the characteristic curved shape. Indeed, features resembling cold fronts are found in numerical simulations of cluster formation (Bialek, Evrard \\& Mohr, 2002, Nagai \\& Kravtsov, 2003). Ablation of the gaseous cloud by the hot gas causes characteristic differential motion of the gas inside the subcluster, which transports the low entropy gas from the subcluster core towards the contact discontinuity, thus enhancing the jump in temperature and surface brightness across the discontinuity (Heinz et al. 2003).  Here we address the question of the front stability. As was pointed out by Vikhlinin et al. (2001a,b), Vikhlinin and Markevitch (2002) the observed fronts are narrow and could be unstable to the Kelvin-Helmholtz (KH) instability. For A3667 however the front appears to be narrow (less than $\\sim$5 kpc in width) up to $\\sim$30\\deg from the stagnation point.  Magnetic fields can act as a stabilizing agent thus allowing indirect estimates of the field strength near the front (Vikhlinin, Markevitch \\& Murray,2001b). On the other hand numerical simulations without magnetic field (and in particular relatively high resolution simulations by Heinz et al. 2003) do not show instability of the front within 20-30\\deg. While this discrepancy could be due to numerical effects, we revisit the question of the front stability below and show that the characteristic convex geometry and finite (small) intrinsic width of the interface may help to stabilize the discontinuity with respect to KH instability.   The structure of the paper is as follows. In Section 2 we derive a simple expression for the growth of the KH instability along the curved interface. In Section 3 we argue that diffusion processes are likely to set an approximately constant width for the interface. In Section 4 we discuss astrophysical applications. The last section summarizes our findings. ",
        "conclusions": "We have shown that accounting for wave number changes along the curved interface and finite (but small) width of the interface significantly reduces the KH instability growth factor. We also argue that if diffusion sets the intrinsic width of the interface then this width does not vary much along the interface (at least for small angles).  For a set of parameters relevant for the front in A3667 the growth factor is not large enough to guarantee nonlinear growth at angles $\\sim$30\\deg, provided that the intrinsic width of the front is of the order of a few percents of the curvature radius (i.e. not much less than the existing observational limit). Therefore dynamically important magnetic fields may not be necessary to stabilize this particular cold front. The growth factor (and therefore possible limits on the magnetic fields) {\\bf depends critically on the width of the interface} and much more weakly on the angular extent over which the front remains sharp. A factor of 2-3 improvement in the observational constraints on the interface width is therefore extremely important.  For a given width of the interface the value of the growth factor does not depend strongly on the absolute value of the shear velocity. This probably explains the ubiquity of fronts both in observations and hydrodynamical simulations."
    },
    "0403/astro-ph0403489_arXiv.txt": {
        "abstract": "s{I review attempts made to determine the properties of neutron stars.  I focus on constraints on the maximum mass that a neutron star can have, and on attempts to measure neutron-star radii. So far, there appears to be only one neutron star for which there is strong evidence that its mass is above the canonical $1.4\\,M_\\odot$, viz., Vela~X-1, for which a mass close to $1.9\\,M_\\odot$ is found. Prospects for progress appear brightest for studies of systems in which the neutron star should have accreted substantial amounts of matter.  While for individual systems the evidence that neutron stars can have high masses is weak, the ensemble appears to show that masses around $1.6\\,M_\\odot$ are possible.  For the radius determination, most attempts have focussed on neutron stars in low-mass X-ray binaries in which accretion has temporarily shut down.  These neutron stars are easiest to model, since they should have pure Hydrogen atmospheres and low magnetic fields.  To obtain accurate radii, however, requires precise distances and very high quality data. } ",
        "introduction": " ",
        "conclusions": "The pessimistic conclusion from the above would be that, since their discovery, neutron stars have not much advanced our understanding of the physics at extreme densities.  Viewing the situation more positively, though, one sees that the mass determinations are now getting accurate enough to be interesting, especially for the radio pulsars with white dwarf companions.  The precision will increase with further timing of the pulsars, helped by optical studies of the white dwarfs.  Furthermore, new avenues are explored in which the thermal emission is modelled and used to derive angular diameters and gravitational redshifts.  The study of low-mass X-ray binaries in quiescence seems particularly promising, particularly once the distances settle down. For globular clusters, this is in progress, while for individual systems in the field it will come in the somewhat longer run, using direct parallax determinations with NASA's Space Interferometry Mission or ESA's GAIA, both expected to be launched in about ten years.  Finally, not mentioned in my talk or this write-up, there are a number of nearby neutron stars with what appear to be purely thermal spectra. Recently, for a number of these, absorption features have been detected, which are likely due to Hydrogen atmospheres in extremely strong magnetic fields, $10^{13}$ to $10^{14}\\,$G.\\cite{vkkd+04} In such fields, absorption might be due to proton cyclotron or neutral Hydrogen; with more than a single line, one may be able to solve for the field strength and the gravitational redshift.  Several are close enough for parallax determinations; combined with angular diameters from atmospheric modelling, this might yield good radii.  Overall, the outlook seems fairly bright.  \\def\\aap,#1,{{\\it Astron.\\ Astroph.{}} {\\bf #1},} \\def\\aaps,#1,{{\\it Astron.\\ Astroph.\\ Supp.{}} {\\bf #1},} \\def\\apj,#1,{{\\it Astroph.\\ J.{}} {\\bf #1},} \\let\\apjl=\\apj \\def\\araa,#1,{{\\it Ann.\\ Rev.\\ Astron.\\ Astroph.{}} {\\bf #1},} \\def\\mnras,#1,{{\\it Mon.\\ Not.\\ R.\\ Astron.\\ Soc.{}} {\\bf #1},} \\def\\nat,#1,{{\\it Nature} {\\bf #1},}"
    },
    "0403/astro-ph0403506_arXiv.txt": {
        "abstract": "I present an Hipparcos color-magnitude diagram (CMD) that is color-coded by transverse velocity $v_\\perp$.  This illustrates the connection between the photometric and kinematic properties of various stellar populations in a manner that is particularly suitable to introductory astronomy courses. ",
        "introduction": "}  Among the many wonderful uses of the Hipparcos catalog \\citep{hip} is that the Hipparcos color-magnitude diagram (CMD) directly demonstrates for students the connections between local stars and the more homogeneous samples found in clusters.  However, since the local population is a mixture of populations of various ages and metallicities, the Hipparcos CMD inevitably appears less distinct than cluster CMDs.  In fact, the Hipparcos catalog contains substantial auxiliary information to distinguish among these populations in the form of proper motions (and hence transverse velocities $v_\\perp$). Early main-sequence stars, being young, tend to be moving more slowly than later main-sequence stars, which are predominantly older.  Halo subdwarfs, which lie below the main-sequence due to their lower metallicities, also tend to travel quite fast relative to the Sun. Thick disk stars are intermediate in both metallicity and kinematics. ",
        "conclusions": ""
    },
    "0403/astro-ph0403026_arXiv.txt": {
        "abstract": "{ We present an X-ray study of the galaxy group RGH~80, observed by \\xmm. The X-ray emission of the gas is detected out to $\\sim 462h^{-1}_{50}$ kpc, corresponding to $\\sim 0.45 r_{200}$. The group is relatively gas rich and luminous with respect to its temperature of $1.01\\pm 0.01$ keV. Using the deprojected spectral analysis, we find that the temperature peaks at $\\sim 1.3$ keV around $0.11r_{200}$, and then decreases inwards to 0.83 keV at the center and outwards to $\\sim 70\\%$ of the peak value at large radii. Within the central $\\sim 60$ kpc of the group where the gas cooling time is less than the Hubble time, two-temperature model with temperatures of 0.82 and 1.51 keV and the Galactic absorption gives the best fit of the spectra, with $\\sim 20\\%$ volume occupied by the cool component. We also derive the gas entropy distribution, which is consistent with the prediction of cooling and/or internal heating models. Furthermore, the abundances of O, Mg, Si, S, and Fe decrease monotonically with radius. With the observed abundance ratio pattern, we estimate that $\\sim 85\\%$ or $\\sim 72\\%$ of the iron mass is contributed by SN Ia, depending on the adopted SN II models. ",
        "introduction": "In the past decade X-ray imaging spectroscopy has become a powerful probe of the physical conditions of the hot ($T^{>}_{\\sim}10^{7}$ K) intergalactic medium (IGM) of galaxy groups. Spatial distributions of gas temperature, density and metallicity, as well as the structures of the gravitational potential and dark matter halos have been measured for numerous galaxy groups. This improves significantly our knowledge of dynamical properties and formation of galaxy groups. Among these, the discovery of similarity breaking of some IGM properties from clusters to groups is crucial for the study of cosmic evolution of the hot IGM in very massive halos of the universe. The well-known example is the deviation of the X-ray luminosity ($L_{\\rm X}$) and temperature ($T$) relation of groups and clusters from the prediction of self-similar model (e.g. Edge \\& Stewart 1991; David \\etal 1993; Wu, Xue \\& Fang 1999; Helsdon \\& Ponman 2000; Xue \\& Wu 2000 and references therein). The break lends strong support to the argument that the IGM of groups was much more affected by non-gravitational processes, such as radiative cooling and/or feedback of star formation, than the IGM of clusters. Another example is the detection of entropy excess in clusters and groups (Ponmam, Cannon \\& Navarro 1999; Lloyd-Davies, Ponman \\& cannon 2000; Xu, Jin \\& Wu 2001), which reinforces the presence of strong non-gravitational effects especially in the inner regions of poor clusters and groups. The reliable $L_{\\rm X}-T$ relations and the entropy profiles obtained from spatially resolved X-ray imaging spectroscopy may allow one for the first time to distinguish between the competing non-gravitational models (Voit \\etal 2002, 2003; Mushotzky \\etal 2003).  X-ray spectra of groups are dominated by emission line features. The wealth of characteristic emission lines provides a robust tool for understanding the metal enrichment processes in the IGM and possibly for constraining the supernova history of galaxy groups (e.g. Finoguenov \\& Ponman 1999). However, an accurate determination of the metal abundances depends strongly on precise measurements of the temperature structure of the hot IGM in groups, as the dominated Fe L-shell lines are highly temperature sensitive (e.g. Buote 2002; Buote \\etal 2003a). The high spectral resolution of \\xmm\\ allows us to obtain the abundances of O and Mg in addition to Si, S and Fe, and provides us a deep insight into the chemical evolution history of groups with the abundance ratio pattern and the metal mass to light ratios (e.g. Xu \\etal 2002; Matsushita, Finoguenov \\& B\\\"ohringer 2003).  In this paper, we report the results from new \\xmm\\ observations of RGH~80, which was first identified by Ramella, Geller \\& Huchra (1989) in the center for astrophysics redshift survey. It was identified as an extended X-ray source in the {\\it ROSAT} All-Sky Survey with an X-ray luminosity of $3.33\\times 10^{43}h^{-2}_{50}$ erg s$^{-1}$ in the 0.1--2.4 keV band within an aperture of $1.0h^{-1}_{50}$ Mpc (NRGs241; Mahdavi \\etal 2000). Global fitting of the {\\it ASCA} SIS and GIS data by Davis, Mulchaey \\& Mushotzky (1999) yielded a mean gas temperature of $1.02\\pm 0.05$ keV, an average abundance of $0.26^{+0.16}_{-0.08}Z_{\\odot}$ for the $\\alpha$ elements and $0.20^{+0.05}_{-0.06}Z_{\\odot}$ for iron. However, Buote (2000) pointed out that a two-temperature spectral model provides a better fit to the {\\it ASCA} spectra of RGH~80 and has a metallicity that is substantially higher than that obtained by the single temperature spectral model, $Z=0.67^{+0.42}_{-0.24}Z_{\\odot}$. We will use the \\xmm\\ data to examine the spatially resolved X-ray properties of the group and explore their physical applications.  In section 2, we describe the observations and data reduction procedure. We present the detailed spectral analysis in section 3. In Section 4 and 5, we derive the gas density profiles, and calculate the three dimensional distributions of gas and dark matter, respectively. We discuss and summarize the implications of our results in section 6 and 7. Throughout this paper, we take $H_0 = 50$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm M} =1$ and $\\Omega_{\\Lambda}=0$. At the group redshift of 0.037, $1\\arcmin$ corresponds to 60.2 kpc. Unless stated otherwise, we adopt $68\\%$ confidence limit for error analysis. ",
        "conclusions": "\\label{sec:conc} We summarize below the main conclusions from our analyses of the \\xmm\\ observations of the galaxy group of RGH~80.  \\begin{itemize} \\item The X-ray emission of the group is detected out to $\\sim 7.67\\arcmin$, or $462h^{-1}_{50}$ kpc, corresponding to $0.45r_{200}$. The group seems to be relatively gas rich and luminous with respect to its temperature ($1.01\\pm 0.01$ keV).  \\item Spectral analysis shows that the temperature profile, which increases from the center and then declines with radius after reaching a plateau around $0.11r_{200}$, follows a universal profile (Mulchaey 2000) and that the abundance profile of each measured element decreases monotonically with radius.  \\item In the central region, the X-ray emission of the gas is better modeled by a two-temperature spectral model, with temperatures of 0.82 and 1.51 keV and the Galactic absorption, than by a single temperature or cooling flow models. Beyond $\\sim 60$ kpc ($1\\arcmin$), the volume filling factor of the cool component, $f(R)$, becomes smaller than $\\sim 0.1$, indicating that the cool gas occupies only a small fraction of the volume therein.  \\item The gas entropy distribution derived from single-temperature assumption deviates from the prediction of the over the whole observed region, The isentropic core at the center expected from preheating model does not show up. Nevertheless, our derived entropy profile resembles what is predicted by radiative cooling model or internal heating model or cooling plus heating model.  \\item Both the abundance ratios of Fe/O and Fe/Si are very high showing a large SN Ia dominance, which is higher than that for M87 (Matsushita \\etal 2004) and comparable to or slightly higher than that for Centaurus cluster (Matsushita \\etal 2004). With the abundance ratio pattern of the hot gas, we estimate that $\\sim 85\\%$ (I99) or $\\sim 72\\%$ (W95;B;Z$_{\\odot}$) of the iron mass is contributed by SN Ia. This SN type fraction remains almost constant against radius. We find that S is significantly overproduced by the two delayed detonation models WDD1 and CDD1, while Mg is slightly underproduced by the four W95 models considered here.  \\end{itemize}"
    },
    "0403/astro-ph0403210_arXiv.txt": {
        "abstract": "We investigate a simple model for globular cluster (GC) formation. We simulate the violent relaxation of initially homogeneous isothermal stellar spheres and show that it leads to the formation of clusters with radial density profiles that match the observed profiles of GCs. The best match is achieved for dynamically unevolved clusters.  In this model, all the observed correlations between global GC parameters are accurately reproduced if one assumes that all the clusters initially had the same value of the stellar density and the velocity dispersion. This suggests that the gas which formed GCs had the same values of density and temperature throughout the universe. ",
        "introduction": "Globular clusters (GCs) are massive, dense clusters of stars, and in many galaxies they are the oldest datable objects. They contain important clues to early star formation in the universe, and to the formation history of our Milky Way and other galaxies. The dynamical evolution of GCs has been studied for many decades, and is considered to be well-understood. However, the formation of GCs is still an open question.  The internal structure and dynamics of Galactic GCs are well described by a family of lowered isothermal models (so-called King models, \\citealt{K66}). The radial profiles of King models have a flat core (characterized by a core radius $r_0$) and a sharp cut-off at a tidal radius $r_{\\rm t}$.  GCs are presently in dynamical equilibrium. Initially they could have formed in some non-equilibrium configuration which relaxed to a King model through violent relaxation \\citep{MH97} followed by a slow evolution in the galactic tidal field.  Currently no models of GC formation explain all the observed properties of these star clusters \\citep{DM94,M00}. Correlations between parameters which are not significantly affected by the dynamical evolution of a cluster (such as half-mass radius, central velocity dispersion and binding energy) are thought to reflect the initial conditions for GC formation. Therefore, any model of cluster formation should aim to address the observed correlations.  In this paper, we investigate the simplest possible model for GC formation: the dynamical evolution of a homogeneous isothermal stellar sphere. Simulations of the collapse of homogeneous stellar sphere with constant velocity dispersion were performed twenty years ago \\citep{vA82} with the intention of reproducing the $r^{1/4}$ density distribution of elliptical galaxies. Those efforts were not very successful --- the collapse produced instead a core-halo structure, somewhat like a GC. In this paper, we compare a family of such models with the observed density profiles of GCs, and attempt to use the family of such models to reproduce correlations between parameters of GCs in the Milky Way.  \\section {MODEL}  We propose a homogeneous isothermal stellar sphere as an initial non-equilibrium configuration for GCs. To test this idea, we evolve stellar spheres with a total mass $M$, an initial density $\\rho_{\\rm ini}$ and an initial velocity dispersion $\\sigma_{\\rm ini}$ for $t_{\\rm end}=10-5000$ initial crossing times $t_{\\rm cross} = (R_{\\rm ini}^3 / M)^{1/2}$, where $R_{\\rm ini} = (3 M / 4 \\pi \\rho_{\\rm ini})^{1/3}$ is the initial radius of the system.  (In this Letter we assume a system of units where the gravitational constant $G = 1$.)  We use the parallel N-body tree-code GADGET \\citep{SYW01} to run the simulations.  The stars in the cluster are represented by $N = 10^5 - 10^6$ equal mass particles. The velocities of the particles have a Maxwellian distribution.  The gravitational potential is softened with a softening length of $\\epsilon = 0.77 R_{\\rm h} N^{-1/3}$, where $R_{\\rm h}$ is the initial half-mass radius of the cluster. (Gravity is softened to minimize numerical noise due to two-body interactions between the particles.) The individual timesteps are equal to $\\sqrt{2\\eta \\epsilon /a}$, where $a$ is the acceleration of a particle, and the parameter $\\eta$ is made small enough to ensure the total energy conservation to better than 1\\%.  This configuration is the simplest realization of possible initial conditions for GCs, and yet results in remarkably good agreement with current cluster parameters.  It is convenient to express the mass of a cluster $M$ in units of the virial mass $M_{\\rm vir}: M = M_{\\rm vir} 10^\\beta$, where $\\beta$ is a mass parameter. For a homogeneous sphere $M_{\\rm vir} = \\sigma_{\\rm ini}^3 (4 \\pi \\rho_{\\rm ini} / 375)^{-1/2}$. The total energy of the system becomes positive (and the system becomes formally unbound) for $\\beta < -0.45$.  We ran a set of 17 different models (see Table~\\ref{table1}) with the same values of $\\rho_{\\rm ini}$ and $\\sigma_{\\rm ini}$, and parameter $\\beta$ ranging from --0.8 to 2. (The corresponding initial virial parameter $\\nu \\equiv 2K/W = 10^{-2\\beta/3}$ is $\\nu$ = 3.4 to 0.046, where $K$ and $W$ are initial kinetic and potential energy.)  Models with $\\beta$ = --0.8 are completely unbound throughout the simulation, whereas models with $\\beta$ = (--0.7, --0.6, --0.5) form a bound cluster, containing less than 100\\% of the total mass, after the initial expansion phase. Models with $\\beta \\geq 0$ initially collapse to a smaller half-mass radius $r_{\\rm min}$ (see Table~\\ref{table1}), and then bounce. All models with $\\beta \\geq -0.7$ experience phase mixing and/or violent relaxation, and within 10 -- 100 crossing times reach an equilibrium state with a flat core and sharply declining density at large radii. The projected surface density profiles of equilibrium clusters closely resemble the profiles of Galactic GCs. The match is the best for the GCs which have experienced little dynamical evolution (see Figure \\ref{surface density}).  Our models fit the observed surface density profiles of GCs very well, despite the lack of a tidal radius as present in King models. We do not include an external tidal field in our simulations, so this lack is expected. The tidal cutoff in surface density is observed in very few GCs \\citep{TKD95}, due to the contamination with background stars at the outskirts of the clusters.  As we will show in Section~\\ref{sec3}, all real GCs should have experienced an adiabatic collapse (when the orbital angular momentum is conserved for individual stars). To make sure that our models are in the same collapse regime as the real clusters, all our models should satisfy the following adiabaticity criterion \\citep*{alp88}:  \\begin{equation} \\sigma_{\\rm ini} \\gtrsim N^{-1/6} (M/R_{\\rm ini})^{1/2}. \\label{ALP} \\end{equation}  \\noindent In our case, this condition can be expressed as $\\beta \\lesssim (1/2)\\log N - (3/2)\\log 5$.  The adiabatic collapse criterion is then $\\beta\\lesssim 1.5$ for $N=10^5$ and $\\beta\\lesssim 2.0$ for $N=10^6$. According to this criterion, all our models undergo an adiabatic collapse (see Table~\\ref{table1}).  In our models, the collapse factor $C \\equiv R_{\\rm h}/r_{\\rm min}$ correlates well with the virial parameter $\\nu$: $C\\propto \\nu^{-\\gamma}$. The value of the exponent $\\gamma =0.95\\pm 0.02$  is very close to the adiabatic value of $\\gamma=1$ \\citep{alp88}.  From our set of models, we obtained the following correlations between the projected half-mass radius $r_{\\rm h}$, projected central velocity dispersion $\\sigma_0$, central surface density $\\Sigma_0$, central density $\\rho_{\\rm c}$, King core radius $r_0$, binding energy $E_{\\rm b}$, and mass $M$ (the uncertainties in the exponents are $\\lesssim 0.01$):  \\begin{eqnarray} r_{\\rm h}  \\sim {\\rm constant},\\quad & \\sigma_0 \\propto M^{0.57}, & \\quad\\Sigma_0 \\propto M^{1.39}, \\nonumber\\\\ \\rho_{\\rm c} \\propto  M^{1.68},\\quad & r_0 \\propto M^{-0.27}, & \\quad E_{\\rm b} \\propto M^{1.95},\\label{model}\\\\ &E_{\\rm b} \\propto \\sigma_0^{3.55}.& \\nonumber \\end{eqnarray}  \\noindent These correlations are valid for $\\beta \\geq 0$, with the last relation being valid for $\\beta \\geq -0.6$. Most of correlations become non-linear in log-log space for $\\beta < 0$ (corresponding to $\\nu > 1$).  Although for our coldest models, $r_{\\rm h}$ scales (as expected) as the initial radius (and hence as $M^{1/3}$), the total variation in $r_{\\rm h}$ over two orders of magnitude in system mass ($\\beta=0-2$) is only $\\sim 0.12$~dex (see Table~\\ref{table1}; please note that $R_{\\rm ini}\\propto M^{1/3}$).  Another very interesting correlation is $r_0\\simeq r_{\\rm min}$ within the measurement errors for $\\beta\\geq 0.2$ (Table~\\ref{table1}). The above correlation can be understood if we rewrite the theoretical result $C \\equiv R_{\\rm h}/r_{\\rm min}\\propto \\nu^{-1}$ \\citep{alp88} for the case of constant $\\rho_{\\rm ini}$ and $\\sigma_{\\rm ini}$: $r_{\\rm min}\\propto M^{-1/3}$. The theoretical relation is very close to the model relation $r_0 \\propto M^{-0.27}$, resulting in $r_0\\propto r_{\\rm min}$ as observed.   We tested the numerical convergence of our results by making two additional runs for $\\beta=1.6$. In the first one, the number of particles was reduced to $N=10^5$, with the optimal value for the softening length of $\\log (\\epsilon/R_{\\rm ini})=-1.88$. Despite the fact that this run had the most poorly resolved ``crunch'' among all our models ($r_{\\rm min}/\\epsilon\\simeq 3.6$), all the derived model parameters were identical to the original run parameters within measurement errors.  In the second run, we again used a reduced number of particles $N=10^5$, but the parameter $\\epsilon$ was made much smaller than the optimal value: $\\log (\\epsilon/R_{\\rm ini})=-2.67$. In this run, the crunch is resolved very well --- $r_{\\rm min}/\\epsilon\\simeq 23$, but again we did not see significant deviations from the original run parameters. The results of the additional runs suggest that the number of particles $N$ and the value of the softening length $\\epsilon$ we use in our runs are adequate for resolving the violent relaxation process. ",
        "conclusions": "We have shown that a simple model of violent relaxation of a constant density, constant velocity dispersion stellar sphere produces clusters which are remarkably similar to the observed Galactic GCs. They have the correct surface density profiles, and a large number of correlations between cluster parameters are accurately reproduced by our models as long as we assume a constant initial density and initial velocity dispersion.  This result implies that all GCs were formed from gas with a universal value of density and temperature. In reality, the constant density, constant velocity dispersion initial setup of our models may correspond to isothermal turbulent cores of giant molecular clouds with a Gaussian density profile. Our simple picture of GC formation suggests that the conditions in the cluster-forming gas must be quite uniform across a large portion of the early universe."
    },
    "0403/astro-ph0403483_arXiv.txt": {
        "abstract": "{ The S235A-B star forming region has been extensively observed in the past from the radio to the near IR, but what was happening in the immediate surroundings of the water maser, placed in between the two nebulosities,  was still unclear because of insufficient resolution especially in the spectral range from the Far IR to the mm, even though there were sound indications that   new young stellar objects  (YSOs) are being formed there. We present here new high resolution maps at mm wavelengths  in different molecules (HCO$^+$, C$^{34}$S, H$_2$CS, SO$_2$ and CH$_3$CN), as well as in the 1.2 and 3.3 mm continuum obtained with the Plateau de Bure interferometer, and JCMT observations at 450 ${\\mu}$m and 850 ${\\mu}$m that unambiguously reveal the presence of new  YSOs placed in between the two \\HII\\ regions S235A and S235B and associated with the water maser. A molecular core and an unresolved source in the  mm and in the sub-mm are centred on the maser, with indication of mass infall onto the core. Two molecular bipolar outflows  and a jet originate from the same position. Weak evidence is found for a molecular rotating disk perpendicular to the direction of the main bipolar outflow. The derived parameters  indicate that one of the YSOs is an intermediate luminosity object ($L \\sim10^3L_{\\odot}$) in a very early evolutionary phase, embedded in a molecular core of $\\sim100~M_{\\odot}$, with  a temperature of 30 K. The main  source of energy for the YSO  could come  from gravitational infall, thus making of this YSO  a rare example of intermediate luminosity protostar representing a link between the earliest evolutionary phases of massive stars and low mass protostars of class 0--I. ",
        "introduction": "S235A and S235B are two close-by (separation $\\sim$1\\amin) small patches of nebulosity ($\\sim$20\\asec\\ and $\\sim$10\\asec\\ in diameter, respectively), so called because they are located about 10\\amin\\ south of the more extended \\HII\\ region S235 (Sharpless~\\cite{S}; Felli \\& Churchwell~\\cite{FC}), but most probably not directly related to it (Israel \\& Felli~\\cite{IF}). About 3.5 arcmin south of S235A-B there is another patch of nebulosity called S235C and identified with a small \\HII\\ region (Israel \\& Felli~\\cite{IF}).  Early interferometric radio continuum observations have shown that S235A is a young \\HII\\ region, even though far beyond the earliest ultracompact (UC) evolutionary phase, while no radio continuum emission comes from S235B (Israel \\& Felli~\\cite{IF}). The interest in these two nebulosities was increased by the fact that an H$_2$O maser is present in the same area (Henkel et al.~\\cite{HHG}; Comoretto et al.~\\cite{COMORE}), a well known  landmark of recent star formation (see e.g. Tofani et al.~\\cite{tof95}). With the resolution available at that time it was not clear whether  the maser was associated with S235A, S235B or simply happened to be in the same  star forming region (Israel \\& Felli~\\cite{IF}). Single dish observations have also revealed the presence of a methanol maser (Nakano \\& Yoshida~\\cite{NY}; Haschick et al.~\\cite{HMB}) and of an SiO maser (Harju et al.~\\cite{HLBZ}).  To ascertain the true relationship between S235A-B and the water maser and establish which is the source of energy supply for the water maser emission, high resolution radio observations  were performed with the VLA in the A configuration in the water line and in the continuum (Tofani et al.~\\cite{tof95}), which led to  the result  that the  H$_2$O maser is well outside the boundary of the S235A \\HII\\ region, away from S235B and that no point-like radio continuum source (down to a flux density of $\\sim$0.5 mJy and with a resolution of 0\\farcs3) is associated with  it, i.e. no trace of a UC \\HII\\ region near the maser was found.  At that time it was  becoming clear that  water masers without associated compact radio continuum emission could trace one of the earliest phases in the formation of a young stellar object (YSO), in which the UC \\HII\\ region ionized by the early type star  (if at all present) is so dense and optically thick to escape detection at radio wavelengths (Codella et al.~\\cite{cfn}, \\cite{ctc}).  The presence of a molecular cloud in the S235A-B region was known from earlier low resolution observations in several molecular lines (Evans \\&  Blair~\\cite{EB}; Lafon et al.~\\cite{LDBL}; Nakano \\& Yoshida~\\cite{NY}). To test the hypothesis that the water maser indicated the location of one or more new  YSOs independent from S235A-B, and get more information on the environment in which the water maser was located, new maps in molecular tracers of high density gas (Cesaroni et al.~\\cite{CFW}) and near-IR observations of the region (Felli et al.~\\cite{FTVW}) were made.  The results confirmed the validity of the hypothesis: the water maser is located at the centre of a high density small diameter molecular core, which peaks just between the S235A and S235B nebulosities. A cluster of near-IR sources  is  also located in the same area. The source with the largest IR excess (called M1) is  placed close to the water maser. The picture before the present observations is summarized in Fig.~\\ref{hcnk} which shows the overlay of the HCN map obtained with  a resolution of 27\\asec\\ (Cesaroni et al.~\\cite{CFW}) with a K-band (2.2 ${\\mu}$m) map of the same area (Felli et al.~\\cite{FTVW}). All the indications suggest  that the water maser,  M1 (both indicated in Fig.~\\ref{hcnk}) and the HCN clump are all part of a much younger episode of star formation with respect to the S235A-B nebulosities, not necessarily directly related to them.  For completeness it should be noted that the VLA observations of the water maser (Tofani et al.~\\cite{tof95}) refer to the component at $ \\sim-60$ \\kms, which happened to be flaring at the time of the  observations (1992) and was the only feature above the noise in the observed bandwidth (from --78 to --40~\\kms). During the continuing patrol of the water maser emission with the Medicina radiotelescope, started in 1988, this velocity component has very seldom emerged again from the noise level, while two other velocity components between  $ \\sim-2.5$  and 2.5 \\kms\\ are frequently observed (Felli et al., in preparation). A partial report of the results of the Medicina patrol is  given in Felli et al.~(\\cite{FTVW}).  In the present work we investigate with better resolution the region around the water maser to study in more detail the YSOs that presumably are  harboured in the centre of the molecular cloud. We present arcsec resolution maps at mm wavelengths  in different molecules that trace high density gas (HCO$^+$, C$^{34}$S, H$_2$CS, SO$_2$ and CH$_3$CN), as well as in the 1.2 and 3.3 mm continuum sensitive to  thermal dust emission.  To obtain the bolometric luminosities of the YSOs and their spectral energy distribution (SED), high resolution far IR observations are needed, since at these wavelengths YSO spectra are  expected to have their maximum. We inspected the IRAS HIRES  maps (Aumann et al.~\\cite{afm}) of the region around  the water maser (which contains IRAS\\,05375+3540) in the four IRAS bands, but the resolution is  inadequate for separating the contribution of S235A and  S235B from that of the YSO.  We also present here archive JCMT observations at 450 ${\\mu}$m and 850 ${\\mu}$m that unambiguously demonstrate that the sub-mm emission peaks on the YSO associated with the water maser and separate its emission  from that of S235A-B.  For consistency with previous work a distance of 1.8 kpc will be assumed (Nakano \\& Yoshida~\\cite{NY}).  \\begin{figure} \\centerline{\\psfig{figure=0905fig1.ps,width=8.8cm}} \\caption[]{Overlay of the Pico Veleta HCN map (contours, resolution 27\\asec) from Cesaroni et al.~(\\cite{CFW}) with the TIRGO K-band image of the S235A-B region (grey scale) from Felli et al.~(\\cite{FTVW}). The VLA position of the water maser from Tofani et al.~(\\cite{tof95}) is indicated with a cross (much larger than the positional uncertainty which is 0.1\\asec). M1, the K-band  source with the largest infrared excess, is also indicated. } \\label{hcnk} \\end{figure} ",
        "conclusions": "The detection of a mm and sub-mm core  at the position of the water maser in S235A-B, as well as of a core  in all the molecular lines  finally settles the question which is the source that provides the energy for the maser emission and indicates that a new YSO (or a cluster of YSOs) is located in this region and is the powering source of the maser emission. The two close-by nebulosities S235A-B do not seem to be directly related to the YSO, even though belonging to the same star forming complex. Consequently, star formation has not been occurring simultaneously in the extended CO molecular cloud, but rather progressing slowly from the outer parts  of the molecular cloud, where S235A and B are placed, to the denser inner part of it, where new YSOs are just being formed.  The present example and a look at the literature on this star forming region indicated the complexity in finding new intermediate luminosity YSOs when they are located in regions where high-mass star formation is already in progress and where their emission is completely obliterated by that of the more evolved regions, especially when the search occurs in the classical domains which were originally used to find them, e.g. cm radio continuum, near IR and  H$\\alpha$ emission. Only water masers, as well as other masers, e.g. methanol (Walsh et al.~\\cite{wmabl}), have remained as good indicators of YSOs (when they are present) and only high resolution mm continuum and molecular line observations are able to pinpoint the YSO emission from that of other close-by more evolved sources (Jenness et al.~\\cite{jsp}).  The detection of two bipolar outflows in S235  is another good indicator of recent star formation. It indicates that more than one YSO might be present or, but more difficult to explain with the present data, multiple episodes of bipolar outflow might originate from the same YSO. Larger scale bipolar outflows detected in this region with lower resolution might be remnants of previous episodes of star formation, but do not seem to be directly related to the two outflows produced by the presently detected YSOs.  Sulphur-bearing molecules (H$_{2}$CS and SO$_{2}$) have been detected, similarly to hot cores. Part of the emission appears related with the outflows, suggesting a shock origin, but most of it clearly arises from the molecular core.  The luminosity of the YSO powering the NE--SW flow is estimated to be in the range 1000--3500~$L_\\odot$, less than that of S235A (1.1 10$^4$~\\lsol) and the dynamical time of the order of 10$^4$~yr, much younger than that of the \\HII\\ region around S235A (the canonical value for well developed \\HII\\ region is 10$^6$~yr). The temperature of the molecular gas of the core ($\\sim$30 K) is smaller than that found in hot cores ($\\geq$100 K) further confirming that the newly detected YSO is intermediate between the massive objects found in hot cores (Kurtz et al.~\\cite{kcchw}) and the low mass class 0 objects (Andr\\'e et al.~\\cite{awb}).  In conclusion, the newly found YSO in between S235A and S235B is a rare example of an intermediate luminosity object  and its morphology may be used to connect early evolutionary phases of massive stars with those of low mass protostars of class 0--I.  We find weak evidence for a possible disk in the C$^{34}$S line,  perpendicular to the direction of the main bipolar outflow. If confirmed, this  is a further proof of how the complex morphology of star forming regions in the earliest  phases can be unveiled only by high resolution molecular observations in high density tracers that can separate the weak emission of the disk from the overwhelming emission of the molecular cloud.  While revealing a new wealth of information, at the same time the latest  observations leave some unanswered questions, as for instance the origin of the thermal jet and the nature of the two-velocity structure of the molecular cloud. Are there multiple YSOs in the region of the mm core  or a single YSO with multiple episodes of outflows? Why there is no thermal jet associated with the  main bipolar outflow? Why is the velocity of the water maser features so different from that of the molecular core? Work is in progress to try to give an answer to these questions."
    },
    "0403/astro-ph0403160_arXiv.txt": {
        "abstract": "We demonstrate a significant ($> 5\\sigma$) correlation between the mean color of \\emph{metal-poor} globular cluster (GC) systems and parent galaxy luminosity. A Bayesian Markov Chain Monte Carlo method is introduced to find the mean color, and is easily generalizable to quantify multimodality in other astronomical datasets. We derive a GC color--galaxy luminosity relation of the form $Z \\propto L^{0.15\\pm0.03}$. When combined with evidence against a \\emph{single} primordial GC metallicity--galaxy luminosity relation for protogalactic fragments, the existence of such a correlation is evidence against both accretion and major merger scenarios as an explanation of the entire metal-poor GC systems of luminous galaxies. However, our relation arises naturally in an \\emph{in situ} picture of GC formation, and is consistent with the truncation of metal-poor GC formation by reionization. A further implication is that the ages of metal-poor GCs in dwarf galaxies constrain the main epoch of galaxy formation in hierarchical models. If the ages of old metal-poor GCs in Local Group dwarfs ($\\ga 11$ Gyr) are typical of those in dwarfs elsewhere, then the bulk of galaxy assembly (at least in clusters and groups) must have occurred at $z \\ga 2.5$, contrary to the predictions of some structure formation models. ",
        "introduction": "A key development in the study of globular clusters (GCs) in external galaxies has been the discovery that most large galaxies have bimodal GC color distributions (e.g., Zepf \\& Ashman 1993; Forbes, Brodie \\& Grillmair 1997; Gebhardt \\& Kissler-Patig 1999; Kundu \\& Whitmore 2001a; Larsen \\etal~2001). These are usually described as blue (metal-poor) and red (metal-rich) GC subpopulations, although additional substructure may be present. The red GC system properties appear to be intimately tied to those of their parent galaxy, suggesting that the red GCs formed along with the bulk of the galaxy field stars (Forbes \\etal~1997; Forbes \\& Forte 2001; Larsen \\etal~2001). In both spirals and ellipticals, they are thought to be associated with the bulge/spheroid component (Forbes, Brodie \\& Larsen 2001).  The blue GCs are among the oldest and most metal-poor stellar systems observable. Therefore, they provide a probe of very early epochs of star formation in the universe. Whether or not the properties of blue GCs correlate with the mass of their parent galaxy has been controversial (Forbes \\etal~1997, C{\\^ o}t{\\' e} \\etal~2000; Burgarella, Kissler-Patig \\& Veronique~2001; Forbes \\& Forte~2001; Larsen \\etal~2001, Lotz \\etal~2003), and no clear demonstration of such a relation exists in the literature. However, the issue is an important one in the context of GC and galaxy formation. If a correlation exists, it implies that the blue GCs, or at least a significant proportion of them, ``knew'' about the galaxy to which they would ultimately belong. This indicates that their formation was affected by the local conditions and that they may have already been contained within the dark matter halo of their final galaxy. The detailed chemical and age structure within the blue GC systems of galaxies of various types and environments would then offer one of the few observational constraints on the properties of the protogalactic clouds which combined to build the galaxies that we observe today. Such a correlation would also rule out any formation mechanism whereby all metal-poor GCs form \\emph{completely} independently of a host galaxy (e.g., Peebles \\& Dicke 1968).  Our aim here is to consider galaxies over a large luminosity range, use only high-quality data, and analyze the photometry in a uniform manner. In this way we will reduce the random and systematic errors that could disguise the existence of a blue GC--host galaxy correlation. In particular, we have added new data on the GC systems of dwarf and low-luminosity elliptical galaxies, and utilized a Bayesian statistical method to find the peak of the blue GC color distribution. ",
        "conclusions": "The classic Fall \\& Rees (1988) classification of GC formation scenarios into primary, secondary, and tertiary has in the last decade given way to a new triad: accretion (e.g., C{\\^ o}t{\\' e}, Marzke, \\& West 1998), major merger (e.g., Ashman \\& Zepf 1992), and \\emph{in situ} (e.g., Forbes \\etal~1997). The primary focus of each of these formation scenarios is the GC systems of giant ellipticals (gEs). However, much of the following discussion applies also to the formation of less-luminous ellipticals (Es) and disk galaxies.  C{\\^ o}t{\\' e} \\etal~(1998, 2000, 2002) were able to produce bimodal metallicity distributions for the GC systems of luminous galaxies under the fundamental assumption that all galaxies have one intrinsic population of GCs, which depends on galaxy luminosity. They used observations of a sample of dwarf and giant galaxies to find a ``primordial'' relation between GC metallicity and parent galaxy luminosity. By varying an assumed mass spectrum of protogalactic fragments, they were able to produce a wide variety of GC metallicity distributions, some of which matched the bimodal form seen in most giant galaxies. It should be noted that they assume dwarf galaxies in the local universe (and especially in the Local Group) are the surviving counterparts of these protogalactic fragments. If instead these gaseous fragments merged themselves out of existence at $\\sim$ the epoch of blue GC formation, then the model is essentially equivalent to an \\emph{in situ} scenario (see below).  The C{\\^ o}t{\\' e} \\etal~(2002)~primordial mean metallicity-$M_V$ relation (see their Figure 1) is $V-I = -0.055M_{V} + 0.04$. This relation predicts that galaxies with $M_{V} = -17.5$ will have one intrinsic population of GCs with $V-I = 1.00$. By contrast, the blue GC relation we derived above and the red GC relation from Larsen \\etal~(2001) predict peaks at $V-I = 0.90$ and $V-I = 1.09$, respectively. In Figure 1 we have overplotted red GC peaks for fifteen galaxies from Larsen \\etal~(2001). To these we have added four less-luminous galaxies (which are in our blue sample as well) for which we have remeasured the red peaks ourselves, using original data from Larsen \\etal~and Kundu \\& Whitmore (2001a, 2001b). The dotted line is a new linear fit to the red peaks, weighted by the number of GCs in each peak. The dashed line is the C{\\^o}t{\\' e} \\etal~relation. While the observations are marginally consistent with the primordial relation at the blue and red ends of the luminosity range, in the intermediate-metallicity region between $M_{V} \\sim -17$ and $-18$ there is a clear, $\\sim 0.1$ mag difference between the observed peaks and the primordial prediction. Thus, the C{\\^ o}t{\\' e} \\etal~primordial relation for a single intrinsic population of GCs appears to be less consistent with current observations than the existence of two separate (blue and red) GC metallicity--luminosity relations.  C{\\^ o}t{\\' e} \\etal~(2002) showed that by using their primordial relation, a wide variety of GC color distributions could be produced. An examination of their Figure 1 suggests that, given appropriate input assumptions, the accretion model could reproduce the blue GC relation. However, the ability of the accretion model to successfully reproduce our observed relation is critically dependent on the validity of the primordial relation. By contrast, a metal-poor GC--galaxy luminosity correlation arises naturally under \\emph{in situ} scenarios for GC formation, e.g., Harris \\& Pudritz (1994) and Forbes \\etal~(1997). In such scenarios, enrichment of the interstellar medium is linked to the depth of the galactic potential, \\emph{independent} of the details of GC formation in a particular scenario.  A blue GC--galaxy relation places limits on the degree to which an intrinsic population of metal-poor GCs in a luminous galaxy has been affected by accretion and/or mergers, with substantial implications for galaxy formation models (see below). A caveat is that the cosmic variance in our model (1$\\sigma$ scatter of $\\sim 0.02$ mag) does not rule out the formation of the metal-poor GC systems of gEs like M49 and M87 from the accretion of several slightly less luminous galaxies with redder-than-average blue GCs. However, our correlation implies that \\emph{in the mean} the blue GC systems of gEs could not have formed primarily from the accretion of less luminous galaxies, especially $M_{V} > -18$ dwarfs, whose blue GCs have mean $V-I \\la 0.90$. The paucity of metal-poor field stars in the halos of luminous E galaxies (e.g., NGC 5128; Harris \\& Harris 2000) is another argument against significant accretion, since such stars would presumably accompany the accreted metal-poor GCs (Harris 2003). We are not arguing that accretion does not take place; witness the existence of the Sagittarius dwarf (Ibata, Gilmore, \\& Irwin 1994), the recently discovered Canis Major galaxy (Martin \\etal~2004; Forbes, Strader \\& Brodie 2004), relic stellar streams in the halo of M31 (Ibata \\etal~2001), and more distant systems (e.g., Forbes \\etal~2003). However, we do suggest that it is not the primary mechanism by which the GC systems of most luminous galaxies are built.  We note that a major merger scenario for the formation of gEs is also difficult to reconcile with our results, since the mean colors of metal-poor GCs in typical Sb/Sc spirals like the Milky Way and M31 are bluer than those of gEs (see Figure 1). The favored site for the formation of blue GCs in the major merger scenario is in pregalactic fragments (Ashman \\& Zepf 1992). Ashman (2003) states ``some variation in mean metallicity [of metal-poor GCs] would not rule out this option provided it did not correlate with properties of the current parent galaxy''. As noted above for the accretion model, while cosmic variance would allow the occasional merger of two $L^{*}$ spirals with redder-than-average blue GCs, the existence of a global correlation suggests that major mergers are unlikely to be the dominant mechanism for the creation of the GC systems of gEs. Independent arguments in favor of this point have been made based upon the relative \\emph{numbers} of metal-poor GCs in spirals and gEs (see Harris 2003 for a summary), though some recent work (e.g., Rhode \\& Zepf 2004)  suggests that these ``specific frequency'' arguments may only pose a problem for the blue GC systems of gEs.  We note that proponents of the merger scenario have themselves argued (see, e.g., the discussion in Ashman \\& Zepf 1998) that major mergers cannot be the \\emph{exclusive} mechanism for for forming gEs.  \\emph{In situ} scenarios have been criticized with straw man comparisons to monolithic collapse models (e.g., Eggen, Lynden-Bell, \\& Sandage 1962), though, for example, Forbes \\etal~(1997) argued that the initial collapse probably involved ``some chaotic merging of many small subunits'' as in the Galactic formation model of Searle \\& Zinn (1978). This idea was reflected in simulations of GC formation by Beasley \\etal~(2002). One of the most significant problems with the \\emph{in situ} scenario is the lack of a definitive physical mechanism for the truncation of metal-poor GC formation. An initial starburst, accompanied by the formation of metal-poor GCs, could conceivably ionize and/or remove a substantial mass of gas from inner regions of the protogalaxy, with the subsequent infall of enriched gas forming a population of metal-rich GCs and the bulk of the galaxy starlight (Forbes \\etal~1997). However, this feedback mechanism would be increasingly ineffective in more massive halos, with little suppression of star formation expected for halos with virial velocities $V > 100$ km/s (Dekel \\& Woo 2003). Thus, feedback is unlikely to be the source of the near-universal metallicity gap observed in the GC systems of massive galaxies with $V \\sim$ several hundred km/s. Cosmic reionization might also result in a rapid truncation of blue GC formation (e.g., Santos 2003), though some instead argue for reionization as the trigger for metal-poor GC formation (Cen 2001). Among the advantages of the reionization hypothesis are its consistency with early structure formation theories, as well as the prediction of the observed differences in metal-poor GC specific frequency between ellipticals and spirals. Since ellipticals (located in high-density peaks) collapse earlier than spirals of similar mass, they have a larger fraction of their final mass assembled and available for GC formation before reionization. By a similar argument, one might expect the metal-poor GC metallicity--galaxy luminosity relation to have a steeper slope in denser environments. A study comparing the blue GC systems of field and cluster galaxies might then offer the intriguing possibility of elucidating the relative roles of galaxy mass and environment in shaping the initial phase of local metal enrichment.  Our findings also have implications for popular hierarchical structure formation models (e.g., Kauffmann, White, \\& Guiderdoni~1993; Cole \\etal~2000; Somerville, Primack, \\& Faber 2001). For example, semi-analytic models generically predict a substantial fraction of mass assembly at $z \\la 1$. However, as already noted, relatively luminous Es cannot have accreted very many dwarf galaxies and their populations of GCs, since this would destroy our observed correlation. The ages of metal-poor GCs in dwarf galaxies place lower limits on the possible epochs of significant accretion. We note that our results do not constrain (i) the accretion of galaxies before the formation of their GC systems, (ii) the accretion of galaxies without GCs (which typically have $M_V > -13$), or (iii)  the formation of field galaxies. However, if the ages of old metal-poor GCs in Local Group dwarfs ($\\ga 11$ Gyr; Johnson \\etal~1999) are typical of those elsewhere (an untested assumption), then the bulk of galaxy assembly in groups and clusters must have occurred at $z \\ga 2.5$. For more likely mean GC ages of $\\sim 12$ Gyr, the primary epoch of mass assembly is pushed to $z \\sim 4$ or higher."
    },
    "0403/astro-ph0403356_arXiv.txt": {
        "abstract": "We present results from fully general relativistic three-dimensional numerical studies of thick-disk accretion onto a rapidly-rotating (Kerr) black hole with a spin axis that is tilted (not aligned) with the angular momentum vector of the disk. We initialize the problem with the solution for an aligned, constant angular momentum, accreting thick disk, which is then allowed to respond to the Lense-Thirring precession of the tilted black hole. The precession causes the disk to warp, beginning at the inner edge and moving out on roughly the Lense-Thirring precession timescale. The propagation of the warp stops at a radius in the disk at which other dynamical timescales, primarily the azimuthal sound-crossing time, become shorter than the precession time. At this point, the warp effectively freezes into the disk and the evolution becomes quasi-static, except in cases where the sound-crossing time in the bulk of the disk is shorter than the local precession timescale. We see evidence that such disks undergo near solid-body precession after the initial warping has frozen in. Simultaneous to the warping of the disk, there is also a tendency for the midplane to align with the symmetry plane of the black hole due to the preferential accretion of the most tilted disk gas. This alignment is not as pronounced, however, as it would be if more efficient angular momentum transport (e.g. from viscosity or magneto-rotational instability) were considered. ",
        "introduction": "\\label{sec:introduction} Quasars, active galactic nuclei (AGN), X-ray binaries, core-collapse supernovas, and gamma-ray bursts (GRBs) are among the most luminous objects in the universe and are therefore of considerable observational interest.  They are also all likely to contain a central, strongly gravitating, compact object and be powered by some form of accretion.  They are therefore also of considerable theoretical interest.  However, the nonlinear, time-dependent, multi-dimensional nature of the physical laws governing these objects makes a study of them from first principles difficult.  While such studies can provide critical insight, there is strong motivation to apply a numerical approach.  Many numerical simulations have been carried out over the past three decades studying the time-dependent evolution of accretion flows around black holes in the hydrodynamic \\citep[e.g.][]{wil72,haw84b,haw91} and MHD \\citep{koi99,gam03a,dev03b} regimes. However, due to the complex nature of this problem, a variety of simplifying assumptions are normally implemented. A common simplification in much of the early work was to assume a non-rotating (Schwarzschild) black hole, often approximated using a pseudo-Newtonian \\citep{pac80} potential.  A notable exception was the work of \\citet{wil72}, which considered the spherical infall of material with a non-zero specific angular momentum toward a Kerr black hole using the full metric, although restricted to two spatial dimensions. Although the assumption of a non-rotating black hole considerably simplifies the problem while still allowing useful insight, this does not appear to be well represented among observed black-hole systems \\citep[e.g.][]{elv02,yu02}, nor does it agree with studies of black hole spin evolution \\citep{gam04,vol04} which suggest that most black holes are likely to be spinning rapidly.  The spin of a black hole (often denoted by the specific angular momentum $a = J/M$) can have important consequences on the accretion rate and the geometry of the accretion flow. One obvious example is the dependence of $r_{ms}$, the radius of the marginally stable circular orbit, on the spin parameter [$r_{ms}=6r_G$ for $a=0$ (non-rotating black hole); $r_{ms}=r_G$ for $a=+M$ (maximal prograde rotation); $r_{ms}=9r_G$ for $a=-M$ (maximal retrograde rotation), where $r_G=GM/c^2$ is the gravitational radius \\citep{bar72}]. Another example is the reversal of flow direction in the vicinity of a rotating black hole for a retrograde accretion disk \\citep[see Figure 6 of][]{dev02}.  Another common simplification in black hole accretion models has been the assumption that the midplane of the accretion flow is aligned with the symmetry plane of the black hole. This has historically been necessitated in numerical work by limitations in computing resources which have often limited simulations to two-dimensional axisymmetry. The justification for this treatment often relies on the Bardeen-Petterson effect \\citep{bar75} through which a spinning black hole may align the inner region of the accretion flow with the symmetry plane of the black hole. However, now that computational resources are available to perform three-dimensional simulations of black-hole accretion, there are many compelling reasons to revisit the question of black-hole - disk alignment. First, the formation avenues for many black-hole - disk systems favor, or at least allow for, an initially tilted configuration (see \\citet{fra01a} for a review of the arguments). Second, although the Bardeen-Petterson effect may align the inner disk, this alignment may not extend nearly as far out as was originally expected [$15-30 r_G$ \\citep{nel00} rather than $100-1000 r_G$].  Many numerical simulations extend to larger radii than this and, therefore, may not be justified in assuming that the midplane of the disk is aligned with the symmetry plane of the black hole. Finally, although there is a residual torque between the tilted outer disk and black hole which will eventually completely align the system \\citep{ree78,sch96,nat98}, this happens over such a long timescale that these systems remain tilted for most of their observational lifetimes.   If an accretion disk is misaligned or tilted, it will be subject to differential Lense-Thirring precession.  This is ultimately what gives rise to the Bardeen-Petterson effect. For an ideal test particle in a slightly tilted orbit at a radius $r$ around a black hole of mass $M$ and specific angular momentum $a$, gravitomagnetic torques cause the orbit to precess at a frequency $\\Omega_{LT}=2aM/r^3$.  Close to the black hole, this is comparable to the orbital frequency [$\\Omega_{Kep}=(M/r^3)^{1/2}$].  However, because of its strong radial dependence, Lense-Thirring precession becomes much weaker far from the hole.  Nevertheless, since it has a cumulative effect, it can build up appreciably over a sufficiently large number of orbits.  In an accretion disk, this cumulative build-up is limited by dynamic responses within the disk. For a steady-state disk, Lense-Thirring precession is expected to be important out to a unique, nearly constant transition radius \\citep{bar75,kum85}.  A comprehensive understanding of tilted accretion disks can be aided by the insights of direct numerical simulations.  An important first step was taken by \\citet{nel00}, who utilized a three-dimensional Newtonian SPH code to study tilted thin accretion disks subject to a post-Newtonian gravitomagnetic torque.  In the current work we present results from three-dimensional, fully relativistic simulations of tilted black-hole accretion flows.  A general relativistic treatment ensures that all important relativistic features, such as the cusp in the potential and the Einstein and Lense-Thirring precessions of the orbits, as well as any higher order couplings of these features, are treated properly. In this first study, we consider relatively simple, inviscid, thick accretion disks. These disks are well studied in the literature and provide a convenient starting point for our work. We should mention, however, that effective angular momentum transport (such as through viscous coupling as considered in \\citet{nel00}), would likely have a considerable impact on our results.  Unless otherwise stated, we use geometrized units ($G=c=1$) throughout this paper.  Units of length are parameterized in terms of the gravitational radius of the black hole, $r_G=GM/c^2$. We also adopt the standard convention in which Greek (Latin) indices refer to 4(3)-space components. ",
        "conclusions": "\\label{sec:discussion}  We have performed the first fully relativistic numerical studies of tilted accretion disks around Kerr black holes.  Such disks are subject to Lense-Thirring precession, resulting in a torque that tends to twist and warp the disk. For the thick disks considered here, we find that the nature of this precession depends primarily on the sound speed in the disk.  Whenever $\\tau_{LT}<\\tau_{cs}$ in the bulk of the disk, it undergoes differential precession out to a transition radius set by $\\tau_{LT} \\approx \\tau_{cs} \\lesssim \\tau_{acc}$, which occurs at a relatively modest distance of $r/r_G \\approx 7$.  In other disk models, such as viscous Keplerian disks, this transition may occur at somewhat larger radii \\citep[$\\sim15-30 r_G$,][]{nel00}. We find that the location of this transition radius does not depend strongly on the size, mass, mass accretion rate, or tilt of the disk. The initial tilt does, however, affect the overall mass accretion rate for our models, since higher tilt models begin further from equilibrium within the black hole potential.  Whenever $\\tau_{LT}>\\tau_{cs}$ for the bulk of the disk, we find that it undergoes near rigid-body precession after a short initial period of differential precession. Such rigid-body precession of a thick disk could have important astrophysical consequences. For instance, this is the basis of at least one model attempting to explain the 106 day variability observed from Sgr A* \\citep{liu02}, the $2.6\\times10^6 M_\\odot$ black hole at the Galactic center. It could also be important for understanding the jet precession observed in Galactic microquasars such as SS 433 and Her X-1.   One obvious shortcoming of this work is that we have ignored angular momentum transport mechanisms that are likely present in most astrophysical disks [e.g. viscosity or the magneto-rotational instability \\citep{bal91}]. Angular momentum dissipation may allow the inner region of the tilted disk to align more quickly with the symmetry plane of the black hole than what we see in the inviscid disks considered here \\citep{bar75,kum85,sch96}. The alignment may also be more efficient in viscous disks, resulting in a sharper transition \\citep{nel00}. A sharp transition from an inner aligned disk to a tilted outer disk may be useful in explaining a variety of phenomena observed in black-hole systems, such as quasi-periodic oscillations \\citep{fra01a} as well as certain spectral features."
    },
    "0403/astro-ph0403430_arXiv.txt": {
        "abstract": "We have used the Australian Long Baseline Array (LBA) to produce milliarcsecond images of five sites of methanol (\\methanol ) maser emission at 6.7~GHz. These are all sites that have linear morphologies at arcsecond resolutions, which have been hypothesised to be due to the masers forming in edge-on circumstellar disks. We find that a simple disk model cannot explain the observations. We discuss various alternatives, and suggest a new model which explains how linear velocity gradients can be produced in methanol masers that arise in planar shocks propagating nearly perpendicular to the line of sight. ",
        "introduction": "\\label{s:intro}  Massive stars play a crucial role in the evolution of galaxies. They are the cauldrons where heavy elements are produced. They provide turbulent energy to the interstellar medium, are involved in the production and destruction of molecular clouds and may regulate the rate of star formation. Despite their importance, the processes leading to massive star formation are poorly understood (e.g. \\citet{garay_99}, \\citet{mckee_03}). Apart from theoretical complexities inherent in such a multiple-process problem \\citep{yorke_02}, the chief obstacle to our understanding is observational -- massive star formation occurs in fewer regions (as they form quickly) that are, on average, further away (than regions of low-mass star formation) and are highly obscured at optical wavelengths (\\citet{churchwell_90}, \\citet{plume_92}).  Interstellar masers of three molecules, \\water , OH and methanol, are both powerful and widely associated with active high-mass star formation regions (\\citet{forster_89}, \\citet{menten_91}). Thus, they have long been recognized as a potentially powerful tool to study the physical conditions and kinematics of massive star formation sites. However, the complex distribution of \\water\\ and OH masers has limited their role in the elucidation of massive star formation. Further, both \\water\\ and OH masers are also found towards other objects such as evolved low-mass stars \\citep{wilson_72} and the centres of galaxies \\citep{santos_79}.  The two strongest class II methanol masers (i.e. those closely associated with OH masers and strong far-infrared emission) were discovered at 12.178593~GHz by \\citet{batrla_87} and at 6.668518~GHz by \\citet{menten_91}. These methanol masers are closely associated with high-mass star forming regions \\citep{norris_88, walsh_97}.  Sensitive searches for 6.7-GHz methanol masers towards low-mass star forming regions \\citep{minier_03}, and nearby active galaxies \\citep{phillips_1998a, darling_03} have failed to detect any. While untargeted searches covering large areas of the Galactic Plane have detected many new 6.7-GHz methanol masers not associated with known massive star forming regions, to date none have been confirmed as being associated with any other type of astrophysical object \\citep{caswell_96,ellingsen_96, szymczak_02}. Thus class II methanol masers prove more efficacious for finding and studying high-mass star forming regions than either \\water\\/ or OH masers.  The brightness of methanol masers allows high resolution interferometric observations to be made, which provide accurate measurements of position, velocity and dimensions of the individual maser components. Previous interferometric observations have generally used connected element interferometers \\citep{norris_88, norris_93, phillips_98, walsh_98} that offer resolutions of the order of arcseconds (which can be improved upon by super resolving the velocity features, see \\citet{phillips_98}, hereafter P98), though some VLBI (Very Long Baseline Interferometry) experiments \\citep{menten_88, menten_92, norris_98, minier_98, moscadelli_99} with milliarcsecond resolution have been performed. VLBI resolution images are vital to uncovering the structure of these masing regions and their paucity comes from the absence of telescopes (let alone arrays of telescopes) able to observe at these frequencies until relatively recently. This situation has now changed with both the Australian Long Baseline Array (LBA\\footnote{The Long Baseline Array is part of the Australia Telescope which is funded by the Commonwealth of Australia for operation as a National Facility managed by CSIRO.}, \\citet{ojha_03}) and the European VLBI Network (EVN) covering the 6.7-GHz transition, while the LBA and the Very Long Baseline Array (VLBA) covers the 12.2-GHz transition.  Many of the observed masing sites have individual masers located along lines or arcs, often with a near-monotonic velocity gradient along the line. It has been suggested \\citep{norris_93, norris_98, phillips_98} that these linear structures trace masers embedded in an edge-on disk surrounding a young star. Values of disk radii and enclosed masses derived from modelling these sources assuming Keplerian rotation, agree with theoretical models of accretion disks around massive stars \\citep{lin_90, hollenbach_94}.  \\citet{walsh_98} found 36 of 97 maser sites to be linearly extended and, while the circumstellar disk hypothesis is consistent with their data, they considered it unlikely as the derived values of the enclosed mass are too high in some cases. They suggested the linear geometry of 6.7-GHz methanol masers resulted from shocks arising behind a smooth shock front. As shocks would not produce the velocity gradient they saw in many of their maser sites, \\citet{norris_93} had rejected shocks as a possible explanation. Pointing out that most maser sites in their larger sample do not show a systematic velocity gradient, \\citet{walsh_98} argue that a velocity gradient is not typical in maser sites. The VLBI observations of \\citet{minier_00} find the linear extent of maser regions to be about 10 times smaller than that seen at lower resolutions. This yields sub-solar masses for the ``massive'' star. Both edge-on shock and edge-on disk models produce linear features across the sky but shock fronts are not normally so well ordered. At milliarcsecond (mas) resolutions we probe resolutions of the order of tens of AU, the perfect scale to investigate these differences.  The case for disks, or the other models, will only become clearer after a significant increase in the sample size of methanol masers imaged at very high resolution. We present the first VLBI images of five sources which have not previously been imaged at mas resolution. This allows us to directly compare conclusions based on the ATCA (Australia Telescope Compact Array) data to those reached with more than two orders of magnitude better resolution.  In Section \\S{\\ref{s:obsdata}} we describe our observations and the data reduction path. Our results are presented in \\S{\\ref{s:results}} and their implications for the current models are discussed in \\S{\\ref{s:discuss}}. \\S{\\ref{s:conclude}} sets out our conclusions.     \\begin{table*} \\centering \\begin{minipage}{150mm}  \\caption{Observed Maser Sites. Absolute positions for the strongest features are from (1) Phillips \\etal (1998) (2) Caswell (1997) and (3) Caswell \\etal (1995). Relative positions are with respect to these absolute positions. Diameters, where calculated, are based on the nearer kinematic distances and the enclosed masses were calculated from a linear least-squares fit in a velocity-major axis diagram under the assumption that the masers arise in an edge-on disk.} \\label{tab:one}  \\begin{tabular}{@{}lclrrrrrrc@{}} \\hline &     Right   &               &Mean LSR   &          &Peak zero {\\em uv} &          &      &  \\\\ Source &Ascension&Declination&Velocity& Distance & Flux Density & Diameter & Mass &Position\\\\ &  (J2000)    &   (J2000) &(km~s$^{-1}$)& (kpc)    &   (Jy)   &   (AU)   & $\\msol$  &  Reference \\\\ \\hline G327.120$+$0.511 &  15:47:32.71 &$-$53:52:38.5 & $-$87.3   & 5.5/8.8  &  27.5    &1900      &17&1 \\\\ G327.402$+$0.444 & 15:49:19.50 &$-$53:45:13.9  & $-$81.8   & 5.1/9.2   &  233     & 239      &0.9&1 \\\\ G327.402$+$0.444E& 15:49:19.72 &$-$53:45:14.6  &          &           &  1.5     &   -      & -\\\\ G327.402$+$0.445 & 15:49:19.33 &$-$53:45:14.4  &          &           &  18      &  1392    & 9.2\\\\ G328.237$-$0.548 & 15:57:58.27  &$-$53:59:23.1 & $-$43.8  &  3.0     &  195     & 463      &0.4&1  \\\\ G328.237$-$0.548W& 15:57:58.20  &$-$53:59:23.2 &         &           &  5.5     & 560      & 2.5 \\\\ G328.254$-$0.532 & 15:57:59.73  &$-$53:58:00.8 & $-37.5$  & 2.6      &  47       &  141    & 0.0\\\\ G328.808$+$0.633 &  15:55:48.50 &$-$52:43:06.64 & $-$45.1 & 3.1      & 148      & 1713     &0.4&2\\\\ G328.809$+$0.633 &  15:55:48.62 &$-$52:43:06.55 &         &           & 29.9     & -        &- & \\\\ G329.029$-$0.205 &  16:00:31.80 &$-$53:12:49.7 & $-$37.3  &  2.6      &  114     &  164     &1.0&3\\\\ G329.031$-$0.198& 16:00:30.33 &$-$53:12:27.4 & $-$45.6   & 3.1       &  5.7     &  612     &  4.2     &  \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*} ",
        "conclusions": "\\label{s:conclude}  We have observed five sites of methanol maser emission, and imaged eleven clusters with milliarcsecond resolution. These high resolution images argue strongly that the linear structures commonly observed are more likely linear shock fronts than Keplerian disks and some of the predictions of the former model are fulfilled. We are, however, not able to provide conclusive evidence of this, and VLBI observations of a large sample is required. We have proposed a further series of experiments in which we will collect full polarisation data which should determine the relationships between the maser spots."
    },
    "0403/astro-ph0403606_arXiv.txt": {
        "abstract": "We analyze interstellar absorption observed towards two subdwarf O stars, JL~9 and LSS~1274, using spectra taken by the {\\em Far Ultraviolet Spectroscopic Explorer} (FUSE).  Column densities are measured for many atomic and molecular species (H~I, D~I, C~I, N~I, O~I, P~II, Ar~I, Fe~II, and H$_{2}$), but our main focus is on measuring the D/H ratios for these extended lines of sight, as D/H is an important diagnostic for both cosmology and Galactic chemical evolution.  We find ${\\rm D/H}=(1.00\\pm 0.37)\\times 10^{-5}$ towards JL~9, and ${\\rm D/H}=(0.76\\pm 0.36)\\times 10^{-5}$ towards LSS~1274 (2$\\sigma$ uncertainties).  With distances of $590\\pm 160$~pc and $580\\pm 100$~pc, respectively, these two lines of sight are currently among the longest Galactic lines of sight with measured D/H.  With the addition of these measurements, we see a significant tendency for longer Galactic lines of sight to yield low D/H values, consistent with previous inferences about the deuterium abundance from D/O and D/N measurements.  Short lines of sight with H~I column densities of $\\log N({\\rm H~I})<19.2$ suggest that the gas-phase D/H value within the Local Bubble is ${\\rm (D/H)_{LBg}}=(1.56 \\pm 0.04) \\times 10^{-5}$. However, the four longest Galactic lines of sight with measured D/H, which have $d>500$~pc and $\\log N({\\rm H~I})>20.5$, suggest a significantly lower value for the true local-disk gas-phase D/H value, ${\\rm (D/H)_{LDg}}=(0.85\\pm 0.09)\\times 10^{-5}$.  One interpretation of these results is that D is preferentially depleted onto dust grains relative to H and that longer lines of sight that extend beyond the Local Bubble sample more depleted material.  In this scenario, the higher Local Bubble D/H ratio is actually a better estimate than ${\\rm (D/H)_{LDg}}$ for the true local-disk D/H, ${\\rm (D/H)_{LD}}$.  However, if ${\\rm (D/H)_{LDg}}$ is different from ${\\rm (D/H)_{LBg}}$ simply because of variable astration and incomplete ISM mixing, then ${\\rm (D/H)_{LD}}={\\rm (D/H)_{LDg}}$. ",
        "introduction": "One of the great successes of the Big Bang theory is that it predicts the light element abundances in the universe with reasonable accuracy. However, the exact abundance predictions depend on the cosmic baryon density, $\\Omega_{b}$.  The abundance most sensitive to $\\Omega_{b}$ is deuterium, so measuring the primordial deuterium-to-hydrogen ratio has become an important method for constraining $\\Omega_{b}$ in cosmological models \\citep[e.g.,][]{amb85,sb01}. A good estimate of the primordial D/H ratio can be obtained from IGM absorption components seen towards distant quasars, since the metal abundance in the IGM is small, indicating that the gas has experienced very little astration.  \\citet{dk03} compile several such measurements \\citep[e.g.,][]{jmo01,mp01,sal02} and report a primordial D/H of ${\\rm (D/H)_{prim}}=(2.78^{+0.44}_{-0.38})\\times 10^{-5}$, although there is significant scatter in the individual measurements\\footnote{Unless otherwise noted the quoted uncertainties for averaged quantities are 1$\\sigma$ standard deviations of the mean, but quoted uncertainties for individual measurements are 2$\\sigma$.}.  This value is consistent with primordial deuterium abundances inferred from recent WMAP measurements of the cosmic microwave background, combined with Big Bang nucleosynthesis calculations \\citep{dr03}.  Since deuterium is gradually destroyed in the interiors of stars, it is expected that the D/H value should be lower in places that have experienced a lot of stellar processing.  \\citet{krs04} find a value of ${\\rm D/H}=(2.2\\pm 0.7)\\times 10^{-5}$ for Complex~C, a high velocity cloud falling onto the Milky Way, which has a low metallicity but presumably has experienced more stellar processing than most IGM material.  The Complex~C D/H measurement may be slightly lower than the ${\\rm (D/H)_{prim}}$ value quoted above, although the uncertainties are too large to provide complete confidence in this result.  The D/H ratio in the local disk region of our galaxy, ${\\rm (D/H)_{LD}}$, should be significantly lower than both the Complex~C and IGM measurements due to significantly more stellar processing.  Comparing ${\\rm (D/H)_{LD}}$ to ${\\rm (D/H)_{prim}}$ provides an excellent indication of the amount of stellar processing experienced by interstellar material in our Galaxy, and therefore provides a useful test of Galactic chemical evolution models.  Many measurements have been made of the gas-phase local-disk D/H ratio in the Galaxy [${\\rm (D/H)_{LDg}}$]. These measurements do not account for deuterium that may be locked into molecules and dust, but it has nevertheless been assumed in the past that ${\\rm (D/H)_{LDg}}\\approx {\\rm (D/H)_{LD}}$.  Unfortunately, the D/H measurements for various Galactic lines of sight have not collectively provided an unambiguous value for ${\\rm (D/H)_{LDg}}$.  There seems to be substantial variation in the ${\\rm (D/H)_{LDg}}$ measurements, at least for longer lines of sight.  However, the D/H measurements at least all suggest that ${\\rm (D/H)_{LDg}}<{\\rm (D/H)_{prim}}$ as one would expect \\citep[e.g.,][]{jll98,hwm02}.  Ultraviolet spectra of interstellar H~I and D~I Lyman-$\\alpha$ absorption lines from the {\\em Hubble Space Telescope} (HST) have provided many accurate Galactic D/H measurements \\citep[e.g.,][]{jll95,bew96,ard97,np97}.  However, the D~I Ly$\\alpha$ line merges with the H~I line for H~I column densities above $\\sim 5\\times 10^{18}$ cm$^{-2}$.  Thus, the HST results only apply for short lines of sight with low H~I column densities. In particular, the HST measurements do not reach beyond the boundaries of the Local Bubble, which is the region within $\\sim 100$~pc of the Sun in which most of the volume consists of very hot, low density, ionized ISM material \\citep[e.g.,][]{sls98,dms99,rl03}.  Analogous to the ${\\rm (D/H)_{LD}}$ and ${\\rm (D/H)_{LDg}}$ quantities defined above, we define ${\\rm (D/H)_{LB}}$ and ${\\rm (D/H)_{LBg}}$ to be the total and gas-phase D/H values for the Local Bubble, respectively. \\citet{jll98} quotes a value of ${\\rm (D/H)_{LBg}}=(1.5\\pm 0.1)\\times 10^{-5}$ based on HST measurements of the Local Interstellar Cloud (LIC) immediately surrounding the Sun.  This is about a factor of 2 lower than the ${\\rm (D/H)_{prim}}$ value quoted above.  More recent measurements from the {\\em Far Ultraviolet Spectroscopic Explorer} (FUSE) have confirmed this result \\citep{hwm02,cmo03}.  There is no convincing evidence for D/H variations within the Local Bubble, but ${\\rm (D/H)_{LBg}}$ will equal ${\\rm (D/H)_{LDg}}$ only if interstellar gas in the Galaxy is relatively homogeneous and well mixed.  Consideration of longer lines of sight is necessary to determine if ${\\rm (D/H)_{LDg}}={\\rm (D/H)_{LBg}}$.  Measuring D/H for longer lines of sight requires access to Lyman lines higher than Ly$\\alpha$.  Observations of these lines have been provided by the {\\em Copernicus} satellite in the 1970s, the Interstellar Medium Absorption Profile Spectrograph (IMAPS) instrument (part of the {\\em ORFEUS-SPAS~II} experiment onboard the STS-80 Space Shuttle Columbia flight in 1996), and more recently by FUSE, which was launched in 1999. Measurements of D/H from these instruments \\citep[e.g.,][]{dgy76,avm77,gs00,cgh03} suggest a significant amount of variability for D/H outside the Local Bubble within the range ${\\rm D/H}=(0.5-2.2)\\times 10^{-5}$ \\citep{hwm02}. When combined with the ${\\rm (D/H)_{prim}}$ measurement, this suggests a deuterium destruction factor of $\\sim 1.3-5.6$ in the local disk region of the Galaxy.  The actual average ${\\rm (D/H)_{LDg}}$ value and the average deuterium destruction factor are presumably somewhere in the middle of the above ranges, but it is not yet clear exactly where.  The spatial variability of the Galactic D/H should disappear if one considers only lines of sight longer than the scale size of the variations, so the longest and highest column density lines of sight should in principle provide the best measurements of ${\\rm (D/H)_{LDg}}$.  In practice, however, the high column densities for long Galactic lines of sight often lead to a very confusing spectrum of blended atomic and molecular absorption lines, making identification and accurate measurement of deuterium lines impossible in many cases.  Nevertheless, considering D/O and D/N measurements with previously published D/H values, \\citet{gh03} suggest that the deuterium abundance for long lines of sight is significantly lower than in the Local Bubble.  However, there are very few measurements for distances longer than 500~pc.  In this paper, we will double the number of D/H measurements in this distance regime by analyzing FUSE observations of two new lines of sight with $d>500$~pc. ",
        "conclusions": "We have analyzed FUSE observations of interstellar absorption for two long lines of sight towards the sdO stars JL~9 and LSS~1274.  Our results are summarized as follows: \\begin{description} \\item[1.] Using two separate measurement techniques, we have measured column densities for many different atomic species and for the $J=0-5$ rotational levels of H$_{2}$.  The results are listed in Table~3. \\item[2.] We find low molecular fractions of $f({\\rm H_{2}})=0.056\\pm 0.012$ and $f({\\rm H_{2}})=0.032\\pm 0.006$ towards JL~9 and LSS~1274, respectively.  The H$_{2}$ gas along these lines of sight is found to have a temperature of $T=89\\pm 6$~K for JL~9 and $T=64\\pm 5$~K for LSS~1274, both very typical values for H$_{2}$ in the Galaxy. \\item[3.] The D/H ratios for the two lines of sight are ${\\rm D/H}=(1.00\\pm 0.37)\\times 10^{-5}$ for JL~9 and ${\\rm D/H}=(0.76\\pm 0.36)\\times 10^{-5}$ for LSS~1274 (2$\\sigma$ uncertainties).  These D/H values are low compared to the Local Bubble value.  When considered with other measurements, these new results provide additional crucial evidence that long lines of sight with high column densities tend to have low gas-phase deuterium abundances. This confirms the results of \\citet{gh03} based on D/O and D/N measurements, and older published D/H values. \\item[4.] We consider our two new D/H measurements in combination with previous Galactic D/H measurements from {\\em Copernicus}, HST, IMAPS, and FUSE.  Collectively, these data suggest that D/H is constant for both low column density [$\\log N({\\rm H~I})<19.2$] and high column density [$\\log N({\\rm H~I})>20.5$] lines of sight, but variable for intermediate columns [$19.2<\\log N({\\rm H~I})<20.5$].  This suggests that regions of constant D/H in the ISM have typical column densities of $\\log N({\\rm H~I})\\sim 19$.  However, no variability is seen for lines of sight with $\\log N({\\rm H~I})>20.5$, perhaps due to the sampling of a large number of these regions for such long sight lines. \\item[5.] The low column density regime [$\\log N({\\rm H~I})<19.2$] represents the Local Bubble.  Our gas-phase Local Bubble D/H value of ${\\rm (D/H)_{LBg}}=(1.56 \\pm 0.04) \\times 10^{-5}$ is consistent with previous measurements \\citep{jll98,hwm02,jll03}. The apparent constancy of D/H at high column densities relies heavily on the two new D/H values for JL~9 and LSS~1274.  Since longer, higher column lines of sight sample more regions of the ISM, we argue that the D/H values for these longer lines of sight provide better estimates of the true gas-phase local-disk D/H ratio [${\\rm (D/H)_{LDg}}$] than the Local Bubble measurements.  For the lines of sight with $\\log N({\\rm H~I})>20.5$ we find ${\\rm (D/H)_{LDg}}=(0.85\\pm 0.09)\\times 10^{-5}$.  This is a factor of 2 lower than ${\\rm (D/H)_{LBg}}$, but is similar to the value derived by \\citet{gh03} from D/O and D/N measurements of lengthy sight lines.  (Note that errors quoted above for ${\\rm (D/H)_{LBg}}$ and ${\\rm (D/H)_{LDg}}$ are 1$\\sigma$ standard deviations of the mean.) \\item[6.] The cause of the observed D/H variability within the Galaxy is uncertain.  We discuss two possible explanations:  1. Variable astration and incomplete mixing in the ISM, and 2. Depletion of deuterium onto dust grains.  If \\#2 is correct, then the low D/H values measured for long lines of sight are due to depletion.  In this case, the total (i.e., gas plus dust) D/H ratio of the local disk, ${\\rm (D/H)_{LD}}$, is at least as high as the Local Bubble value of ${\\rm (D/H)_{LBg}}=(1.56 \\pm 0.04) \\times 10^{-5}$ instead of being at the low gas-phase value of ${\\rm (D/H)_{LDg}}=(0.85\\pm 0.09)\\times 10^{-5}$. However, if \\#1 is correct, then ${\\rm (D/H)_{LD}}={\\rm (D/H)_{LDg}}$. Deuterium destruction factors can be computed by comparing the Galactic and primordial D/H ratios.  Scenario \\#1 suggests a destruction factor of $3.3\\pm 0.6$, while scenario \\#2 suggests a value of $1.8\\pm 0.3$.  The latter is more consistent with the predictions of Galactic chemical evolution models. \\end{description}"
    },
    "0403/astro-ph0403095_arXiv.txt": {
        "abstract": "We present a detailed analytical study of the propagation of ultra high energy (UHE) particles in extragalactic magnetic fields. The crucial parameter which affects the diffuse spectrum is the separation between sources. In the case of a uniform distribution of sources with a separation between them much smaller than all characteristic propagation lengths, the diffuse spectrum of UHE particles has a {\\em universal} form, independent of the mode of propagation. This statement has a status of theorem. The proof is obtained using the particle number conservation during propagation, and also using the kinetic equation for the propagation of UHE particles. This theorem can be also proved with the help of the diffusion equation. In particular, it is shown numerically, how the diffuse fluxes converge to this universal spectrum, when the separation between sources diminishes. We study also the analytic solution of the diffusion equation in weak and strong magnetic fields with energy losses taken into account. In the case of strong magnetic fields and for a separation between sources large enough, the GZK cutoff can practically disappear, as it has been found early in numerical simulations. In practice, however, the source luminosities required are too large for this possibility. ",
        "introduction": "\\label{introduction}  The propagation of UHECR protons and nuclei with $E\\gsim 1\\times 10^{19}$~eV in the large-scale intergalactic magnetic field (IMF) remains an open problem, mainly because the knowledge of the IMF is still very poor. The possibilities vary between rectilinear propagation in a weak field and diffusive propagation in a strong magnetic field. The experimental data on IMF and the models of origin of these fields do not allow at present to choose even between the two extreme propagation regimes mentioned above.  Most reliable observations of the intergalactic magnetic field are based on the Faraday rotation of the polarized radio emission (for the reviews see Kromberg (1994), Vall\\'e (1997), Carilli and Taylor (2002)) . The upper limit on the Faraday rotation measure (RM) in the extragalactic magnetic field, obtained from the observations of distant quasars, gives an upper limit of ${\\rm RM} < 5~{\\rm rad/m}^2$. It implies an upper limit on the extragalactic magnetic field on each assumed scale of coherence length (Kromberg (1994), Vall\\`e (1997), Ryu et al. (1998)). For example, according to Blasi et al. (1999a) for an inhomogeneous universe $B_{l_c} < 4$~nG on a scale of coherence $l_c= 50$~Mpc.  According to observations of the Faraday rotations the extragalactic magnetic field is strongest, or order of $1~\\mu$G, in clusters of galaxies and radiolobes of radiogalaxies (Vall\\'e (1997), Kromberg (1994), Carilli and Taylor (2002)). The largest scale in both structures reaches $l_c \\sim 1$~ Mpc. Most probably various structures of the universe differ dramatically by magnetic fields, with very weak field in voids and much stronger in the filaments (Ryu et al. (1998)). Superclusters seem to be too young for the regular magnetic field to be formed in these structures on a large scale $l_c \\sim 10$~Mpc.  In case of hierarchical magnetic field structures in the universe, UHE protons with $E> 4\\times 10^{19}$~eV can propagate in a quasi-rectilinear regime. Scattering of UHE protons occurs mostly in galaxy clusters, radiolobes and filaments. Deflections of UHE protons can be large for some directions and small for the others. The universe looks like a leaky, worm-holed box, and correlation with the sources can be observable (see Tinyakov and Tkachev (2001)), where correlations of UHECR with BLLacs are found). Such a picture has been suggested by Berezinsky et al. (2002b).  A promising theoretical tool to predict the IMF in large scale structures is given by magneto-hydrodynamic (MHD) simulations. The main uncertainty in these simulations is related to the assumptions concerning the seed magnetic field.  The MHD simulations of Sigl et al. (2004) and Sigl et al. (2003) favor the hierarchical structure with strong magnetic fields. Assuming an inhomogeneous seed magnetic field generated by cosmic shocks through the Biermann battery mechanism, the authors obtain a $\\sim 100$~nG magnetic field in filaments and $\\sim 1$~nG in voids. In some cases they consider IMF up to a few micro Gauss as allowed. In these simulations UHECR are characterized by large deflection angles, of the order of $20^{\\circ}$, at energies up to $E\\sim 10^{20}$~eV (Sigl et al. (2003), Sigl et al. (2004)). Thus, the scenario that emerges in these simulations seems to exclude the UHECR astronomy. These simulations have some ambiguity related to the choice of magnetic field at the position of the observer (Sigl et al. (2003), Sigl et al. (2004)). The authors consider two cases: a strong local magnetic field $B\\sim 100$~nG and a weak field $B \\ll 100$~nG. The different assumptions about the local magnetic field strongly affects the conclusions about UHECR spectrum and anisotropy.  The essential step forward in MHD simulations has been made recently. In Dolag et al. (2003) the Local Universe is simulated with the observed density and velocity field. This eliminates the ambiguity for the local magnetic field, that is found to be weak. The seed magnetic field, used in this simulation, is normalized by the observed magnetic field in rich clusters of galaxies. The results of these constrained simulations indicate the weak magnetic fields in the universe of the order of $0.1$~nG in typical filaments and of $0.01$~nG in voids. The strong large-scale magnetic field, $B\\sim 10^{3}$~nG, exists in clusters of galaxies, which, however, occupy insignificant volume of the universe. The picture that emerges from simulations of Dolag at el. (2003) favors a hierarchical magnetic field structure characterized by weak magnetic fields. UHE protons with $E> 4\\times 10^{19}$~eV can propagate in a quasi-rectilinear regime, with the expected deflection angles being very small $\\le 1^{\\circ}$. However, until direct observational evidences for this picture becomes available, an alternative case of propagation in strong magnetic fields, with diffusion as an extreme possibility, can be hardly excluded.  This case has been studied in Sigl et al. (1999), Lemoine et al. (1999), Stanev (2000), Harari et al. (2002), Yoshiguchi et al. (2003), Deligny et al. (2003). An interesting features found in these calculations are small-angle clustering of UHE particles as observed by Hayashida et al. (1996), Hayashida et al. (1999), Uchiori et al. (2000), Glushkov and Pravdin (2001), and absence of the GZK cutoff in the diffusive propagation, when the magnetic field is very strong. Many aspects of diffusion of UHECR have been studied in numerical simulation by Casse et al. (2002).  We shall illustrate UHECR propagation in strong magnetic fields by the calculations by Yoshiguchi et al. (2003). The authors performed MC simulations for propagation in random magnetic field with the Kolmogorov spectrum of turbulent energy density $w_k \\propto k^n$, with $n=-5/3$. The basic (largest) coherent scale is chosen as $l_c$ with the  mean field $B_0$. Numerically these parameters vary in the range (1 - 40)~Mpc for $l_c$, and (1 - 100)~nG - for $B_0$. The sources are taken as galaxies from Optical Redshift Survey catalog with absolute magnitude $M_B$ brighter than some critical value $M_c$. The calculated quantities are the energy spectrum, anisotropy and small-angle clustering.  The observed small-angle clustering and absence of the GZK cutoff in the AGASA observations can be reproduced in the case of strong magnetic field $B \\geq 10$~nG.  Diffusive propagation of extragalactic UHECR has been studied already in earlier work. The stationary diffusion from Virgo cluster was considered by Wdowczyk and Wolfendale (1979), Giller et al. (1980) and non-stationary diffusion from a nearby source was studied by Berezinsky et al. (1990a), Blasi and Olinto (1999b) using the the Syrovatsky solution (Syrovatskii (1959)) of the diffusion equation. In this case the GZK cutoff can be absent. A very similar problem was considered again more recently by Isola et al. (2002)).  In this paper we shall study how propagation influences the diffuse energy spectrum of UHECR. We shall prove the theorem that if distance between sources  is much smaller than all propagation lengths, the spectrum has the same universal form independent of the mode of propagation. For diffusion in magnetic fields we shall demonstrate how the spectra converge to the universal one when the separation between sources diminishes. Finally, we shall obtain, with the help of the Syrovatsky solution, the spectra for strong magnetic field. In this case with large enough separation between sources, the GZK cutoff becomes weak or absent, and we will discuss the physical explanation of this phenomenon. ",
        "conclusions": "\\label{conclusions}  We have performed a formal study of the propagation of UHE particles, using an analytic approach. We demonstrated that the distance between sources is a crucial parameter which strongly affects the diffuse energy spectrum.  We have proved that, for a uniform distribution of sources, when the separation between them is much smaller than all characteristic propagation lengths, most notably the diffusion length $l_d(E)$ and the energy attenuation length $l_{\\rm att}(E)$, the diffuse spectrum of UHECR has a {\\em universal} form, independent of the mode of propagation. This statement has the status of a theorem and is valid for propagation in strong magnetic fields. The proof is given using particle number conservation during propagation and also using the kinetic equation for the propagation of UHE particles. In particular the exact solution to the kinetic equation (\\ref{kin-eq1}) for a {\\em homogeneous} ({\\em i.e.} continuous and distance-independent) distribution of the sources gives exactly the same spectrum as in the method using particle-number conservation. Note that in Eq.(\\ref{kin-eq1}) the diffusion term is absent due to homogeneous distribution of the sources.  Another proof of the theorem is given using the diffusion equation (\\ref{diff-eq}), and its exact solution (Syrovatskii (1959)) for a single source in the case of time-independent energy losses. We calculated the diffuse flux putting the sources at the vertexes of a big lattice with size $a$ and with separation $d$ between vertexes (sources). The results of the calculations are shown in Fig. \\ref{fig2}. One can see that when the distance between sources diminishes from 50~Mpc to 10~Mpc, the spectra converge to the universal spectrum (full curves) and at $d=3$~Mpc they become identical to the universal spectrum. This result is confirmed by analytic calculations. When the separation between sources is small, summation over the sources can be replaced by an integration and the corresponding spectrum (\\ref{diffuse}) coincides exactly with the universal spectrum.  In this paper we have studied the diffusive propagation of UHE particles in intergalactic space, which is considered as a turbulent magnetic plasma with baryonic gas density $n_b$ and temperature $T$, and with two basic turbulence scales $l_c$, where external energy is injected, and $l_{\\rm min}$, where turbulent energy is dissipated. Turbulent motion is considered to be random pulsations at different scales, described by an ensemble of magneto-hydrodynamic (MHD) waves $B_{\\lambda}\\exp [ i(-\\omega t + \\vec{k}\\vec{r})]$ with different frequencies $\\omega$ and different wave numbers $\\vec{k}$. Diffusion arises due to resonant scattering of particles in magnetic fields of MHD waves. At ``low'' energies, when the Larmor radius $r_L$ is much smaller than the basic scale $l_c$, the diffusion coefficient can be calculated, provided the spectral energy density of the turbulent plasma, $w(k)$, is known as a function of $k$. Such calculations for the Kolmogorov spectrum of turbulence result in the diffusion coefficient given by Eq.(\\ref{D-num}).  It is interesting to note, that the diffusion in static magnetic fields can be calculated by this method as the limiting case, when wave frequency $\\omega \\to 0$. The diffusion coefficient for static magnetic fields is given by Eq.(\\ref{static}).  The diffusion coefficient in the high energy limit, when $r_L \\gg l_c$, can be reliably calculated as the process of multiple scattering. The diffusion coefficient for this extreme case is given by Eq.(\\ref{D_asymp}). The diffusion coefficients in the low energy regime ($r_L\\ll l_c$), given by Eq.(\\ref{D-num}), and in the high energy regime ($r_L\\gg l_c$), given by Eq.(\\ref{D_asymp}), match each other well.  In practical calculations of the diffuse fluxes we characterize the magnetic configuration by three parameters ($B_0,~l_c,~d$), where $B_0$ is the mean magnetic field on the basic scale $l_c$ and $d$ is the separation of sources. We put the sources at the vertexes of a lattice of total size $a$, for which we typically used $a= 300$~Mpc. Increasing $a$ does not change the fluxes. The fluxes from the individual sources were found as the Syrovatsky solution (\\ref{syr-sol}), and the diffuse flux was found by summation over all sources, given by Eq.(\\ref{Sy-diffuse}). As a representative case for a strong magnetic field, we considered the configuration ($B_0,~l_c$)=(100~nG,~1~Mpc).  An important feature of the diffusion model is that the observed diffuse flux is produced by nearby sources (see Fig. \\ref{fig3} and \\ref{fig4}). The maximum distance $r_{\\rm max}$ depends strongly on the maximum acceleration energy $E_{\\rm max}$. For example, for $E_{\\rm max}=1\\times 10^{21}$~eV and $(B_0,l_c)=$ (100~nG,~1~Mpc), the maximum distance is less than 70~Mpc, while for $E_{\\rm max}=1\\times 10^{22}$~eV this distance is 200~Mpc. For smaller magnetic fields these distances are larger (see Fig. \\ref{fig3} right panel). The small radius $r_{\\rm max}$ of the region, which provides the dominant contribution to the observed diffuse flux of UHECR, imposes a constraint on the diffusion models, which will become more severe when/if particles with higher energies will be observed. As an example, let us consider the representative configuration (100~nG,~1~Mpc) and highest observed energy $E=3\\times 10^{20}$~eV. In this case we have $E_c=1\\times 10^{20}$~eV, $l_d(E)\\approx 10$~Mpc and $l_{\\rm att}=21$~Mpc. To avoid rectilinear propagation from nearby sources we should impose the separation $d > l_d(E)\\approx 10$~Mpc. For a maximum generation energy  of $E_{\\rm max}=1\\times 10^{21}$~eV, the maximum distance is $r_{\\rm max} < 70$~Mpc  and is only marginally consistent with $d> 10$~Mpc. The UHECR sources at $r<70$~Mpc with $d>10$~Mpc and with maximum acceleration energy $E_{\\rm max}\\sim 1\\times 10^{21}$~eV tentatively imply AGN, whose luminosities satisfy the energy requirement for the observed UHECR fluxes in the case of weak magnetic fields. However, for the case of a weak GZK cutoff (diffusion in strong magnetic fields and large separation $d$) the required luminosities are very high $L_p \\sim 10^{47}$ erg/s.  \\begin{figure}[t!] \\begin{center} \\begin{tabular}{ll} \\includegraphics[width=0.45\\textwidth]{fig8a.eps} & \\includegraphics[width=0.45\\textwidth]{fig8b.eps} \\\\ \\end{tabular} \\caption{ Comparison of the analytic  diffusive spectrum with the Monte Carlo simulation by Yoshiguchi et al. (2003). The spectra are given in the case of Kolmogorov diffusion with $(B_0,l_c)=(100~ {\\rm nG},1~{\\rm Mpc})$ and for a single source placed at distance $10$~Mpc (left panel) and $30$~Mpc (right panel). } \\label{fig8} \\end{center} \\end{figure}  The calculated diffuse energy spectra are shown in Fig. \\ref{fig5} for the Bohm  diffusion coefficient and for the complex generation energy spectrum (\\ref{complex}). In case of configuration with a very strong magnetic gield (1000~nG,~1~Mpc,~30~Mpc) the spectrum has a weak GZK cutoff (the solid curve in  Fig. \\ref{fig5}). This spectrum agrees reasonably with the AGASA excess, but requires very large source luminosity $L_p=1.5\\times 10^{47}$ erg/s. When magnetic field diminishes to 100~nG the GZK cutoff appears (dotted line), as it should according to the propagation theorem.  A weak GZK cutoff in the case of diffusive propagation is explained by the flux suppression at $E<E_{\\rm GZK}$ cutoff due to longer propagation time from the source (see Section \\ref{diff_diff} for the detailed explanation).  The spectra calculated in this paper in the diffusion approximation are compatible, for the relevant parameters, with the numerical simulations in Yoshiguchi et al. (2003). As an example we present in Fig. \\ref{fig8} our spectra from a single source calculated in the diffusion approximation for magnetic configuration ($B_0,~l_c$)=(100~nG,~1~Mpc) with the Kolmogorov spectrum of turbulence for a distance to the source of 10~Mpc (left panel) and 30~Mpc (right panel). For the distance 10~Mpc we also plot the spectrum for the rectilinear propagation shown by the dotted curve. The two curves intersect at $E_* \\sim 3\\times 10^{20}$~eV, as they should provided by the condition $l_d(E)=r$, where $r$ is the distance to the source. At $E>E_*$ the particles propagate rectilinearly  and the flux is given by the dashed line. This spectrum is compared with the numerical simulations of Yoshiguchi et al. (2003) for the same magnetic field configuration (Yoshiguchi (2004)), shown by the crosses. Since the calculations of Yoshiguchi et al. (2003) are not normalized, we equate the fluxes at energies with rectilinear propagation as shown in Fig. \\ref{fig8}. One can observe considerable disagreement at $E \\sim 8\\times 10^{19}$~eV, where it is possible to  suspect a transitional regime between quasi-rectilinear and diffusive propagation in numerical simulations. For the distance to the source r=30~Mpc (right panel) the agreement is much better, probably because at these distances the diffusion regime is reached in the numerical simulations.  The study of this paper is not intended to be a realistic one. We consider the diffusive propagation of UHE particles using an analytic approach with the aim of understanding the basic properties of propagation in strong magnetic fields. A realistic propagation should be studied within the hierarchical model of different magnetic fields in different large scale structures, as it is done in simulations (Sigl et al. (2003), Sigl et al. (2004), Dolag et al. (2003)). However, we are able to reach some practical conclusions.  Our study is focused on the  spectra of UHECR. We demonstrated that the crucial parameter is the source separation $d$. For a wide class of magnetic field configurations, when the separation $d$ is smaller than the diffusion length $l_d$, the spectrum is universal, i.e. the same as in case of rectilinear propagation. For the simulation of Dolag et al. (2003) the spectrum must be universal. The simulations of Sigl et al. (2003), Sigl et al. (2004) and Yoshiguchi et al. (2003) include diffusive and intermediate regimes. According to our calculation, in the diffusion regime very strong magnetic fields, $B_0 \\sim 1000$~nG,  and large separation between sources are needed to produce a spectrum with weak (or absent) GZK cutoff. In this case very high source luminosities are required, $L_p > 1\\times 10^{47}$~erg/s. It disfavors diffusion in strong magnetic fields as the  explanation for the AGASA excess at high energies."
    },
    "0403/astro-ph0403576_arXiv.txt": {
        "abstract": "In order to confirm the anisotropic effect of resonance photons in hot accretion columns on white dwarfs in magnetic cataclysmic variables, proposed by Terada \\etal (2001), systematic studies with {\\it ASCA} of 7 polars and 12 intermediate polars are performed. The equivalent widths of He-like Fe K$_\\alpha$ line of polars are found to be systematically modulated at their spin periods in such a way that it increases at the pole-on phase. This implies that the anisotropic mechanism is commonly operating among polars. On the other hand, those of intermediate polars are statistically consistent with being unmodulated with an upper limit of 1.5 times modulation. This may be due to a different accretion manner like aurora curtain \\citep{Rosen_88_Curtain_Model} so that the plasma becomes optically thin along the horizontal axis also for the resonance lines, or because of larger optical depths for Compton scattering if the emission regions have the same coin-like shapes as polars. ",
        "introduction": "\\label{section:introduction} A magnetic cataclysmic variable (MCV) is a close binary consisting of a white dwarf (WD) and a mass-donating companion star. Matter that spills over the Roche lobe of the companion star accretes on a magnetic pole of the WD, forming a post-shock hot plasma (Patterson 1994 and references therein). According to the systematic X-ray study of MCVs by {\\it Ginga}, their X-ray spectra can be described by optically thin thermal bremsstrahlung continua with Fe lines \\citep{Ishida_91_PhD}. These X-ray characteristics are approximately the same for both subgroups of MCVs, polars and intermediate polars (IPs). Plasma diagnostics with atomic lines, performed with {\\it ASCA} (\\cite{Fujimoto_97_asca_exhya}; \\cite{Ezuka_99}), confirmed the multi-temperature structure of the plasma as had been predicted by \\citet{Hoshi_73} and \\citet{aizu}.  Considering the typical size ($\\sim 10^{7}$ cm) and density ($\\sim 10^{15 \\mbox{--} 16}$ cm$^{-3}$) of an accretion column of an MCV, the plasma is not completely optically thin; it is estimated to be marginally optically thick for Compton (or Thomson) scattering with an optical depth of order 0.1. Under such conditions, the atomic line structure is expected to be suppressed by 10\\%--20\\%, because the line photons would experience large shifts reaching $\\sim$ 1 keV as they Compton scatter off hot electrons. At the same time, the plasma is optically thick for resonant photons of atomic lines from abundant heavy elements, including Fe in particular \\citep{Fujimoto_97_asca_exhya}. Then, the line structure of resonance transitions would be more suppressed, because the resonant photons would be effectively trapped in the column, and escape most easily when they Compton scatter out of the resonance. Nevertheless, we usually observe many ionized atomic lines in the X-ray spectra of MCVs \\citep{Ezuka_99}, and apparent metal abundances derived from the equivalent width (EW) of these lines are not much lower than one times solar. In addition, some MCVs called Pole-on line emitters (POLES; Paper I), including AX~J2315$-$0592 \\citep{Misaki_96_axj2315}, RX~J1802.1+1804 \\citep{Ishida_98_rxj1802_asca}, and AX~J1842.8$-$0423 \\citep{Terada_1999_paper1}, exhibit extremely strong He-like Fe-K lines, having an EW of $\\sim 4000$ eV, implying apparent Fe abundances of $\\sim$ three times solar.  In order to solve the puzzle of extremely strong Fe lines, we have developed a scenario of Fe-K line collimation along the vertical axis of the accretion column incorporating resonance scattering (\\cite{Terada_1999_paper1}; \\cite{Terada_2001_paper2}, hereafter Paper I; \\cite{Terada_2002_phD}). The anisotropy occurs when the accretion column has a flat shape, and is augmented by the strong longitudinal velocity gradient in the accretion flow, which invalidates the resonance condition for line photons when they propagate along the vertical axis. Since the Fe ions are heavy, their thermal resonance width (typically $\\sim3$ eV for a 10 keV plasma) is easily exceeded by the vertical Doppler shift ($\\sim10$ eV from the bottom to the top of the column). Therefore, the resonant trapping of Fe-K photons is preferentially reduced along the vertical axis, and hence the resonance line photons are collimated along the vertical direction. This scenario can explain both the lack of line destruction and the enigmatic POLE phenomenon.  We have numerically confirmed the anisotropic propagation of resonance photons via extensive Monte Carlo simulations (Paper I), employing the analytic solution to the post-shock flow by \\citet{aizu}. The results show that the directional emissivity of the resonant Fe-K line photons is indeed enhanced up to a factor of $\\sim2\\mbox{--}2.5$. In addition, we have confirmed, through {\\it ASCA} observations of the polar V834 Centauri, that the EW of its He-like Fe K$_\\alpha$ line increases during pole-on phases of its rotation (Paper I). The proposed mechanism has thus been confirmed through calculation, and observation of a few particular polars.  According to our Monte Carlo simulations, the anisotropic propagation of resonance Fe-line photons should occur for a rather wide range of plasma parameters in the accretion column (Paper I). We hence expect to observe the rotational modulation of the resonant Fe-K line EW, not only in the particular case of V834 Centauri, but also of MCVs in general. In order to examine this conjecture, here we systematically analyze the {\\it ASCA} data of 7 polars and 12 IPs. ",
        "conclusions": "\\label{section:discussion} In order to confirm the general validity of the anisotropic transfer scenario of resonance photons in MCVs pointed out in Paper I, we have systematically analyzed the {\\it ASCA} data of 7 polars and 12 intermediate polars. Through phase-resolved analysis of the ionized Fe K$_\\alpha$ lines, the EWs of the He-like Fe lines of the polars have been found to be systematically enhanced in their pole-on phases. Although the significance is only 90\\%, and the individual cases were insignificant except for V834 Centauri (Paper-I), the statistical trend is in support of our scenario. To better visualize this result, we plotted in figure \\ref{fig:polar_theta_xi} left panel the EWs of the He-like Fe lines of the polars relative to the phase averaged value, as a function of $\\theta$. There,we also plot the results of our Monte Carlo simulations based on representative parameters (see caption). Thus, the observed spin-phase dependence of the Fe-K line EWs is consistent with the calculation, although within rather large errors.  The observed He-like Fe line is in fact a blend of resonance, forbidden, and intercombitation lines, where the latter two do not suffer the resonance effect. Since the H-like line is a pure resonance line, we may expect to observe a larger spin modulation for this line compared to that of the He-like blend. According to our Monte Carlo simulations (Paper I; \\cite{Terada_2002_phD}), however, the expected enhancement of the H-like Fe line is actually smaller than that of the He-like resonance line, because the former photons are preferentially produced in a higher portion of the post-shock column, which has lower densities and smaller velocity differences than those at the bottom of the column \\citep{aizu}. As presented in figure \\ref{fig:sim_rslt_he_h_temp}, we expect almost the same enhancement of H-like Fe line as that of the He-like blend. Furthermore, the H-like Fe-K line is subject to rather large measurement errors, due to its relatively poor statistics and its possible confusion with the Fe K$_\\beta$ emission line from He-like iron. As a result, the measurements of the H-like Fe line EWs (figures \\ref{fig:polars_ew}, \\ref{fig:intermediate_polars_ew}, and \\ref{fig:polar_theta_xi} right) are consistent with the prediction.  As for IPs (figure \\ref{fig:intermediate_polars_ew}), the modulations of both line species have been found to be statistically insignificant. One possible explanation is that the accretion streams of IPs have a curtain-like shape \\citep{Rosen_88_Curtain_Model}, with a much-reduced optical depth in the lateral direction, so that the anisotropic resonance effect is suppressed. Alternatively, the emission region may have a column shape like polars, and the lack of spin modulation in the Fe line EWs may be attributed to higher electron densities $n_e$ in the plasma, as is indicated by their larger volume emission measure (VEM, described as $\\int n_e^2 dV$, where $V$ is the plasma volume) than those of polars. Then, the enhanced Compton scattering by hot electrons reduces the anisotropic effect on the resonance photons (see figure 8 in Paper I). In fact, by comparing figure \\ref{fig:polars_ew} with figure \\ref{fig:intermediate_polars_ew}, we find that the IPs exhibit systematically smaller EWs of Fe K lines than the polars.  In order to examine the latter possibility for the IPs, we plot in figure \\ref{fig:mcv_emission_measure} VEMs of the polars and IPs listed in table \\ref{tbl:target_table}, which shows that the VEMs of IPs are one-to-two orders of magnitude larger than those of polars. Then, adopting the same discussion as in section 5.2 of Paper I, we estimate that the accretion column of an IP has an order-of-magnitude higher electron density of $n_{\\rm e}\\sim10^{17}$ cm$^{-3}$, an order-of-magnitude lower column height of $h \\sim 10^6$ cm, and several times larger column radius of $r \\sim 10^7$ cm, all compared with polars. Under such conditions, the optical depth to Compton scattering is $\\sim 1.0$, whereas it is only 0.4 for polars. Therefore, the interpretation may actually work.  The Fe-K line photons provide the best diagnostics of the resonance effects, because Fe ions are massive enough for the thermal Doppler effect to fall significantly below the bulk-motion Doppler shift. Although the {\\it ASCA} data of MCVs are rather limited both in statistics for phase-resolved analyses and energy resolution, a quantum jump in this research subject may be realized by the {\\it ASTRO-E 2} satellite to be launched in 2005, with its energy resolution reaching $\\sim$ 10 eV around the Fe-K$_\\alpha$-line range for the first time.  \\bigskip Finally, we thank the members of the {\\it ASCA} team for spacecraft operation and data acquisition."
    },
    "0403/hep-ph0403043_arXiv.txt": {
        "abstract": "We study the cosmological consequences of generic scalar fields like moduli which decay only through gravitationally suppressed interactions.  We consider a new production mechanism of gravitinos from moduli decay, which might be more effective than previously known mechanisms, and calculate the final gravitino-to-entropy ratio to compare with the constraints imposed by successful big bang nucleosynthesis (BBN) etc., taking possible hadronic decays of gravitinos into account. We find the modulus mass smaller than $\\sim 10^4$ TeV is excluded. On the other hand, inflation models with high reheating temperatures $T_{R,\\rm inf} \\sim 10^{16}$ GeV can be compatible with BBN thanks to the late-time entropy production from the moduli decay if model parameters are appropriately chosen. ",
        "introduction": "\\label{sec:introduction}  One of the consequences of local supersymmetry or supergravity is the existence of gravitinos, the superpartner of gravitons, whose natural mass scale is the weak scale $\\order (0.1-10)$ TeV.  If they are not the lightest supersymmetric particle (LSP), they decay into LSP and other  high energy particles such as photons, neutrinos, quark-antiquark pairs or gluons  after/during the  big-bang nucleosynthesis (BBN) epoch. Such decay products may change the light element abundances by  changing the baryon-to-entropy, neutron-to-proton ratios, or destroying the produced elements,  which may result in a significant discrepancy between yields and observations. This is called  the ``gravitino problem''. A lot of authors have studied a variety of constraints on the decaying gravitinos from BBN~\\cite{Weinberg:zq,Krauss:1983ik,Lindley:1984bg,Khlopov:pf,Ellis:1984eq,Juszkiewicz:gg,Ellis:1984er,Audouze:be,Kawasaki:1986my,Scherrer:1987rr,Dominguez:1987,Reno:1987qw,Dimopoulos:1987fz,Ellis:1990nb,Kawasaki:1994af,Kawasaki:1994bs,Protheroe:dt,Holtmann:1998gd,Jedamzik:1999di,Kawasaki:2000qr,Kohri:2001jx,Cyburt:2002uv,Kawasaki:2004yh}~\\footnote{% Cosmological constraints on stable gravitinos are studied in Refs.~\\cite{Pagels:ke,Berezinsky:kf,Moroi:1993mb} and references therein. }. %   In the standard inflationary cosmology, gravitinos are produced by the scattering processes of thermal particles at the reheating epoch just after inflation.  The yield parameter $y_{\\g}$ of gravitinos, which is the  ratio of  their number density to entropy density, $y_{\\g}=n_{\\g}/s$, is approximately given as a function of the reheating temperature after inflation, $T_{R,\\rm inf}$, as \\cite{Kawasaki:1994af,Moroi:1995fs,Bolz:2000fu}~\\footnote{% The error of this approximation formula is  within  $\\sim 5\\%$ for $T_{R} = 10^{6}$ -- $10^{14}$ GeV, and $\\sim 25\\%$ for $T_{R} = 10^{2}$ -- $10^{19}$ GeV~\\cite{Moroi:1995fs}. }% \\begin{eqnarray} \\label{eq:y32_scattinf} y_{3/2,\\rm inf} \\simeq 1.5 \\times 10^{-12} \\left(\\frac{T_{R,\\rm inf}}{10^{10} \\gev}\\right) . \\end{eqnarray} Hence it is customary to express the constraint on their abundance imposed by BBN as that on the reheating temperature $T_{R,\\rm inf}$. For example, if gravitino mass is equal to 0.1 TeV  we find an upper bound   $T_{R,\\rm inf}\\sim 10^6$GeV,  which imposes a constraint on model building of inflation.  For example, hybrid inflation \\cite{hybrid} is difficult to reconcile with this low reheating temperature, in which the inflaton is typically coupled to gauge fields and decays rapidly.  So it is preferable to have another inflation after hybrid inflation as suggested by recent observational data \\cite{YY}.   We would like to point out, however, that these constraints on $T_{R,\\infl}$ have all been obtained under the implicit assumption that $y_{\\g, \\infl}$ remains constant until their lifetime, \\begin{eqnarray} \\label{eq:lifetime_grav} \\tau_{3/2} \\simeq 4 \\times 10^{5} \\sec \\frac1{N_{\\rm G}} \\left(\\frac{\\mgrav}{1 \\ \\tev}\\right)^{-3}, \\end{eqnarray} where we assume that  the gravitino decays into a massless gauge boson and a gaugino, $N_{\\rm G}$ is the number of the generators of the gauge group, and $m_{3/2}$ is  the gravitino mass.  In this paper we argue that both the denominator and the numerator of $y_{\\g}=n_{\\g}/s$ are subject to change between the reheating epoch after inflation and their decay time, apart from the dilution due to cosmic expansion which does not change the ratio $y_{\\g}$ itself.  We then derive more appropriate constraints imposed by BBN.   In supergravity or superstring theories there appear a number of long-lived scalar fields which decay only through gravitational interactions, such as moduli, dilatons, or Polonyi field, which are referred to collectively as the modulus $\\phi$ hereafter.  The modulus $\\phi$ starts coherent field oscillation to dissipate its energy density as the Hubble parameter $H$ becomes smaller than its mass $\\mphi$. Because its dissipation rate is smaller than that of radiation, the universe turns to be matter dominated well before the lifetime of $\\phi$, when it decays producing not only huge amount of entropy, causing what is called the moduli problem \\cite{Coughlan:1983ci,Banks:1993en,deCarlos:1993jw},  but also undesirable particles for cosmology.   So far a number of groups have studied cosmological constraints on the modulus decays depending on the properties of each decay product, {\\it e.g.}, for LSPs not to close the universe~\\cite{MYY,Kawasaki:1995cy,Moroi:1999zb}, and  for radiations  to complete thermalization~\\cite{Kawasaki:1999na}.   The effects of decaying moduli on the gravitino problem are twofold. One is that the entropy produced by their decay dilutes primordial gravitino abundance, which is a good news to relax constraints on inflation model building.  The other is that these unwanted particles may also be  produced directly by the decay of moduli $\\phi$. Including these two effects, it is not apparent for us to find the allowed region of model parameters to avoid both the gravitino problem and moduli problem. So far no one has considered this type of scenario with decaying moduli. Therefore, we comprehensively study the effects in this work.  The constraint by gravitinos produced directly by the decay of moduli was first studied by Hashimoto et  al.\\ \\cite{Hashimoto:1998mu}, who considered the case that the  modulus decays into two gravitinos. In the present paper, however, we point out a more efficient mechanism of direct gravitino production from moduli decay, which is the mode that $\\phi$ decays into its superpartner modulino, $\\tilde{\\phi}$, and a gravitino, $\\psi_{\\g}$, namely, $\\phi \\to \\tilde{\\phi} + \\psi_{3/2}$. We incorporate this decay mode and consider cosmological constraints on the masses of gravitinos and moduli from BBN and other observations.  The fate of the decaying moduli depends on  its mass $m_{\\phi}$. In \\cite{Banks:1993en,deCarlos:1993jw}, it was presumed that moduli fields acquire masses through supersymmetry breaking and then their masses are comparable to masses of superparticles. In this case the modulus field is long-lived and the reheating at the modulus decay takes place with the temperature much below 1MeV, spoiling the success of the BBN. (See, however, \\cite{MYY}).  Recently it was realized that a mechanism to stabilize the moduli fields is operative in the compactification with non-zero NS and RR fluxes in  certain string theories~\\cite{Giddings}. This makes most of the moduli fields very  massive, typically around the string scale.  Still there are  some moduli which are not stabilized. In particular, the modulus field which is responsible to determine the size of the compactification is not stabilized in the flux compactification. With the ignorance of possible mechanisms on mass generation, we take the mass of the modulus field as a free parameter in this paper. ",
        "conclusions": "\\label{sec:conclusion}  We have studied the effects of decaying modulus oscillation on the cosmological gravitino problem.  We have considered a new direct production mechanism of gravitinos from modulus decay, namely, a decay mode of modulus into a gravitino and a modulino.  The width of this decay mode can be larger than the other mode into two gravitinos which has been studied in \\cite{Hashimoto:1998mu}, if the coupling constant is of the same order of magnitude with $\\acoup \\gtrsim 1$.   Comparing our yield of gravitinos with the constraints imposed by BBN and the relic LSPs, which are decay products of gravitinos, we have obtained a constraint on the masses of gravitinos and modulus.  As a result we have found that due to the above-mentioned direct production of gravitinos from decaying modulus, the modulus mass with $\\mphi < 10^4$TeV is excluded, even when the branching ratio into hadrons is minimal.  On the other hand, we have also found that wide range of $\\mgrav$ and $\\mphi$ are still allowed even if the reheating temperature after inflation is as high as $T_{R,\\rm inf}=10^{16}$ GeV and the effects on the hadronic decay of the gravitinos are taken into account, thanks to the dilution of primordial gravitinos due to the entropy production associated with modulus decay.  Thus in order to study cosmological consequences of gravitinos, it is important to analyze not only their abundance right after inflation but also their subsequent dilution due to late-time entropy production, as well as late-time production from scalar condensates with only gravitationally suppressed interactions including a dilaton and a Polonyi field."
    },
    "0403/nucl-th0403056_arXiv.txt": {
        "abstract": "{In this work, we analyze the effect of charge in compact stars considering the limit of the maximum amount of charge they can hold. We find that the global balance of the forces allows a huge charge ($\\sim 10^{20}$ Coulomb) to be present in a neutron star producing a very high  electric field ($\\sim 10^{21}$ V/m). We have studied the particular case of a polytropic equation of state and assumed that the charge distribution is proportional to the mass density. The charged stars have large mass and radius as we should expect due to the effect of the repulsive Coulomb force with the M/R ratio increasing with charge. In the limit of the maximum charge the mass goes up to $\\sim 10$ M$_\\odot$ which is much higher than the maximum mass allowed for a neutral compact star. However, the local effect of the forces experienced by a single charged particle, makes it to discharge quickly. This creates a global force imbalance and the system collapses to a charged black hole.} ",
        "introduction": "In 1924, Rosseland\\cite{ross24} studied the possibility of a self  gravitating  star  on   Eddington's  theory  to  contain  a  net charge where the star  is modeled by a  ball of hot ionized gas (see also Eddington \\cite{ed26}).  In such a system  the electrons  (lighter particles)  tend  to rise  to the  top because  of  the  difference  in  the partial  pressure  of  electrons compared to that of ions (heavier particles).  The motion of electrons to the top and further escape from the star is stopped by the electric field created  by the charge  separation. The equilibrium  is attained after some  amount of electrons escaped leaving  behind an electrified star  whose net positive  charge is  of about  100 Coulomb  per solar mass, and building an interstellar  gas with a net negative charge. As shown by Bally  and Harrison \\cite{bally78}, this result  applies to any bound system whose size is smaller  than the Debye length of the surrounding media. The  conclusion is that a  star formed by  an initially neutral gas cannot acquire a net  electric charge larger than about $100$C per solar mass. It is expected that the sun holds some amount of net charge due to the much more frequent escape of electrons than that  of protons.  Moreover, it is also expected that  the escape would stop when the electrostatic  energy of an electron $e\\Phi$ is of the order  of its thermal  energy $kT$. This  gives for a ball  of hot matter with the  sun radius, a net charge  $Q \\sim 6.7\\times 10^{-6}T$ (in Coulomb). Hence, the escape  effect cannot lead to a net electric charge much larger than a few hundred Coulomb for most of the gaseous stars.  For Newtonian stars, the net charge of 100 C per solar mass is obtained by the balance between the  electrostatic  energy  $eQ/r$ and the gravitational  energy $mM/r$ (Glendenning\\cite{glen00}). However, for very compact stars, the high density and the relativistic effects  must be taken into account \\cite{bek71}.  In  a strong  gravitational  field, the general relativistic effects  are felt and the star  needs more charge to be in equilibrium. Moreover, for  very compact stars,  the induced electric field  can be substantially higher  than in the case  of the sun.  For instance, the same amount of charge yields an electric field approximately $10^{9}$ times larger at  the surface of a neutron star than  at the surface of the sun.  So, even a relatively  small amount of  net charge on compact  stars can induce  intense electric  fields whose  effects may become  important to  the structure  of the  star. This  fact deserves further investigation.   The general relativistic analog for charged dust stars was discovered by Majumdar \\cite{maj47} and by Papapetrou \\cite{papa47}, and further discussed by Bonnor \\cite{bon75} and several other authors\\cite{ivanov02c}. Study for  the stability of charged fluid  spheres have been  done by Bekenstein\\cite{bek71}, Zhang et al.\\cite{zhang82},  de Felice \\& Yu\\cite{fel95}, Yu \\& Liu\\cite{yu00},  de Felice  et al.\\cite{fel99},  Anninos \\& Rothman\\cite{ani01} and others. This  was indirectly  verified by Zhang et al.\\cite{zhang82} who found  that the structure of a neutron star, for a degenerate relativistic fermi gas, is significantly  affected by the electric charge  just when  the charge density  is close to  the mass  density (in  geometric units).  In the investigations by de  Felice et al., and by  Anninos \\& Rothman, they assumed that the charge  distribution followed particular functions of the radial coordinate, and they  were mostly interested in the extreme $Q=M$ case.  Our basic consideration to incorporate charge into  the system is in the form of trapped  charged particles  where the  charge goes  with  the positive value. The effect of charge does not depend on its sign by our formulation. The energy density  which appears from the electrostatic field will {\\it add up} to the total energy density of the system, which in turn  will  help  in the  {\\it  gaining}  of  the  total mass  of  the system. The modified Tolman-Oppenheimer-Volkoff (TOV) equation now has extra terms due to the presence of  the Maxwell-Einstein  stress tensor. We solve the modified TOV equation for  polytropic equation of state (EOS) assuming  that the  charge density goes with the  matter density and discuss the  results. The formation of this extra charge inside the star is however left open. A mechanism to generate charge asymmetry for charged black holes has been suggested recently by Mosquera Cuesta et al. \\cite{herman03} and the same may be applied for compact stars too .  This article is arranged in the following way. In Section \\ref{sec:rel-form}, we show the basic formalism for the modified TOV. In Section \\ref{sec:poly}, we used this modified TOV on a polytropic EOS, discuss the results and the stability of the charged stars. Finally we make our conclusions in Section \\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc} In our study, we have shown that a high density system like a neutron star can hold huge charge of the order of 10$^{20}$ Coulomb considering the global balance of forces. With the increase of charge, the maximum mass of the star recedes back to a lower density regime. The stellar mass also increases rapidly in the critical limit of the maximum charge content, the systems can hold. The radius also increases accordingly, however keeping the M/R ratio increasing with charge. The increase in mass is primarily brought in by the softening of the pressure gradient due to the presence of a Coulombian term coupled with the Gravitational matter part. Another intrinsic increase in the mass term comes through the addition of the electric energy density to the mass density of the system.  The inside electric field of the charged stars are very high and crosses the critical field limit for pair creation (Bekenstein \\cite{bek71}). However, this issue is debatable because the critical field has been calculated for vacuum and one does not really know what the value will be in a high density system. The stability of the charged stars are however ruled out from the consideration of forces acting on individual charged particles. They face enormous radial repulsive force and leave the star in a very short time. This creates an imbalance of forces and the gravitational force overwhelms the repulsive Coulomb and fluid pressure forces and the star collapses to a charged black hole.  Finally, these charged stars are supposed to be very short lived, and are the intermediate state between a supernova collapse and charged black holes.  SR acknowledges the FAPERJ for research support, MM for the partial CNPQ support, and JSPL and VTZ for the hospitality of Observat\\'{o}rio Nacional, Rio de Janeiro.  \\small"
    },
    "0403/astro-ph0403338_arXiv.txt": {
        "abstract": "During the APPLES parallel campaign, the HST Advanced Camera for Surveys has resolved a distant stellar system, which appears to be an isolated dwarf galaxy. It is characterized by a circularly symmetric distribution of stars with an integrated magnitude $m_{F775W}$ = 20.13 $\\pm$ 0.02, a central surface brightness $\\mu_{F775W} \\simeq$ 21.33 $\\pm$ 0.18 mag~arcsec$^{-2}$ and a half-light radius of $\\simeq$ 1.8 arcsec. The ACS and VLT spectra show no evidence of ionized gas  and appear dominated by a 3 Gyr old stellar population. The OB spectral type derived for two resolved stars in the grism data and the systemic radial velocity (V$_{hel} \\simeq$ 670 km~s$^{-1}$) measured from the VLT data give a {\\it fiducial} distance of $\\simeq$ 9 $\\pm$ 2 Mpc. These findings, with the support of the spatial morphology, would classify the system among the dwarf spheroidal (dSph) galaxies. Following the IAU rules, we have named this newly discovered galaxy APPLES~1. An intriguing peculiarity of APPLES~1 is that the properties (age and metallicity) of the stellar content so far detected are similar to those of dSph galaxies in the Local Group, where star formation is thought to be driven by galaxy interactions and mergers. Yet, APPLES~1 seems not to be associated with a major group or cluster of galaxies. Therefore, APPLES~1 could be the first example of a field dSph galaxy with self-sustained and regulated star formation and, therefore, would make an interesting test case for studies of the formation and evolution of unperturbed dSph galaxies. ",
        "introduction": "In the hierarchical framework dwarf galaxies are considered to be the building blocks of larger baryonic structures, and observations indeed show that they are the most abundant galaxies in the Universe. The SDSS (York 2000) and 2dFGRS (Colless et al. 2001) surveys have shown that the galaxy luminosity function keeps rising down to the faintest magnitude probed (Blanton et al. 2001, Norberg et al. 2002), with the faint end of the galaxy luminosity function depending on environment. Hoyle et al. (2003) have found that SDSS dwarf galaxies fainter than M$_r \\simeq$ -15  are not generally present in {\\it voids} but preferentially populate {\\it overdense regions} (i.e. clusters; cf. Mo et al. 2004, Sabatini et al. 2003). \\par\\noindent Cold Dark Matter (CDM) theories predict mass functions for subhaloes with a very steep slope at the low mass end, virtually independent of the mass of the parent halo (Kravtsov et al. 2004a). While this is true for the Virgo cluster (Sandage et al. 1985, Moore et al. 1999), it is dramatically inconsistent with the observed galaxy content of the Local Group in which far fewer galaxies have been detected than one would na\\\"ively expect for CDM cosmogonies (Klypin et al. 1999; Moore et al. 1999). A number of mechanisms for suppressing dwarf galaxy formation have been included in CDM theories, such as supernova-driven winds (Dekel \\& Silk 1986) and reionization (Bullock et al. 2000, Tully et al. 2002, Kravtsov et al. 2004b). These authors have suggested that reionization could shape the environment dependence of the faint end of the galaxy luminosity function, so that galaxy clusters assembled before reionization would be dwarf-galaxy rich in contrast with clusters formed after reionization (the ``squelching'' effect).  Counts of dwarf galaxies in galaxy groups and clusters indicate that they constitute the majority of the in-situ population, e.g. 80$\\%$ in the Virgo cluster (Binggeli et al. 1985) and 90$\\%$ in the Local Group (Mateo 1998). The best studied dwarf galaxies are in the Local Group, where they have been classified (cf. Ferguson \\& Binggeli 1994) into dwarf ellipticals (dE: with a compact morphology and a high central surface brightness), spheroidals (dSph: more diffuse structures with a low central surface brightness) and irregulars (dIrr: with an irregular morphology and a significant content of ionized and neutral gas). dEs and dSphs are ``gas-free'' systems (Mateo 1998). The differences observed in the kinematic properties and chemical abundances (the latter being inferred from  the well-known luminosity - metallicity relation for dwarf galaxies where dIrrs are too metal-poor for their luminosity as compared to dSphs) rule out that dIrrs naturally evolve into dSphs and dEs (Binggeli 1994, Grebel et al. 2003). Such a transition is rather induced by the environment, via tidal interaction with a closeby, more massive galaxy and via ram-pressure with the intergalactic medium in the Local Group, with the final effect of depleting dIrrs of their original gas. Indeed, observations indicate that dwarf ellipticals outnumber dwarf irregulars in the Local Group (Mateo 1998), in Hickson compact groups (Zepf et al. 1997) and in the Virgo cluster (Sandage et al. 1985), while dwarf irregulars in the field dominate the galaxy luminosity function at M$_B <$ -18 (Lilly et al. 1996).  Dwarf galaxies in the Local Group (as well as dwarf galaxies studied so far in the M81 group, Caldwell et al. 1998) exhibit complicated star-formation histories and are characterized by a mixture of old- and intermediate-age stars (dEs and dSphs) at different metallicities, sometimes with traces of a more recent starbust (few Myrs) like in dIrrs (Grebel 1997, Mateo 1998, Tolstoy 2003, Grebel et al. 2003). Although van den Bergh (1994) suggested that dSph galaxies nearest to the Milky Way are on average older than more distant ones, there is no solid evidence for their star-formation history to be dependent on distance from the Galaxy or M31 (Mateo 1998). Similarly, there is no clear understanding of the mechanisms driving star formation in the Local Group dwarfs. Galactic outflows induced by supernova explosions, ram-pressure stripping, tidal stripping, galaxy threshing, galaxy harassment and mergers are all thought to play a role in controlling the star-formation activity and gas content in satellite dwarf galaxies (Grebel 1997, Moore et al. 1999, Bekki et al. 2001). There is no doubt that the dynamical evolution of a galaxy group ``washes out'', in a sense, the pristine properties of the group dwarf galaxies. Galaxy harassment in groups and clusters of galaxies can disturb large disk galaxies and transform them into dwarfs (Moore et al. 1999, Gnedin 2003a,b). Galaxy interactions in compact groups are even able to produce an ex-novo generation of dwarf galaxies, the so called tidal debris dwarf galaxies (cf. Palma et al. 2002).  Have the truly pristine galaxy building blocks been found yet? These are the unperturbed dwarf galaxies which most resemble the galaxy building blocks at high redshifts. In the nearby Universe, Tucana and Cetus, two dSph galaxies, are located at about 880 kpc from the Milky Way and are considered practically unbound to the Local Group. Unlike the satellite dSphs, they do not show any recent burst of star formation and their stellar populations date back to $\\sim$ 10 Gyr ago (Mateo 1998, Grebel 1997). Therefore, they may well represent primordial dwarf galaxies, which have evolved as predicted by Dekel \\& Silk (1986), i.e. dSphs descend from gaseous systems which lost their gas because of a single burst of star formation.  In this paper we report on the first detection of a gas-poor dwarf galaxy which has undergone recent star formation and, yet, likely lies in the field, far away from any overdensity environment. We argue that this new galaxy, hereafter referred to as APPLES~1, may be located at a fiducial distance of $\\sim$ 9 Mpc and with an absolute magnitude of M$_{F775W} \\simeq$ -9.6. This would make APPLES~1 possibly one of the faintest dwarf galaxies resolved in the field. Its case would then set important constrains on feedback and on how star formation occurs in the smallest size scales of dark matter.  The discussion on the nature of APPLES~1 develops following the timeline along which more data became available for this galaxy: first, the survey images alone which allowed us to study the morphology and structure of APPLES~1, and finally the ACS and VLT spectra which allowed us to estimate the distance and stellar content of this galaxy. ",
        "conclusions": "Among  the data  collected for  the APPLES  Parallel Program  with the HST/ACS  Wide Field  Channel, we  have discovered  a  resolved stellar system whose properties  are consistent with the object  being a dwarf galaxy. At a  distance of $\\simeq$ 9 $\\pm$ 2 Mpc (based  on its radial velocity and the spectral type of its resolved sources) APPLES  1 seems to belong to  the field  population  of galaxies and not associated with any major cluster or group of galaxies.  In the  available ACS  image, APPLES~1  looks quite like  a circularly symmetric  distribution of  stars, with a  half-light  radius of $\\simeq$ 1$''.$8 or 78 pc (at the fiducial distance of $\\sim$ 9 Mpc) and  a total $m_{F775W}$  magnitude of  20.1, which  translates into M$_{F775W} \\simeq$ -9.6 ($\\pm$ 0.5 due to the uncertainty on the distance). At face value, the size  of APPLES~1 could remind  of  the  ultra-compact  dwarf  galaxies  (UCD)  discovered  by Drinkwater et al. (2003) in the Fornax cluster. Unlike APPLES~1, these dwarf galaxies are dominated by  old stellar populations (Gregg et al. 2003) and are believed to be the result of galaxy threshing as modeled by  Bekki  et  al.   (2001),  whereby nucleated  galaxies  lose  their envelope  and  are reduced  to  their  only  nucleus by  gravitational interactions with a more massive galaxy.   The observed  surface brightness  profile  of APPLES  1 is  well fitted by a Sersic profile with a scale length of $\\simeq$ 0$''.$35, a central surface brightness  of $\\simeq$ 21.3 mag~arcsec$^{-2}$ and a Sersic index  $n$ = 2.3. Stiavelli et al.   (2001) have shown that the Sersic  indices computed for  25 dwarf ellipticals in  the Virgo and Fornax clusters range between  0.62 and 4. The authors have also pointed out  that the  best fitted galaxies  are characterized  by a slope intermediate between  $n$ = 1 and $n$ = 4,  i.e. a range which well matches  the Sersic index  of APPLES~1.  The  non-detection of ionized gas  (i.e.  H$\\beta$,  [OIII], H$\\alpha$) in  the ACS  and VLT spectra  together  with  the  spatial  morphology  would  likely classify APPLES~1 among the dSph galaxies.  Inquiries of the HIPASS database show  that no HI  gas has been  detected at the  position and radial velocity (V$_{hel}$ = 674 km~s$^{-1}$) of APPLES~1. At the distance  of APPLES~1,  HIPASS turns  out to  have a  mass detection limit of $\\simeq$ 10$^8$ M$_o$. So far, accurate HI masses have been measured  only for  nearby dwarf  galaxies; for  example,  the dwarf ellipticals NGC  185 and NGC 205  have a HI content  of about 10$^5$ M$_o$ (Young  \\& Lo 1997),  while the HI  mass in the  dSph galaxies Sculptor  and  Phoenix  ranges   between  10$^3$  and  10$^5$  M$_o$ (Carignan et al. 1998, St-Germain et al. 1999).   The  comparison  of the  Lick  line indices  as  measured  in the  VLT spectrum  with the  stellar population  models by  Bruzual  \\& Charlot (2003) gives for the core of APPLES~1 a luminosity-weighted age of $1-5$~Gyr.  We have also estimated the age and metallicity of the APPLES~1 core  by fitting  the continuum  slope of  its VLT  spectrum  with the chemical enrichment models  by Ferreras \\& Silk (2000,  2003). In this way, we have derived a  luminosity-weighted age of $\\simeq$ 3 Gyr with a burst  duration of $\\sim$  2 Gyr, a  luminosity-weighted metallicity $\\log Z/Z_\\odot=-0.2$~dex with a  metallicity spread of $\\simeq 0.2$~dex. The best  fitting model gives  for the  stellar population  dominating the optical  a mass-to-ligh  ratio M/L$_V$  = 1.4  and a  stellar  mass of $\\simeq 8\\times 10^{6}$ M$_{\\odot}$.  The infall of  gas, which feeds the ongoing burst, has apparently started at a redshift $z_F\\sim 3$ and has a  timescale of about 5~Gyr.  The efficiency at which  the core of APPLES~1  forms stars  is about 0.7~Gyr$^{-1}$, typical  of late-type galaxies. The presence of OB stars and an integrated spectra which matches an F type star argues against a simple stellar population, whose best fit requires a very young age ($\\lesssim 1$~Gyr) and supersolar metallicities, which are in conflict with the weakness of the \\ion{Ca}{2} $H$ and $K$ lines, and with the relatively weak Mgb feature. Hence, the spectra strongly suggests an extended period of star formation over several Gyr.    As  shown in Mateo (1998)  and Grebel (1997),  the star formation history of dSphs  in the Local Group is  quite complicated. Five out of  23  dSphs listed  in  Mateo (1998)  present  one  burst of  star formation, now observed in the  form of RGB stars.  The remaining 18 dSphs  are characterized  by multiple  bursts, seven  of  which show main-sequence stars with  age = 0 Gyr. Bursts  like the one observed in APPLES~1  (i.e. dominated by a 3 Gyr  old stellar population) are known to  occur in a  number of dSph  and dIrr/dSph galaxies  in the Local Group: NGC~205, Fornax,  Carina, Sextans, DDO~210 and Pegasus. In particular,  the VLT spectrum  of the core  of APPLES~1  is quite similar to  that of the  core of NGC~205  (Bica et al. 1990)  with a stellar  population  comparably  metal-rich  and  young.  Therefore, APPLES~1  would be  consistent with the  dSph galaxies in  the Local Group. But there are four important differences with the Local Group dSphs: \\par\\noindent {\\it i)} APPLES~1 is apparently an isolated system, \\par\\noindent {\\it ii)} two nearby dSphs, Tucana and Cetus, very likely unbound to the Local Group, do not show any recent burst of star formation and their stellar populations date back to $\\sim$ 10 Gyr ago (Mateo 1998, Grebel 1997). \\par\\noindent {\\it iii)} NGC~205 probably was a disc galaxy and its interaction with M31 has removed its gas with the net effect of ``shrinking'' NGC~205 to its bulge. \\par\\noindent {\\it iv)} no old stellar component has been detected in the data currently available for APPLES~1. \\par\\noindent Tidal interactions, ram-pressure, collisions, merger events and ionization from the UV background are thought to regulate the depletion of gas (primordial and ejected by stars during their evolution) and the star formation in satellite dSphs (cf. Grebel et al. 2003, Dong et al. 2003). Therefore, Grebel et al. (2003) suggest that Tucana and Cetus could be runaways or have been stripped of their gas by ram-pressure along their orbit if the intergalactic medium of the Local Group were clumpy. Alternatively, these dwarf galaxies could have been more massive in the past and had been heated and truncated by the tidal interactions with the massive neighbors at high redshift, as suggested by Kravtsov et al. (2004b). Certainly, a better knowledge of their orbital parameters would greatly help in understanding their origin. If these ruled out the occurrence of any major interactions with the Milky Way, then Tucana and Cetus would have evolved as predicted by Dekel \\& Silk (1986) with supernova explosions removing their gas. The key point, here, is to determine the full star formation history experienced by APPLES~1. For example, the best fitting model of chemical enrichment (as well as the Lick line indices) indicates that the VLT spectrum of the core of APPLES~1 is dominated by young and intermediate-age stars with a negligible contribution from old stars to the optical light. This does not rule out that old stars, if present, could contribute significantly to the total mass of APPLES~1. Is there a $\\sim$ 10 Grys old population in APPLES~1 as typically observed in dwarf galaxies in the Local Group, or is APPLES~1 experiencing the very first burst of star formation similarly to GR~8, a dIrr galaxy in the Local Group with a stellar population younger than 1 Gyr (Mateo 1998)? In this second hypothesis, APPLES~1 would resemble a dwarf galaxy in the predictions of Dekel \\& Silk (1986) which formed out of the intergalactic medium (IGM) at $z$ = 3 and has lost its gas through stellar winds after a single burst of star formation. The metallicity derived for APPLES~1, though, would require a significantly enriched IGM, whereas the intergalactic metallicity measured at $z \\simeq$ 3 is about -1 dex (Pettini 2004). Significantly deeper and multi-wavelength imaging with HST is necessary to detect the older and fainter stars to re-construct the full star-formation history of APPLES~1. The same observations will also provide a more accurate measure of the distance of this galaxy. If these observations confirm the distance estimated for APPLES~1 in this paper and reveal the presence/absence of older stellar components, they will set the challenge of explaining what regulates star formation in a isolated dSph galaxy."
    },
    "0403/astro-ph0403612_arXiv.txt": {
        "abstract": "This paper summarizes the results of 13 years of timing observations of a unique binary pulsar, PSR B1259$-$63, which has a massive B2e star companion. The data span encompasses four complete orbits and includes the periastron passages in 1990, 1994, 1997 and 2000. Changes in dispersion measure occurring around the 1994, 1997 and 2000 periastrons are measured and accounted for in the timing analysis. There is good evidence for a small glitch in the pulsar period in 1997 August, not long after the 1997 periastron, and a significant frequency second derivative indicating timing noise.  We find that spin-orbit coupling with secular changes in periastron longitude and projected semi-major axis ($x$) cannot account for the observed period variations over the whole data set. While fitting the data fairly well, changes in pulsar period parameters at each periastron seem ruled out both by X-ray observations and by the large apparent changes in pulsar frequency derivative. Essentially all of the systematic period variations are accounted for by a model consisting of the 1997 August glitch and step changes in $x$ at each periastron.  These changes must be due to changes in the orbit inclination, but we can find no plausible mechanism to account for them. It is possible that timing noise may mask the actual changes in orbital parameters at each periastron, but the good fit to the data of the $x$ step-change model suggests that short-term timing noise is not significant. ",
        "introduction": "PSR B1259$-$63 was discovered in a large-scale high frequency survey of the Galactic plane \\cite{jlm+92}. It is unique because it is the only known radio pulsar in orbit about a massive, main-sequence, B2e star \\cite{jml+92}. PSR B1259$-$63 has a short spin period of $\\sim$48~ms and moderate period derivative of $2.28\\times10^{-15}$, implying a characteristic age of only 330~kyr from spin-down by magnetic dipole radiation. It has an orbital period of $\\sim$1237 days, and an eccentricity of 0.87, the longest orbital period and largest eccentricity of any of the known binary radio pulsars. The inclination of the orbit to the plane of the sky is $i\\sim36\\arcdeg$.  The companion, SS 2883, is a 10th magnitude star with a mass of about 10~$\\msolar$ and a radius of 6~$\\rsolar$. Typical of Be stars as a class, it has a hot, tenuous polar wind and a cooler, high density, equatorial disk. The mass loss rate of the B2e star is $\\sim 10^{-6}$~$\\msolar$.  Johnston et al. (1994)\\nocite{jml+94} observed H$\\alpha$ emission lines from the disk at 20 stellar radii (R$_{*}$), just inside the pulsar orbit of 24 R$_{*}$ at periastron. The density of the disk material is high near the stellar surface ($10^{8}-10^{10}$~cm$^{-3}$), and falls off as a power-law with distance from the star. The disk is likely to be highly tilted with respect to the pulsar orbital plane, and PSR B1259$-$63 is eclipsed for about 40 days as it goes behind the disk.  Assuming the B2e star is rotating at $\\sim$70 per cent of its break up velocity \\cite{por96}, SS 2883 then has an equatorial velocity $\\sim$280~kms$^{-1}$.  The spin-induced oblateness of the star implies an additional $1/r^3$ gravitational potential term in the interaction with the pulsar, known as quadrupole gravitational moment in typical binary system. This effect introduces an apsidal motion and precession of the orbital plane if the spin of the companion is not aligned with the orbit angular momentum, characterised by $\\dot \\omega$ and $\\dot x$ respectively, where $\\omega$ is the longitude of periastron passage and $x$ is the projected pulsar semi-major axis \\cite{lbk95}.  In an eccentric binary system, the passage of the pulsar through periastron excites tidal motions on the companion star which then interact with the orbital motion. In the case of PSR J0045$-$7319 \\cite{kbm+96}, tidal effects may account for the observed evolution in binary parameters (e.g. Lai 1997)\\nocite{lai97}. The tidal interactions may be enhanced by large factors if there is a resonance with oscillation modes of the star \\cite{ws99}. For PSR J0045$-$7319, at periastron the pulsar is only $\\sim 4$ R$_{*}$ from the companion compared to 24 R$_{*}$ for PSR B1259$-$63. Tidal effects are strongly dependent on the separation of the two stars, and so may be insignificant in PSR B1259$-$63.   Interactions between the pulsar and the disk of the Be star may also be important. If there is accretion from the disk on to the neutron star, the pulsar may be spun up or slowed down. For accretion to occur, the bow shock, where there is pressure balance between the pulsar wind and the Be-star disk gas, must lie inside the accretion radius, which depends on relative velocity of the pulsar and the accreting gas. For PSR B1259$-$63, accretion is unlikely \\cite{mjl+95,ktn+95,ta97}, with X-ray observations \\cite{hck+99} most easily interpreted as originating from a bow shock well outside the pulsar magnetosphere. Disk interactions may also affect pulsar orbit directly through frictional drag.  Estimates \\cite{mjl+95,wjm+98} of variations in the Keplerian parameters $P_b$, $\\omega$ and $e$ suggest the effect is negligible.  Timing models for binary pulsars have evolved over the years, mainly to take into account general relativistic effects \\cite{bt76,dd86}. Wex (1998)\\nocite{wex98} extended these models to deal with long-period systems with main sequence companions. His model, the so-called MSS model, accounts for short-term periodic orbital effects and for the long-term secular effects caused by classical spin-orbit coupling.  The latest paper on timing of PSR~B1259$-$63 \\cite{wjm+98} analysed observations from 1990 January to 1996 December which covered the 1990 and 1994 periastron passages. Wex et al. (1998) were unable to derive a unique timing model for the system, partly because of the large timing noise and partly due to the long eclipses of the pulsed emission.  They put forward two possible solutions for the timing. In one, timing noise dominated, its effects were removed through higher order derivatives of the spin frequency, and there was no orbital precession or tidal effects.  In the second model, which they preferred, significant values of $\\dot \\omega$ and $\\dot x$ were obtained. These were interpreted as being due to orbital precession, and a large tilt between the rotation axis of the Be star and the orbital angular momentum vector was inferred. Two different fits were obtained for $\\dot \\omega$ and $\\dot x$ depending on the number of (integer) rotations added during the 1994 eclipse. The preferred solution then relied on the physical interpretation of the values of $\\dot \\omega$ and $\\dot x$.  Since then, we have obtained a further six years of timing data on this pulsar to give a total data span of 13 years, including two further periastron passages in 1997 and 2000.  Our main findings are: (1) that there is strong evidence for a pulsar period glitch in 1997 August, about 94 days after the 1997 periastron, (2) that the Wex et al. (1998) timing solution which included orbital precession does not remain valid through the 1997 and 2000 periastrons, (3) that a solution with jumps in $\\nu$ and $\\dot \\nu$ at each periastron, while fitting the data reasonably well, is unlikely to be physically viable, and (4) that the data are best explained by a combination of the 1997 August pulsar glitch and steps in projected orbit semi-major axis ($x$) at each periastron.  However, the possible presence of timing noise and the long-duration eclipses conspire to make timing of the pulsar difficult. In Section~\\ref{sec:obs} we discuss details of the observations and data analysis techniques, Section~\\ref{sec:prf} describes the pulse profile evolution from 0.66 to 13.6~GHz and Section~\\ref{sec:ddm} describes the dispersion measure variations observed around the 1994, 1997 and 2000 periastrons. The timing results are described in Section~\\ref{sec:tim} with fits of the various possible models, and the possible interpretations are discussed in Section~\\ref{sec:dis}. ",
        "conclusions": "We have presented an analysis of 13 years of timing data encompassing four periastron passages of PSR B1259$-$63. Significant changes in dispersion measure occurring near each periastron were accounted for in the data analysis. It is clear that secular changes in orbit parameters due to spin-orbit coupling does not give a good description of the data. The effect of timing noise is difficult to quantify because of coupling with the orbital parameters, a problem exacerbated by the long-duration eclipses which occur at each periastron. Consequently, we cannot rule out the possibility that our measured parameters are affected by timing noise. However, excepting the glitch which occurred near MJD 50691, the fact that within each orbit the observed residuals show variations with characteristic timescale of the same order as the orbital period and significant shorter-term variations only near the periastrons strongly suggests that changes in orbital parameters at periastron are the dominant effect. This is reinforced by the excellent fit to the data by a model containing the MJD 50691 glitch and steps in projected semi-major axis at each periastron. It is clear that these steps must result from step changes in the orbit inclination. However, we are unable to offer any plausible mechanism to account for the observed changes."
    },
    "0403/hep-th0403069_arXiv.txt": {
        "abstract": "We present a new mechanism for creating the observed cosmic matter-antimatter asymmetry  which satisfies all three Sakharov conditions from one common thread, gravitational waves.   We generate lepton number through the gravitational anomaly in the lepton number current.  The source term comes from elliptically polarized gravity waves that are produced during inflation if the inflaton field contains a CP-odd component.  In simple inflationary scenarios, the generated matter asymmetry is very small.  We describe some special conditions in which our mechanism can give a matter asymmetry of realistic size. ",
        "introduction": "As far into the universe as we can see, there is an excess of matter over anti-matter.  The recent determination of the cosmological parameters from the cosmic microwave background by the WMAP experiment gives the baryon density of the universe as \\cite{WMAP} \\beq {n_b\\over n_\\gamma} = (6.5\\pm 0.4)\\times 10^{-10}\\ . \\eeq{baryond} This is a small number, but at the same time it is large enough to be a puzzle for models of particle physics.  A baryon excess this large cannot be produced in the early universe within the Standard Model of particle physics \\cite{HSather}.  Models that can explain the baryon excess typically involve exotic nonstandard physics, CP violating couplings in the Higgs or supersymmetry sectors or in the couplings of the heavy neutral leptons associated with neutrino mass \\cite{troddenreview}.  In this paper, we introduce a new mechanism for the creation of the matter-antimatter asymmetry, one associated with gravitational fluctuations created during cosmological inflation.  The conditions for generating a matter-antimatter asymmetry were stated by Sakharov almost forty years ago~\\cite{Sakharov}. First, baryon number should be violated. Second, CP should be violated.  Third, these symmetry violations should be relevant at a time when the universe is out of thermal equilibrium. Since the 1980's, it has been realized that the standard weak interactions contain processes, mediated by {\\it sphaelerons}, which interconvert baryons and leptons and are thermally activated at temperatures greater than 1~TeV.  Thus, we can also create the baryon asymmetry by creating net lepton number at high temperature through out-of-equilibrium and CP-asymmetric processes \\cite{KRS,FY}. Scenarios of this type are known as {\\it leptogenesis}.  The out-of-equilibrium conditions can be created at a phase transition or through late decay of massive particles.  The most attractive choice for a phase transition is that associated with electroweak symmetry breaking.  However, that phase transition is probably not sufficiently strongly first-order. This is known to be an obstacle to baryogenesis in the Standard Model, and most of the allowed parameter space in minimal supersymmetry is also already excluded. Particle decay asymmetries are loop-suppressed and therefore require relatively large CP-violating phases.  Such large phases are strongly constrained in supersymmetry \\cite{SUSYphases} though they still could appear in the neutrino Yukawa couplings that are used in the  Fukugita-Yanagida scenario for leptogenesis \\cite{FY}.   In any event, there is good reason to seek more effective sources of CP-violating out-of-equilibrium physics. ",
        "conclusions": ""
    },
    "0403/astro-ph0403342_arXiv.txt": {
        "abstract": "We have developed a new capacitive transimpedance amplifier (CTIA) that can be operated at 2 K, and have good performance as readout circuits of astronomical far-infrared array detectors. The circuit design of the present CTIA consists of silicon p-MOSFETs and other passive elements. The process is a standard Bi-CMOS process with 0.5 $\\mu$m design rule. The open-loop gain of the CTIA is more than 300, resulting in good integration performance. The output voltage swing of the CTIA was 270 mV. The power consumption for each CTIA is less than 10 $\\mu$W. The noise at the output showed a 1/f noise spectrum of 4 $\\mu$V/${\\sqrt {\\rm Hz}}$ at 1 Hz. The performance of this CTIA nearly fulfills the requirements for the far-infrared array detectors onboard ASTRO-F, Japanese infrared astronomical satellite to be launched in 2005. ",
        "introduction": "Establishment of cryogenic readout circuits is one of the most important technologies for astronomical application in the far-infrared wavelength region (30 - 300 $\\mu$m), because preamplifier circuits must be placed in the very vicinity of the detectors at a cryogenic temperature (below 4 K) in order to suppress electromagnetic interference (EMI) and microphonic noises. Recently, larger-formatted, two-dimensional (2D) array detectors are required for the far-infrared astronomical observation, which implies the development of larger-formatted, preamplifier array operable at the cryogenic temperature.   However, at cryogenic temperatures below 4.2 K, most of semiconductor devices such as Si JFET and bipolar transistors cannot be used due to the carrier freeze-out effect. The most probable silicon device that can be used at lower than 4.2 K is the silicon metal-oxide-semiconductor field-effect transistor (MOSFET). Of course, we can use materials other than silicon, such as GaAs, Ge, etc. However, the silicon MOSFET has a great advantage that the mass production technology established for commercial usage is available for larger-formatted array circuits. For example, in the field of near-infrared astronomy, silicon MOSFETs have already been widely used as readout circuits of near-infrared detector arrays at 15 - 80 K. Additionally, silicon MOSFETs can be operated with low power dissipation, and very suitable as parts of cryogenic readout circuit. On the other hand, silicon MOSFETs frequently show anomalous current-voltage characteristics such as \"kink\" and \"hysteresis\" below liquid helium temperature [1].   Several approaches have been dedicated to develop cryogenic readout circuits until now. The cryogenic readout circuits used for ISOPHOT [2] onboard the Infrared Space Observatory (ISO) [3] contained an amplifier with both n-MOSFETs and p-MOSFETs that was designed to eliminate anomalies mainly from n-MOSFET [4]. For MIPS (The Multiband Imaging Photometer for SIRTF) on board the Space Infrared Telescope Facility (SIRTF) [5], the cryogenic readout circuits called CRC696 are employed [6]. Special wafers that has very thin, lightly doped epitaxial layers over degenerate substrate was used for CRC696, and it solved the problem of DC instability of the readout circuit at less than 20 K [7].   On the other hand, several unique efforts have been successfully made for development of cryogenic readout circuit technology in Japan [8, 13]; p-MOSFETs of a special structure, which can be fabricated by a standard BiCMOS process, show good performances below liquid helium temperature as follows. The drain resistance and tranceconductance of the p-MOSFETs at the drain current of -1 $\\mu$A and at the drain-source voltage of -1 V are $\\sim$2 M$\\Omega$ and $\\sim$35 $\\mu$S, respectively [8]. The noise voltage at 1 Hz is 1.9 $\\mu {\\rm V_{rms}}$/${\\sqrt {\\rm Hz}}$ [8]. Moreover a cryogenic amplifier consisting of this p-MOSFET has an open-loop gain of more than 1000. On the basis of the previous work, the present work is focused in developing a capacitive transimpedance amplifier (CTIA) by using this cryogenic amplifier developed by Hirao et al. [8]. Moreover, a 60-channel CTIA array with low-pass filters and multiplexers (MPXs) are also designed by the same technology for far-infrared Ge:Ga array detectors of the Far Infrared Surveyor (FIS) [9] on board ASTRO-F [10]. Therefore, by adopting this new MOSFET structure, highly developed, commercially based silicon MOS technology enables us to make a more complicated cryogenic preamplifier array. Moreover, it is also demonstrated that we will be able to make highly integrated LSI that can be operated even at 2 K by using the same technology. In the far-infrared and submillimetric wavelength regions, we will be able to make much more integrated array detectors. In section II, we describe the design and fabrication. The results of the performance measurement and discussion are presented in section III and IV, respectively. ",
        "conclusions": "We have achieved the development of the CTIA electronics that can be operated below the liquid-helium temperature. The CTIA has shown good performances of low power consumption, low voltage noise, and sufficient open-loop gain, which are suitable for high accuracy astronomical observations in the far-infrared region.  By using this design, we fabricated 60-channel readout preamplifier arrays for the Japanese infrared astronomical satellite, ASTRO-F. This preamplifier array consists of CTIA with a reset switch and a low-pass-filter per each channel. This array was successfully bumped with a Ge:Ga monolithic photoconductor array into a complete far-infrared array detector. Details are described in [11].  The present work contributes to the establishment of electronics operated at very low temperatures, even at 2 K. We employed a standard process and commercially available, standard silicon wafers without any special requirements. Therefore, widely used LSI technologies could be directly applied to fabricate larger-scale cryogenic integrated circuits: the present 60ch-CTIA design (Fig. 4b) contains 1,800 transistors per chip. Moreover, having two differential inputs, the cryogenic amplifier used in the CTIA is an ordinary operational amplifier. The present work has advantage in these points, which are important for a broader range of application, for example, those for bolometers and super tunneling junction devices.  \\subsection{Comparison to the Previous Work} Cryogenic readout circuits had so far been developed for other infrared astronomical missions, ISO, SIRTF, and SOFIA (the Stratospheric Observatory For Infrared Astronomy) [21]. The present work has clearly improved in several points for that of ISOPHOT [18], and is comparable with that of MIPS onboard SIRTF [6] and that of AIRES onboard SOFIA [20]. Table IV summarizes the performances of these four types of readout circuits.   The first point is the number of readout circuits in one chip. We have developed two-dimensional preamplifier array that has 3 $\\times$ 20 CTIA circuits in one chip. This preamplifier array has the largest format in readout circuits for far-infrared detectors. Using the 60-channel preamplifier array, Fujiwara et al. [11] first succeeded in developing direct hybrid two-dimensional array detectors.   The second point is the amplifier gain of the CTIA. The amplifier gain of CTIA directly relates to the constancy of the responsivity of detector; low amplifier gain may introduce changes in the bias voltage. The first CTIAs for Ge:Ga detectors, used as the readout circuits of ISOPHOT, did not have an enough amplifier gain to keep the detector bias voltage constant, and introduced responsivity changes in each integration operation [18]. Young et al. [6] developed the CTIA of very large amplifier gain, CRC696, and obtained the linearity of better than 1 \\%. Erickson et al. [22] and Young et al. [23] improved the CRC696 for optimization for high background condition of SOFIA, and named SBRC190. The amplifier gain of SBRC 190 is 500. This value is lower than that of CRC696 because it employed the ac-coupling input through 10 pF. We have developed a CTIA whose amplifier gain is more than 300. This value is sufficient for compensating the change of responsivity of Ge:Ga detectors during integration, although the value of amplifier gain may be smaller than that of Young et al. [6] [23] and Erickson et al. [22].   The third point is the noise performance. As shown in Table IV, the read noise in electron number of the present work (1000 ${\\rm e^{-}}$) is larger than those of the previous work. However, this is primarily due to the larger feedback capacitance used in the present CTIA. If compared in voltage (converted from the number of electrons with the feedback capacitance), the noise performance of the present work (16 $\\mu$V) is better than those of the previous work.   The final point is the dynamic range. This parameter determines the range of brightness to be observed per unit integration time. The dynamic range is generally defined as the ratio of the well depth to the read noise (e.g., [14]). Since both the well depth and the read noise change linearly with the feedback capacitance when the noise voltage is independent of the capacitance (it is the case if current noise is negligible), the dynamic range can be directly compared although the feedback capacitances of these readout circuits are different from each other. The dynamic ranges are listed in Table 4. The value of the present work is the second largest. Consequently, the readout circuit will have larger dynamic range as well as better read noise performance than those of MIPS and ISOPHOT.  \\subsection{Improvement of the Noise Performance} In the present work, the noise performance does not fully satisfy the requirement of FIS. In order to improve the signal-to-noise (S/N) ratio, we changed the feedback capacitance into smaller values (3 pF). By changing the feedback capacitance from 10 pF to 3 pF, kTC noise, which is inversely proportional to the square root of feedback capacitance, increases from 1.7 $\\mu$V to 3.0 $\\mu$V in this case. However, as suggested in section III-B, kTC noise does not contribute to the S/N ratio. A preliminary result shows that the noise spectrum is comparable with that in Fig.7, and thus the current noise is negligible. As a result of this improvement, three times higher S/N ratio is expected, and then, the detection limit will be acceptable for the FIS onboard ASTRO-F."
    },
    "0403/astro-ph0403174_arXiv.txt": {
        "abstract": "We present results from a study of high- and low-ionization emission line ratios as a function of projected velocity for a sample of eight active galactic nuclei (AGNs). Our results are based on analysis of high signal-to-noise optical and \\emph{Hubble Space Telescope (HST)} UV spectra. Comparing the emission line ratios to those predicted by photoionization models indicates that the physical conditions responsible for the high-ionization emission lines are consistent with a wind, whereas those of the low-ionization lines are consistent with a virialized disk. ",
        "introduction": "It is evident from reverberation mapping studies that high- and low-ionization emission lines in AGNs show different time-lag responses to continuum variability. Presumably, they are generated at different distances from the central engine (Korista et al; 1995, Santos-Lle\\'{o} et al.; 1997, among others). However, it is still a matter of debate whether the kinematics of the high- and low-ionization regions are fundamentally different; e.g., is the high-ionization gas associated with a wind and the low-ionization gas with a disk? In this study we look for evidence of virialized motion from the emission lines of both high- and low-ionization gas of AGNs. We do this by looking at line ratios as a function of projected velocity across the emission line profiles. ",
        "conclusions": "The results shown in Figure 1 are consistent with the Balmer line emitting gas being virialized. The lower velocity low-ionization gas sees a smaller $\\it{\\Phi_H}$ in accord with the inverse-square law. Figure 2, however, is evidence of fundamentally different kinematics for the high-ionization gas that produces C\\,IV. In this case, the independence of physical conditions with respect to projected velocity, averaged over the sample, is consistent with gas motion that is not virialized.  These results are consistent with the high-ionization lines arising in a wind and the low-ionization lines arising in a disk."
    },
    "0403/astro-ph0403497_arXiv.txt": {
        "abstract": "{ We present the first near-infrared $K$-band long-baseline interferometric measurement of the nucleus of the prototype Seyfert\\,2 Galaxy NGC~1068 with resolution $\\lambda/B \\sim$\\,10\\,mas obtained with the Very Large Telescope Interferometer (VLTI) and the two 8.2\\,m diameter Unit Telescopes UT\\,2 and UT\\,3. The adaptive optics system MACAO (Multi Application Curvature Adaptive Optics) was employed to deliver wavefront-corrected beams to the $K$-band commissioning instrument VINCI. A squared visibility amplitude of 16.3\\,$\\pm$\\,4.3\\,\\% was measured for NGC~1068 at a sky-projected baseline length of 45.8\\,m and azimuth angle 44.9\\,deg. This value corresponds to a FWHM of the $K$-band intensity distribution of 5.0\\,$\\pm$\\,0.5\\,mas (0.4\\,$\\pm$\\,0.04\\,pc at the distance of NGC~1068) if it consists of a single Gaussian component. Taking into account $K$-band speckle interferometry observations (Wittkowski et al. \\cite{wittkowski}; Weinberger et al. \\cite{weinberger}; Weigelt et al. \\cite{weigelt}), we favor a multi-component model for the intensity distribution where a part of the flux originates from scales clearly smaller than $\\sim$\\,5\\,mas ($\\lesssim$\\,0.4\\,pc), and another part of the flux from larger scales. The $K$-band emission from the small ($\\lesssim$\\,5\\,mas) scales might arise from substructure of the dusty nuclear torus, or directly from the central accretion flow viewed through only moderate extinction. ",
        "introduction": "\\begin{table*} \\caption{Calibration sequence of NGC~1068. JD is the Julian date of the observation, $N$ the number of processed interferograms, $B$ and $Az$ the projected baseline length and azimuth angle (east of north), and $\\mu^2$ the obtained coherence factors with their statistical errors. The interferometric efficiency (IE) given in italic characters is the value adopted for the calibration of the NGC~1068 data.} \\label{Cal_ngc1068} \\begin{tabular}{ccccclll} \\hline JD & $N$ & B\\,(m) & $Az$ & $\\mu^{2} \\pm$ stat. & IE $\\pm$ stat. $\\pm$ syst. & $V^2 \\pm$ stat. $\\pm$ syst. & Target\\\\ \\hline 2452948.717 & 250 & 45.79 & 44.87 & $0.054 \\pm 0.017$ & $\\it 0.340 \\pm 0.009 \\pm 0.001$ & $0.158 \\pm 0.049 \\pm 0.001$ & \\object{NGC~1068} \\\\ 2452948.722 & 214 & 46.00 & 45.13 & $0.062 \\pm 0.031$ & $\\it 0.340 \\pm 0.009 \\pm 0.001$ & $0.182 \\pm 0.091 \\pm 0.001$ & \\object{NGC~1068} \\\\\\hline 2452948.761 & 407 & 46.62 & 45.40 & $0.319 \\pm 0.005$ & $0.348 \\pm 0.006 \\pm 0.001$ & & \\object{HD\\,20356}\\\\ 2452948.767 & 438 & 46.64 & 45.62 & $0.304 \\pm 0.005$ & $0.331 \\pm 0.005 \\pm 0.001$ & & \\object{HD\\,20356}\\\\ 2452948.772 & 416 & 46.63 & 45.80 & $0.315 \\pm 0.006$ & $0.344 \\pm 0.007 \\pm 0.001$ & & \\object{HD\\,20356}\\\\ \\hline \\end{tabular} \\end{table*} The Seyfert galaxy NGC~1068 harbors one of the brightest and closest active galactic nuclei (AGN). Active galaxies appear as types 1 and 2, where the spectra of the former exhibit broad and narrow emission lines, and those of the latter show only narrow lines. Antonucci \\& Miller (\\cite{antonuccimiller}) suggested that the broad-line emission regions are located inside an optically and geometrically thick disk and that central continuum and broad-line photons are scattered into the line-of-sight by free electrons above and below the disk. Depending on the observer's viewing angle, the broad-line emission region is either obscured or not. This suggestion is now widely accepted and has evolved to the so-called ``unified scheme of AGN'' (e.g., Antonucci \\cite{antonucci}).  Various theoretical models of the postulated dusty tori have been presented by, for instance, Krolik \\& Begelmann (\\cite{krolikbegelmann}), Pier \\& Krolik (\\cite{pierkrolik1}), Granato \\& Danese (\\cite{granato}), Efstathiou \\& Rowan Robinson (\\cite{efstathiou1}), Manske et al. (\\cite{manske}), Nenkova et al. (\\cite{nenkova}), and Vollmer et al. (\\cite{vollmer}). Owing to the lack of spatially resolved observations, the models were usually compared to the integrated spectrum of the core of NGC~1068. It turned out that torus models with a wide variety of geometries, spatial extensions, and optical depths are consistent with this spectrum. In addition, the nature of the central emission source itself could not be well constrained because the flux spectrum of the very inner nuclear engine could not be separated from the emission of surrounding material.  Several high-resolution infrared observations of NGC~1068 showing a compact central IR core and surrounding structure were carried out by, for instance, Thatte et al. (\\cite{thatte}), Rouan et al. (\\cite{rouan}), Wittkowski et al. (\\cite{wittkowski}), Weinberger et al. (\\cite{weinberger}), and Bock et al. (\\cite{bock}). Wittkowski et al. (\\cite{wittkowski}) presented both a $K$-band visibility function up to spatial frequencies corresponding to a baseline of 6\\,m, as well as the first 76\\,mas resolution $K$-band image of the nucleus of NGC~1068 obtained by bispectrum speckle interferometry. The compact central IR core was resolved with a Gaussian FWHM of $\\sim$\\,30\\,mas (2\\,pc). New $K$-band and $H$-band bispectrum speckle interferometry observations were performed by Weigelt et al. (\\cite{weigelt}).  Very recently, the first interferometers consisting of 8--10\\,m class telescopes started operations, and they have already succeeded in observing AGN with much higher spatial resolutions. The potential of optical/infrared interferometry to investigate AGN was recently discussed by Wittkowski et al. (\\cite{wittkowski2}). Swain et al. (\\cite{swain}) reported the first $K$-band interferometric observations of the Seyfert\\,1 galaxy NGC~4151 obtained with the 85\\,m baseline of the Keck Interferometer. Jaffe et al. (\\cite{jaffe}) reported on the first mid-infrared interferometric observation of the dusty torus of NGC~1068 using VLTI/MIDI.  In the present letter, we report on the first $K$-band long-baseline interferometric observation of NGC~1068. ",
        "conclusions": "We have obtained $K$-band interferometric observations of the nucleus of the Seyfert\\,2 galaxy NGC~1068, and hereby show that the correlated magnitude of this AGN up to a baseline of at least $\\sim$\\,50\\,m is such that it can be studied with the VLTI at near-infrared wavelengths.  Together with other observations, we conclude that a $K$-band flux of $\\gtrsim$\\,50\\,mJy originates from scales clearly smaller than about 5\\,mas or 0.4\\,pc and another part of the flux from larger scales. Our VLTI/VINCI measurement alone sets an upper limit of $\\le$\\,6.7\\,mas at the 3\\,$\\sigma$ level to the Gaussian FWHM of the very compact component. The origin of this newly constrained small-scale emission can be interpreted as substructure of the dusty torus, as for instance part of a clumpy inner cavity or distinct clumps forming the torus, or as direct emission from the central accretion flow viewed through only moderate extinction in the $K$-band."
    },
    "0403/astro-ph0403204_arXiv.txt": {
        "abstract": "It has recently been stated that the warp structure observed around the star $\\beta$ Pictoris may be due to four planets embedded in its debris disk \\cite{wj03}. It, therefore, becomes important to investigate for what range of parameters, and for how long such a multibody system will be dynamically stable. We present the results of the numerical integration of the suggested planetary system for different values of the mass and radii of the planets, and their relative positions and velocities. We also present a statistical analysis of the results in an effort to understand the relation between the different regions of the parameter-space of the system and the duration of the orbital stability of the embedded planets. ",
        "introduction": "Beta Pictoris, a type A5 IV star at a distance of $\\sim$19 pc from the Earth, is one of the youngest close stars with an approximate age of  $12^{+8}_{-4}$ Myr \\cite{zuckerman01}. Observational evidence indicates that this star is surrounded by a planetary debris disk \\cite{st84,ho85, weinberger03}. The close proximity of such a young circumstellar disk has made $\\beta$ Pictoris an ideal candidate for the study of the evolution of protoplanetary disks and planetary system formation.  In the past several years, there have been a number of reports of the detection of symmetric warps in the $\\beta$ Pictoris debris disk \\cite{wj03, weinberger03,burrows95, mouillet97, heap00}. Models have best explained these warps as gravitational perturbations caused by planetary companions \\cite{burrows95, mouillet97, heap00}. Among these reports, the disk warps discovered by \\citet{wj03} were explained as the edge-on projection of debris rings orbiting the central star.  Models of the flux density allow some parameters of these rings to be determined via $\\chi^2$ fitting. \\citet{wj03} proposed a multiple planet system to account for ring formation, noting that all adjacent rings are in mean-motion resonances.  In consideration of these findings, we undertake here a study of the dynamical evolution of a multibody system similar to that proposed by \\citet{wj03}.  We perform a statistical analysis of the stability of randomly generated systems within a portion of the total available parameter-space and briefly analyze the relationship between the parameters of the system and its orbital stability. ",
        "conclusions": "\\label{conclusions}  We have analyzed the stability of the proposed five-body planetary system embedded in the $\\beta$ Pictoris debris disk \\citep{wj03} for over 14000 initial conditions. Our results indicate that the majority of systems became unstable between a time of 1 to 10 million years. There were only 8 unique simulations that remained stable for the entire 50\\,Myr integration time. These systems contained planets with masses less than 2.4 times the mass of Jupiter.   \\begin{theacknowledgments} This work is partially supported by the Carnegie Institution of Washington Internship Program, and also an REU Site for Undergraduate Research Training in Geoscience, NSFEAR-0097569 for JHC, and NASA Origins of the Solar System Program under grant NAG5-11569, and also the NASA Astrobiology Institute under Cooperative Agreement NCC2-1056 for NH. \\end{theacknowledgments}"
    },
    "0403/astro-ph0403032_arXiv.txt": {
        "abstract": "We have performed a fully three-dimensional general relativistic magnetohydrodynamic (GRMHD) simulation of jet formation from a thin accretion disk around a Schwarzschild black hole with a free-falling corona. The initial simulation results show that a bipolar jet (velocity $\\sim 0.3c$) is created as shown by previous two-dimensional axisymmetric simulations with mirror symmetry at the equator.  The 3-D simulation ran over one hundred light-crossing time units ($\\tau_{\\rm S} = r_{\\rm S}/c$ where $r_{\\rm S} \\equiv 2GM/c^2$) which is considerably longer than the previous simulations. We show that the jet is initially formed as predicted due in part to magnetic pressure from the twisting the initially uniform magnetic field and from gas pressure associated with shock formation in the region around $r = 3 r_{\\rm S}$.  At later times, the accretion disk becomes thick and the jet fades resulting in a wind that is ejected from the surface of the thickened (torus-like) disk. It should be noted that no streaming matter from a donor is included at the outer boundary in the simulation (an isolated black hole not binary black hole). The wind flows outwards with a wider angle than the initial jet. The widening of the jet is consistent with the outward moving torsional Alfv\\'{e}n waves (TAWs). This evolution of disk-jet coupling suggests that  the jet fades with a thickened accretion disk due to the lack of streaming material from an accompanying star. ",
        "introduction": "Relativistic jets have been observed in active galactic nuclei (AGNs) (e.g. Urry \\& Pavovani 1995; Ferrari 1998; Blandford 2002), in microquasars in our Galaxy (e.g. Mirabel \\& Rodriguez 1999), and it is believed that they originate in the regions near accreting (stellar) black holes and neutron stars \\citep{mku01}.  To investigate the dynamics of accretion disks and the associated jet formation, we have performed jet formation simulations using a full 3-D GRMHD code. This magnetic-acceleration mechanism has been proposed not only for AGN jets, but also for (nonrelativistic) protostellar jets (see \\citet{mku01}). Kudoh, Matsumoto, \\& Shibata (1998) found that the terminal velocity of a jet is comparable to the rotational velocity of the disk at the foot of the jet in nonrelativistic MHD simulations, and that a relativistic jet should be formed near the event horizon. It should be noted that the mechanism of jet generation in neutron star binaries may be different from black hole binaries.  Recently, an ultra-relativistic outflow was observed from an accreting stellar-mass neutron star (Fender et al. 2004). Fender et al. (2004) have concluded that the generation of highly relativistic outflows does not require properties unique to black holes, such as an event horizon. It seems plausible that production of ultrarelativistic jets may be a generic feature of accretion onto neutron stars such as Cir-X and Sco X-1 at or near to the Eddington limit. This result differs from the suggestions that the velocities of jets would be limited to the surface escape speed (Livio 1999).  The fluorescent iron line (K$_{\\alpha}$) seen in the X-ray spectrum of many active galactic nuclei (AGNs), and particularly in Seyfert galaxies, has provided us with an extremely useful probe of the region close to an accreting black hole. This is an important diagnostic with which to study Doppler and gravitational redshifts, thereby providing key information on the local kinematics of the material (see the reviews by Fabian et al. 2000 and Reynolds \\& Nowak 2003). The simultaneous X-ray/IR flares from a black hole/relativistic jet system are the observational evidence revealing a time-dependent interaction between the jet and inner disk (e.g. Ueda et al. 2002). In regions where disk-jet interactions take place near a black hole, the plasma and magnetic fields interact in a complicated manner; therefore a 3-dimensional GRMHD simulation such as we perform here is required to investigate dynamics near the horizon and ultimately to couple predicted dynamics to the iron line.  \\citet{ksk99} have investigated in 2-D, the dynamics of an accretion disk initially threaded by a uniform magnetic field in a non-rotating corona (either in a state of steady fall or in hydrostatic equilibrium) around a non-rotating (Schwarzschild) black hole. The numerical results show that matter in the disk loses angular momentum by magnetic braking, then falls into the black hole.  The disk falls faster in this simulation than in the non-relativistic case because of general-relativistic effects that are important below $3 r_{\\rm S}$, where $r_{\\rm S} \\equiv 2GM/c^2$ is the Schwarzschild radius. A centrifugal barrier at $r = 2 r_{\\rm S}$ strongly decelerates the infalling material. Plasma near the shock at the centrifugal barrier is accelerated by the ${\\bf J} \\times {\\bf B}$ force and forms bipolar jets. Inside this {\\it magnetically driven jet}, the gradient of gas pressure also generates a jet above the shock region ({\\it gas-pressure driven jet}). This {\\it two-layered jet structure} is formed both in a hydrostatic corona and in a steady-state falling corona. Koide et al. (2000) have also developed a new GRMHD code using a Kerr geometry and have found that, with a rapidly rotating ($a \\equiv J/J{\\max} = 0.95$, $J_{\\max} = GM^{2}/c$, $J$ and $M$ are angular momentum and mass of black hole) black-hole magnetosphere. the maximum velocity of the jet is (0.3 - 0.4) $c$. Recently the extraction of rotational energy from plasma in the ergosphere has been investigated (Koide et al. 2002; Koide 2003). All of these GRMHD simulations were made assuming axisymmetry with respect to the $z$-axis and mirror symmetry with respect to the plane $z = 0$.  It is known that the axisymmetric assumption suppresses azimuthal instabilities.  Magnetic fields are expected to play an essential role in the dynamics of most relativistic objects such as active galactic nuclei, X-ray binaries, gamma-ray bursts, and quasars. In accretion disks, the most important role of the magnetic field is angular momentum transport, which twists the field, and results in torsional Alfv\\'{e}n waves (TAWs) and associated jet formation (Koide et al. 1999). Based on the importance of magnetorotational instability (MRI) in accretion disks \\citep{bah98}, the nonlinear growth of these instabilities has been investigated by three-dimensional global MHD simulations of black hole accretion disks \\citep{haw00,igu00,haw01,hak01,mhm00,hbs01, arc01,raa01,hak02,hab02,kah02,mam03,igu03} using a pseudo-Newtonian potential \\citep{pacy80}. \\citet{mam03} investigated black hole accretion flow using three-dimensional global resistive MHD simulations. They conclude that the reconnection in the innermost region releases enough energy to explain the origin of X-ray bursts from black hole candidates.  Recently, more GRMHD simulations have been developed by \\citet{dev02,dev03a,gam03}. In particular, \\citet{gam03} have employed the numerical scheme HARM (High Accuracy Relativistic Magnetohydrodynamics) in which the conservation schemes for GRMHD are used \\citep{ksk98,ksk99,kmsk00}.  A nonconservative scheme for GRMHD following a ZEUS-like approach has been developed also by \\citet{dev02}.  They also compare the code performance of 1-D and 2-D test cases.  \\citet{dev03b} have investigated the properties of accretion flows with tori in the Kerr metric using 3-D GRMHD simulations. The MRI in torus has also been investigated using Schwarzschild and Kerr metrics (De Villiers, Hawley, \\& Krolik 2003; Hirose, et al. 2004, De Villiers, et al. 2004; Krolik, Hawley, \\& Hirose 2004).  Different studies have chosen various initial magnetic fields. We have chosen an initial magnetic field from the Wald's exact analytic solution for a Kerr black hole immersed in a uniform magnetic field (Wald 1974; Punsley 2001; Komissarov 2004). De Villiers, Hawley \\& Krolik (2003) have performed simulations in which the initial conditions for the accretion disks and magnetic fields are appropriate for MRI research. They have investigated the properties of accretion flows in the Kerr metric including bi-conical outflows in the quarter sphere with periodic boundary conditions in the azimuthal direction. These simulations use the torus accretion disks as an initial condition. The inner edge of the accretion disk is located at $r = 15 r_{\\rm S}$, and the inner edge moves into the innermost stable radius within one orbit (De Villiers, Hawley, \\& Krolik 2003). In these simulations, the initial field consists of axisymmetric (azimuthal) field loops, laid down along an isodensity surface within the torus, which is different from the initial uniform magnetic field used in simulations by Koide, Shibata, \\& Kudoh (1999) and Nishikawa et al. (2002, 2003). The initial weak loop magnetic fields contribute to the generation of the MRI and the associated jet. However, the structure of magnetic fields piled up near the black hole in their simulations is different from the magnetic fields which are twisted and pinched (not wound) by the free-falling corona and accretion disk (Koide, Shibata, \\& Kudoh 1998, 1999; Koide et al. 2000; Nishikawa et al. 2002, 2003). The absence of a free-falling corona (Bondi \\& Hoyle 1944; Bondi 1952; Michel 1972) as well as the initial loop magnetic field in (De Villiers, Hawley \\& Krolik 2003) results in different dynamics of the accretion disk and associated jet formation, compared to the case with the free-falling corona and initial uniform magnetic field in Koide, Shibata, \\& Kudoh (1999) and Nishikawa et al. (2002, 2003).  In this paper we present a fully three-dimensional, GRMHD simulation of jet formation with a thin accretion disk.  Our motivation for this type of simulation is to determine the parameters necessary for relativistic jet formation and the resulting interaction and instabilities found between the accretion disk and the black hole.  The simulation we present here was performed using the same parameters as in \\citet{ksk99} in order to determine the physical differences resulting from a full 3-D versus a 2-D simulation with axisymmetry and  mirror symmetry at the equator. The three-dimensional simulation allows us to study the evolution of jet formation because it is run for longer light-crossing times than the previous two-dimensional simulations.  We find that at the later stages of the simulation the accretion disk becomes thick and a wind is formed with a much wider angle than the collimated jet formed at the earlier stage.  In Section 2, we describe the governing GRMHD equations.  The simulation model and numerical method is described in Section 3. Our new simulation results are presented in Section 4 followed by a summary and discussion in Section 5. ",
        "conclusions": "We have presented a fully 3-D GRMHD simulation of a Schwarzschild black hole accretion disk system with moderate resolution.  Because no deliberate asymmetric perturbation was introduced, the present 3D simulation remains largely axisymmetric. While we found some numerical differences between this simulation and previous two-dimensional simulations, our results reaffirm the previous result of a two-layered jet. Our results suggest that the jet develops as in the 2-D case. Our longer simulation has allowed us to study jet evolution.  At the end of this simulation the jet fades and a weak wind is generated by a thickened accretion disk. This phenomena was not observed in the two-dimensional simulations because they did not run for this duration.  In this simulation at the outer boundary of the accretion disk no matter is injected, therefore after accreted matter is ejected from the accretion as a jet, the power to generate a jet is dissipated and the jet is switched to a wind. However, if streaming matter is injected at the outer boundary of accretion disk, transient changes may be controlled by accretion rates with instabilities and be related to the state transitions. Such disk-jet coupling in black hole binaries is reviewed by Fender (2002, 2003) and Fender, Belloni, \\& Gallo (2004). Black hole binaries exhibit several different kinds of X-ray `state'. The two most diametrically opposed, which illustrate the relation of jet formation to accretion, can be classified as low/hard/off and high/soft states (Fender 2003). These two states provide different luminosity. Pellegrini et al.\\ (2003) have discussed a nuclear bolometric luminosity and an accretion luminosity $L_{\\rm acc}$ in terms of the accretion rate $\\dot{M}$ and jet power. Clearly these issues can be investigated by further 3-D GRMHD simulations, and future simulations will investigate jet formation with different states of the black hole including streaming material from an accompanying star (e.g., Miller et al.\\ 2003).  Additionally, we plan to study the development of instabilities along the azimuthal direction such as the proposed accretion-ejection instability (AEI) (e.g. Tagger \\& Pellat 1999; Caunt \\& Tagger 2001; Rodriguez et al. 2002; Vani\\'{e}re, Rodriguez \\& Tagger 2002),  the magnetorotational instability (MRI) (e.g. Balbus \\& Hawley 1991, 1998), and the ``screw'' (current-driven helical kink) instability (e.g. Li 2000).  Such simulation studies may indicate how the observed variabilities in AGN jets (Blandford 2003; Mirabel 2003) and microquasars (e.g. Eikenberry et al. 1998; Ueda et al. 2002; Nobil 2003; Rau, Greiner, \\& McCollough 2003) can be produced. Initial magnetic field configurations affect the dynamics of accretion disks and the associated phenomena such as jet formation and its variabilities, and thus we plan to study different initial field configurations.  The Fe K$\\alpha$ fluorescent emission line in active galactic nuclei (AGN) is believed to originate in a relativistic accretion disk around a black hole (e.g. see reviews Fabian et al. 2000; Reynolds \\& Nowak 2003). The X-ray continuum emission of MCG-6-30-15 is highly variable (e.g. Vaughan \\& Fabian 2003). The source spectrum and variablity can be explained by two-components consisting of a variable power law component (PLC) and an almost constant reflection-dominated component (RDC) containing  the iron line (Fabian \\& Vaughan 2003). The region where the flare (or multiple flares) which created the PLC must be located close to the rotation axis and within 3-4 gravitational radii to reproduce the highly variable behavior, (Miniutti et al. 2003). Since this behavior is most likely powered by magnetic field from the disk, possibly linking to the hole (e.g. Blandford \\& Znajek 1977; Wilms et al. 2001), and perhaps forming the base of jet, it is essential to perform GRMHD simulations. These issues are being investigated with 3-D GRMHD simulations with a Kerr black hole with streaming matter from an accompanying star."
    },
    "0403/astro-ph0403518_arXiv.txt": {
        "abstract": "We propose a new imaging gamma-ray detector in the  MeV region. By measuring the directions and energies of not only a scattered gamma ray but also a recoil electron, the direction of an incident gamma ray would be essentially reconstructed event by event. Furthermore, one of two measured (zenith and azimuth) angles of a recoil electron gives us an additional redundancy which enables us to reject the  background events by kinematic constraints. In order to measure the track of a recoil electron, the micro Time Projection Chamber($\\mu$-TPC) has been developed, which can  measure the successive positions of the track of charged particles in a few hundred micron meter pitch. The $\\mu$-TPC consists of the new type of a gas proportional chamber: micro PIxel gas Chamber ($\\mu$-PIC) which is one of wireless gas chambers and expected to be robust and stable. Using this $\\mu$-TPC and the Anger camera for the detection of a scattered gamma ray, we have obtained the first gamma-ray image by the full reconstruction of the direction of gamma rays event by event. ",
        "introduction": "Nowadays gamma-ray astronomy has become a very promising field of astronomy. However the MeV region of gamma rays is still uncultivated, although this region would surely show us new aspects of high energy phenomena in the universe. In fact, COMPTEL in Compton Gamma-Ray Observatory had observed for 9 years, and found about 30 celestial objects emitting MeV gamma rays\\citep{schon2002}. However, taking into account of more than two hundreds discoveries of GeV gamma-ray emitters by EGRET, it is expected that more than a hundred celestial objects would be detected in the MeV region. This is almost due to the nonexistence of a reliable imaging method of MeV gamma rays scattered via Compton process in the matter.  Recently multi-Compton telescopes have eagerly been studied as a next generation of the detector launched in satellites\\citep{Kurfess2003}. Here we propose a new imaging detector of MeV gamma rays measuring both the directions and energies of not only a scattered gamma ray but also a recoil electron, which enables us to reconstruct the direction of an incident gamma ray event by event. Figure \\ref{comp} shows the schematic view of the planned detector based on this concept for a balloon-borne experiment, which consists of a 30cm cubic $\\mu$-TPC surrounded by a scintillation pixel detector. Also it is pointed out that  one of two measured angles of a recoil electron gives us an  additional redundancy which enables us to reject almost all the  background events by the kinematic constraint.  Classical Compton telescopes measure the energies of both the scattered gamma ray and recoil electron, and  one direction of the scattered gamma ray. Since this method losses the direction of the recoil electron, only one angle ($\\phi$) of the incident gamma ray can be determined. Furthermore there remains no redundancy to reject the background. In the MeV region, not only strong celestial diffuse gamma rays but also the huge background gamma rays from the albedo and the satellite itself by irradiation of cosmic rays come to the detector in space. In addition, internal radioactivities in the detector emit both gamma rays and beta decay electrons faking a real gamma ray in Compton detectors. Good reviews of the analysis for the background in COMPTEL were presented in this workshop\\citep{Sch2003, Ryan2003}. In order to achieve a sensitivity ten times better than COMPTEL, another powerful method for the background rejection seems to be inevitable for the next  MeV gamma-ray imaging detector.  Here we introduce the new type of a gamma-ray imaging detector based on a wireless gas detector, and the results of its performance study are presented.     \\begin{figure} \\includegraphics*[width=0.5\\linewidth]{MeV_gamma_scale.eps}\\hspace{6mm} \\includegraphics*[width=0.35\\linewidth]{MeV_axes2_bw.eps} \\caption{ Schematic view of the planned detector for a balloon-borne experiment, which consists of a 30cm cubic $\\mu$-TPC surrounded by a scintillation pixel detector(left figure). In the right figure, definitions of the angles are shown.} \\label{comp} \\end{figure} ",
        "conclusions": ""
    },
    "0403/astro-ph0403296_arXiv.txt": {
        "abstract": "We present a method designed to identify the spectral type of young (less than $\\sim30$ days after explosion) and nearby ($z \\simlt 0.05$) supernovae (SNe) using their broad-band colors. In particular, we show that stripped-core SNe (i.e., hydrogen deficient core-collapse events, spectroscopically defined as SNe~Ib and SNe~Ic, including broad-lined SN 1998bw-like events) can be clearly distinguished from other types of SNe. Using the full census of nearby SNe discovered during the year 2002, we estimate the impact that prompt multi-band photometry, obtained by $1$~m class telescopes, would have on the early identification of stripped-core events. Combining this new approach with ongoing spectroscopic follow-up programs, one can expect $\\sim 20$ nearby, stripped-core events to be identified, each year, around, or before, maximum light. Follow-up studies, including prompt, multi-epoch optical spectroscopy and spectropolarimetry, as well as radio and X-ray observations, could greatly increase our understanding of these events, and shed new light on their association with cosmological gamma-ray bursts. ",
        "introduction": "Several types of supernovae (SNe) are recognized, and are most commonly determined by the attributes of the SN optical spectra. The most basic division is between hydrogen-rich (type II) SNe, and hydrogen deficient (type I) SNe (Minkowski 1941). Type I events are further divided into SNe Ia, which usually show the prominent Si II absorption trough near $6150$~\\AA, SNe Ib, which lack this feature but show prominent He lines, and SNe Ic, whose spectra show neither (see Filippenko 1997 for a review).  Direct detection of the progenitor stars on pre-explosion images has established that type II SNe result from the explosion of massive stars, presumably following the gravitational collapse of the stellar core (SN 1987A, White \\& Malin 1987; SN 1993J, Aldering et al. 1994; SN 2003gd, Van Dyk, Li, \\& Filippenko 2003, Smartt et al. 2004). The existence of transition objects, which appear to be SNe II in spectra obtained shortly after explosion, but later evolve and become very similar to SNe Ib and SNe Ic (SNe IIb; e.g., Filippenko 1997), as well as the apparent association of SNe Ib and Ic with young stellar populations (e.g., Van Dyk, Hamuy, \\& Filippenko 1996), suggest that these events also result from the core collapse of massive stars. These stars are assumed to have lost most (in the case of SNe IIb) or all (in SNe Ib and Ic) of their hydrogen envelope, as well as, in the case of SN Ic progenitors, most of their helium, either through stellar winds or due to binary interaction. In contrast, SNe Ia are frequently found among old stellar populations, e.g., in elliptical galaxies, and are unlikely to be associated with very young progenitor stars. These events are commonly assumed to be the result of a thermonuclear runaway explosion of a white-dwarf star, which reaches the Chandrasekhar mass either through accretion from, or a merger with, a binary companion.  Of all the various types of SNe, the ones that are probably the least studied are SNe Ib and Ic, which will be referred to henceforth as stripped-core events. Contributing factors probably include the intrinsic rareness, relatively low luminosity, and heterogeneity of these events, along with the absence of a high-profile science driver for such work, as provided, for example, by cosmological distance estimation for studies of SNe Ia. Most of the progress in this field was due to the (rare) occurrence of bright stripped-core events in nearby galaxies, e.g., SN 1994I in M51. The only recent comprehensive study of stripped-core events was conducted by Matheson et al. (2001).  This situation is perhaps about to change, following the emerging connection between stripped-core SNe and long-duration gamma-ray bursts, and in particular, the associations between GRB 980425 and SN 1998bw (Galama et al. 1998) and GRB 030329 and SN 2003dh (Stanek et al. 2003; Hjorth et al. 2003; Matheson et al. 2003) -- see Lipkin et al. (2003) for a recent review. The amount of overlap between GRBs and stripped-core SNe is intriguing. Are all GRBs associated with a SN? Are all stripped-core SN explosions triggered by, or associated with, a GRB or GRB-like mechanism (i.e., a non-isotropic, jetted explosion; Khokhlov et al. 1999) ? Early-time studies of stripped-core SNe may provide the answer to this last quesion. ",
        "conclusions": "The emerging connection between stripped-core SNe and GRBs is likely to revitalize the study of this class of SNe. The rareness and relative faintness of these events has so far held back advancements in systematic studies of their properties, which would probably help understanding their relation to GRBs and the underlying physics. In particular, studies possible only while these SNe are young (i.e., at, or preferably before, maximum light) have been hindered by frequent delays between the discovery, spectroscopic type determination, and the announcement of the type in IAU circulars.  To facilitate prompt classification of young stripped-core SN candidates using broad-band colors, we have added the spectra of SN 1998bw, associated with GRB 980425, and the similar SN 1998bw-like SN 2002ap, to the spectroscopic data base of Poznanski et al. (2002). We have applied the methods of these authors to this expanded data base, and have studied in detail methods to classify young (around peak, or before) and nearby ($z\\le0.05$) SNe. We have shown that stripped-core events can be efficiently selected using four-band photometry focused on the bluer bands. We have also demonstrated that the required photometric accuracy can be obtained with a $1$~m class telescope, and that reddening effects do not significantly compromise our analysis. Application of this method to the SN population discovered by current SN surveys will probably double the number of young stripped-core SNe identified at or before peak magnitude, and, combined with current SN spectroscopy programs, will supply about $20$ new targets each year for detailed follow-up studies in radio, optical, and X-ray bands."
    },
    "0403/astro-ph0403405_arXiv.txt": {
        "abstract": "Braneworld black holes existing today may be of primordial origin, or may even be produced in high energy particle collisions in the laboratory and in  cosmic ray showers as well. These black holes obey a modified mass-radius relationship compared to standard Schwarzschild black holes.  Using the variational principle we calculate the bending angle of a light ray near the horizon of a braneworld black hole in the weak field limit. We next derive the expressions of several lensing quantities like the Einstein radius and the magnification for a point light source. These expressions are modified compared to  the lensing quantities for standard Schwarzschild black holes and contain the scale of the extra dimensions. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403633_arXiv.txt": {
        "abstract": "We study the statistics of large-separation gravitational lens systems produced by non-spherical halos in the Cold Dark Matter (CDM) model.  Specifically, we examine how the triaxiality of CDM halos affects the overall lensing probabilities and the relative numbers of different image configurations (double, quadruple, and naked cusp lenses).  We find that triaxiality significantly enhances lensing probabilities by a factor of $\\sim$2--4, so it cannot be ignored.  If CDM halos have central density slopes $\\alpha \\lesssim 1.5$, we predict that a significant fraction ($\\gtrsim$20\\%) of large-separation lenses should have naked cusp image configurations; this contrasts with lensing by isothermal ($\\alpha \\approx 2$) galaxies where naked cusp configurations are rare.  The image multiplicities depend strongly on the inner density slope $\\alpha$: for $\\alpha=1$, the naked cusp fraction is $\\gtrsim$60\\%; while for $\\alpha=1.5$, quadruple lenses are actually the most probable.  Thus, the image multiplicities in large-separation lenses offer a simple new probe of the inner density profiles of dark matter halos.  We also compute the expected probabilities and image multiplicities for lensed quasars in the Sloan Digital Sky Survey, and argue that the recent discovery of the large-separation quadruple lens SDSS J1004+4112 is consistent with expectations for CDM. ",
        "introduction": "The Cold Dark Matter (CDM) model of structure formation naturally predicts the existence of strong gravitational lens systems with image separations of $\\sim\\!10''$ or even larger.  Observations of massive clusters of galaxies have revealed many systems of ``giant arcs'' representing lensed images of background galaxies \\citep{lynds86,soucail87,luppino99,gladders03,zaritsky03}. However, until recently all lensed quasars and radio sources had image separations $<\\!7''$ corresponding to lensing by galaxies, despite some explicit searches for lenses with larger separations \\citep{phillips01,ofek01}.\\footnote{\\citet{miller04} recently reported six pairs of quasars in the Two-Degree Field (2dF) Quasar Redshift Survey that are candidate lenses with image separations on the scale of an arcminute; but none of the candidates has been confirmed, and theoretical arguments by \\citet{oguri03a} indicate that it would be quite surprising if any of the systems are lenses.}  Lensing of quasars by clusters was finally observed with the recent discovery and confirmation of SDSS J1004+4112, a quadruple lens with an image separation of $14\\farcs62$ found in the Sloan Digital Sky Survey \\citep{inada03,oguri04}.  This lens confirms an important prediction of the CDM model; indeed, the lensing probability inferred from the discovery is in agreement with reasonable values of the cosmological parameters \\citep{oguri04}.  The statistics of large-separation lenses can be used to place constraints on the density profile of dark halos \\citep*{maoz97, keeton01c,keeton01a,wyithe01,takahashi01,sarbu01,li02,oguri02a, oguri02b,huterer04,kuhlen04}, or determine the abundance of massive dark halos \\citep{narayan88,wambsganss95,kochanek95,nakamura97, mortlock00,oguri03a,lopes04,chen04}.  The discovery of the first such lens suggests that these statistics can be a practical tool to study structure formation in the universe.  The statistics of giant arcs are also known as a good probe of clusters \\citep*{bartelmann98,meneghetti01,molikawa01,oguri01,oguri03b, wambsganss04,dalal04,maccio04}, and in fact lensed arcs and quasars complement each other in several ways.  For instance, in lensed quasar surveys one first identifies source quasars and then checks whether they are lensed, while in searching for lensed arcs one selects massive clusters and then searches for lensed arcs in them.  In other words, surveys for arcs are biased toward high mass concentrations, while lensed quasars probe random lines of sight.  Clusters selected by the presence of lensed quasars could, in principle, differ from those selected as having giant arcs.  In addition, lensed quasars have three advantages over lensed arcs in statistical studies.  First, quasars can be regarded as point sources, while sources for arcs are galaxies whose intrinsic sizes and shapes are important but unobservable.  Second, the number and configuration of images in a quasar lens system is unambiguous. Third, the redshift distribution of arc sources is poorly known (and controversial; see \\citealt{oguri03b}; \\citealt{wambsganss04}; \\citealt{dalal04}), while the redshift distribution of quasars is well known.  In all previous analytic work on the statistics of large-separation lensed quasars, the lens objects were assumed to be spherical. However, in the CDM model dark halos are not spherical at all but triaxial \\citep[e.g,][hereafter JS02]{jing02}.  It is already known that triaxiality has a significant effect on the statistics of lensed arcs, from both analytic \\citep[][hereafter OLS03]{oguri03b} and numerical (\\citealt*{meneghetti03a}; \\citealt{dalal04}) points of view.  In the statistics of normal lensed quasars, triaxiality (or ellipticity) has been thought to mainly affect the image multiplicities, with only small changes to the total lensing probability \\citep*{kochanek96,keeton97,evans02,chae03,hk04}. However, that conclusion is based on nearly-singular isothermal lens models, and the situation may be quite different for the less concentrated mass distributions of the massive halos that create large-separation lenses.  Moreover, only triaxial modeling allows us to study image multiplicities, and to consider whether it is statistically natural that the first known large-separation lens is a quadruple.  The structure of this paper is as follows.  In \\S\\ref{sec:theory} we review the triaxial dark halo model and its lensing properties (originally presented by JS02 and OLS03).  In \\S\\ref{sec:prob} we show our general results, while in \\S\\ref{sec:sdss} we customize our predictions to the SDSS quasar sample.  We summarize our conclusions in \\S\\ref{sec:sum}.  Throughout the paper, we assume a $\\Lambda$-dominated cosmology with current matter density $\\Omega_M=0.3$, cosmological constant $\\Omega_\\Lambda=0.7$, dimensionless Hubble constant $h=0.7$, and normalization of matter density fluctuations $\\sigma_8=0.9$. ",
        "conclusions": "\\label{sec:sum}  The dark matter halos predicted by the CDM model are triaxial rather than spherical, which has a significant effect on the statistics of large-separation gravitational lenses.  Triaxiality systematically enhances the lensing probability by a factor of $\\sim$4 if dark halos have an inner density profile with $\\alpha=1$, or a factor of $\\sim$2 if $\\alpha=1.5$.  The effects may be even more dramatic at very large image separations ($\\theta \\gtrsim 100''$), although such lenses are very sensitive to the most triaxial halos and so the predictions are not as reliable.  Thus, triaxiality must be added to the list of important systematic effects that need to be included in calculations of large-separation lens statistics.  (Some of the other effects are the inner density profile and the shape of the distribution of concentration parameters, as found by previous studies of large-separation lens statistics using spherical halos.)  Triaxial modeling allows us to predict the image multiplicities for large-separation lenses.  We found that the multiplicities depend strongly on the density profile: for $\\alpha=1$, lenses are dominated by naked cusp image configurations; while for $\\alpha=1.5$, quadruple configurations are the most probable. Double lenses, which are dominant among normal arcsecond-scale lenses, are subdominant in both cases. Note that cusp lenses can be distinguished from doubles by the presence of a third image comparable in brightness to the other two, and by the configuration of image positions.  The differences can be ascribed to the different mass density profiles, and they indicate that the multiplicities of large-separation lenses will provide a qualitatively new probe of the central density profiles of massive dark matter halos (and hence a new test of CDM).  We have computed lensing probabilities and image multiplicities for the SDSS quasar sample.  We predict that for both $\\alpha=1$ and $1.5$ most of the large-separation lenses should be quadruples or cusps.  The fractions of quadruple lenses at separations of $\\theta=15''$ are $\\sim$0.2 and $\\sim$0.4 for $\\alpha=1$ and $1.5$, respectively.  Thus it is not surprising that the first large-separation lens discovered is a quadruple.  In addition, we computed the expected number of large-separation lenses in both SDSS and CLASS, and found that the data are consistent with the CDM model with $\\sigma_8 \\sim 1$, in agreement with other measurements \\citep[e.g.,][and references therein]{spergel03}. Thus, the discovery of SDSS J1004+4112 can be interpreted as additional support for CDM on non-linear cluster scales.  \\lastpagefootnotes The prediction that triaxial halos produce significant fractions of quadruple and cusp lenses should be kept in mind when considering samples of candidate large-separation lenses.  For example, \\citet{miller04} found six large-separation double lens candidates in the 2dF Quasar Redshift Survey, but no quadruple or cusp lens candidates.  Even accounting for small number statistics, our results suggest that such a high fraction of doubles would be inconsistent with CDM at more than 3$\\sigma$,\\footnote{If the predicted double fraction is $f_2$, then the Poisson probability of having $N$ doubles and no quadruples or cusps is $\\mathcal{L}(f_2) \\propto (f_2)^N$.} and that it would be surprising if many of the six candidates are genuine lens systems.  We note that the triaxial dark halo model we adopted in this paper can be improved in several ways. First, we assumed that the axis ratios of the triaxial ellipsoids are constant with radius. However, JS02 showed that the axis ratios decrease slightly toward the halo centers: $a/c$ decreases by $\\sim\\!0.2$ as the mean radius decreases from $\\sim\\!0.6 r_{\\rm vir}$ to $\\sim\\!0.06 r_{\\rm vir}$.  Since strong lensing is most sensitive to the inner parts of dark halos, it is possible that we have actually underestimated the effects of triaxiality on the statistics of large-separation lenses.  On the other hand, including baryons (which were neglected in the simulations of JS02) would tend to make dark halos rounder by $\\Delta(a/c) \\sim 0.1$--$0.2$ because of the isotropic gas pressure.  While our theoretical model is much more realistic than the simple spherical model, we have still made several simplifying assumptions.  One is that we have neglected substructure in dark halos.  The galaxies in massive cluster halos do not have a large effect on the statistics of lensed arcs \\citep{meneghetti00}, but it is not obvious whether or not they would affect large-separation lenses.  Substructure can affect the image multiplicities for isothermal lenses \\citep{cohn04}, so it should be considered for CDM halos as well.  Another effect we have neglected is the presence of a massive central galaxy in a cluster.  \\citet*{meneghetti03b} claim that central galaxies do not have a large effect on arc statistics.  However, because our results depend on the inner slope of the density profile, and a central galaxy effectively increases the concentration, this effect should be considered. A third phenomenon we have neglected is cluster merger events. Indeed, mergers can change the shapes of critical curves and caustics substantially, and thus have a great impact on lensing cross sections \\citep{torri04}.  To estimate the effect on large-separation lens statistics, we would need a realistic model of the cluster merger event rate and the physical conditions of merger events.  Addressing these various issues to improve the accuracy of the theoretical predictions is beyond the scope of this paper, but is certainly of interest for future work."
    },
    "0403/astro-ph0403319_arXiv.txt": {
        "abstract": "We present a study of variable faint optical sources discovered by comparing the Sloan Digital Sky Survey (SDSS) and the Palomar Observatory Sky Survey (POSS) catalogs. We use SDSS measurements to photometrically recalibrate several publicly available POSS catalogs (USNO-A2.0, USNO-B1.0, GSC2.2 and DPOSS). A piecewise recalibration of the POSS data in 100 arcmin$^{2}$ patches (one SDSS field) generally results in an improvement of photometric accuracy (rms) by nearly a factor of two, compared to the original data. In addition to the smaller core width of the error distribution, the tails of the distribution become much steeper after the recalibration. These improvements are mostly due to the very dense grid of calibration stars provided by SDSS, which rectifies the intrinsic inhomogeneities of Schmidt plates. We find that the POSS I magnitudes can be improved to $\\sim$0.15 mag accuracy, and POSS II magnitudes to $\\sim$0.10 mag accuracy. The smallest final errors are obtained with the GSC2.2 catalog for which they approach 0.07 mag at the bright end.  We use the recalibrated catalogs for the $\\sim$2,000 deg$^2$ of sky in the SDSS Data Release 1 to construct a catalog of $\\sim$60,000 sources variable on time scales 10-50 years. We find that at least 1\\% of faint optical sources appear variable at the $>$0.2 mag level, and that about 10\\% of the variable population are quasars, although they represent only 0.25\\% of all point sources in the adopted flux-limited sample. A series of statistical tests based on the morphology of SDSS color-magnitude and color-color diagrams, as well as visual comparison of images and comparison with repeated SDSS observations, demonstrate the robustness of the selection methods. Candidate variable sources are correctly identified in more than 95\\% of cases, and the fraction of true variables among the selected candidates is as high as 73\\%. We quantify the distribution of variable sources in the SDSS color-color diagrams, and the variability characteristics of quasars. The observed long-term variability (structure function) is smaller than predicted by the extrapolation of the power-law measured for short time scales using repeated SDSS imaging (0.35 vs. 0.60 mag for SDSS-POSS I, and 0.24 vs. 0.35 mag for SDSS-POSS II, rms). This turn-over in structure function suggests that the characteristic time scale for quasar variability is of the order one year. The long-term ($\\ga$1 year) quasar variability decreases with luminosity and rest-frame wavelength similarly to the short-term ($\\la$1 year) behavior. We also demonstrate that candidate RR Lyrae stars trace the same halo structures, such as the Sgr dwarf tidal stream, that were discovered using repeated SDSS observations. We utilize the POSS-SDSS selected candidates to constrain the halo structure in the parts of sky for which repeated SDSS observations do not exist. ",
        "introduction": "The time domain represents another dimension, in addition to spectral and spatial, in the exploration of celestial objects. Despite the importance of variability phenomena, the properties of optically faint variable sources are by and large unknown. There are about 10$^9$ stars brighter than $V=20$ in the sky, and at least 3\\% of them are expected to be variable at a few percent level (Eyer, 1999). However, the overwhelming majority are not recognized as variables even at the brightest magnitudes: 90\\% of variable stars with $V<12$ remain to be discovered (Paczy\\'{n}ski, 2000). Paczy\\'{n}ski (1996) lists striking examples of the serious incompleteness in the available samples of variable stars: eclipsing binaries of Algol type, and contact binaries (W UMa stars) are incomplete fainter than $V \\about 12$, and RS CVn type binaries are complete only to $V \\about 5$. Another vivid example of serious selection effects is the sky distribution of RR Lyrae stars: objects listed in the 4$^{th}$ General Catalogue of Variable Stars (the main resource for variable stars) are distributed in isolated square patches with the size and shape of the Schmidt plates used to discover them.  The discrepancy between the utility of variable stars and the available observational data has prompted several contemporary projects aimed at regular monitoring of the optical sky. The current state-of-the-art also greatly benefited from past and present microlensing searches (Paczy\\'{n}ski, 2001). Some of the more prominent surveys in terms of the sky coverage, depth, and cadence are:  \\begin{itemize} \\item The Faint Sky Variability Survey (Huber, Howell, Everett \\& Groot, 2002) is a very deep ($V < 24$) $BVI$ survey of 23 deg$^2$ of sky, containing about 80,000 sources \\item The QUEST Survey (Vivas et al. 2001) monitors 600 deg$^2$ of sky to a limit of $V=21$. \\item ROTSE (Akerlof et al. 2003) monitors the entire observable sky twice a night to  a limit of $V=15.5$. \\item OGLE (most recently OGLE III; Udalski et al. 2002) monitors \\about 100 deg$^2$ towards the Galactic bulge to a limit of $I=20$. Due to a very high stellar density towards the bulge, OGLE II has detected over 200,000 variable stars (Wo\\'{z}niak et al. 2002). \\end{itemize}  These and other surveys have demonstrated that in addition to variable stars, there are many other exciting photometrically variable objects in the sky. For example, ROTSE~I detected an optical flash generated by a gamma ray burst at a redshift of 1.6, the most luminous optical source ever measured ($V=9$, $M_V=-36.4$, Vestrand et al. 2002). The detection of such optical flashes may place strong constraints on the physical mechanisms responsible for gamma ray bursts. Similarly, the variability of quasars offers significant clues for the origin of their emission (e.g. Trevese, Kron \\& Bunone, 2001).  Recognizing the outstanding importance of variable objects, the last Decadal Survey Report highly recommended a major new initiative for studying the variable sky, the Large Synoptic Survey Telescope (LSST). The LSST\\footnote{There are currently two designs considered for implementation: a distributed aperture approach (Pan-STARRS, Kaiser et al. 2002) and a single large-aperture telescope (Tyson 2002).} will offer an unprecedented view of the faint variable sky: according to the current designs it will scan the entire accessible sky every three nights to a limit of $V\\about 24$. Compared to any other survey currently available, the data from LSST will be revolutionary. Yet, at least half a decade or more will elapse before the first photons are detected by the LSST. Meanwhile, the already available Palomar Digital Sky Surveys (POSS I and POSS II, for references see Appendix A) and ongoing Sloan Digital Sky Survey (see Section 2.2) can be used to study the properties of faint ($r \\sim 20$) optical sources, and here we present such a study.  The comparison of the POSS and SDSS surveys allows studies of the long term variability with time scales of up to half a century. By necessity, such studies are based on a small number of observations of the same objects to constrain the ensemble properties of a sample of sources, as opposed to studying well-sampled light curves for a small number of objects. The lack of detailed information for individual objects is compensated in some ways by the large sample size. In addition, the 5-band accurate SDSS photometry can be used for very detailed source classification; e.g. separation of quasars and stars (Richards et al. 2002), spectral classification of stars to within 1-2 spectral subtypes (Lenz et al. 1998, Finlator et al. 2000, Hawley et al. 2002), and even remarkably efficient color selection (thanks to accurate $u$ band photometry) of the low-metallicity G and K giants (Helmi et al. 2002) and horizontal branch stars (Yanny et al. 2000, Ivezi\\'{c} 2003a).   However, when using only a few epochs of observations, the robustness of variability detection critically depends on the stability of the photometric errors. While the SDSS photometric errors are small ($\\sim$0.02) and well behaved (Ivezi\\'{c} et al. 2003b), older photographic POSS data can have large errors (tenths of mag) due to intrinsic inhomogeneities of Schmidt plates and the lack of a sufficient number of calibration stars. This problem can be alleviated to some extent by using photometric measurements of stars in the SDSS to recalibrate POSS catalogs. In Section 2 we describe such a recalibration method and demonstrate that photometric errors in the POSS catalogs can be decreased by a factor of $\\sim$2 (rms), with a significant improvement in the behavior of the error distribution tails. In Section 3, we use SDSS data and recalibrated POSS catalogs to select variable objects in $\\sim$2000 deg$^2$ of sky from the SDSS Data Release 1 (Abazajian et al. 2003). In the same section we discuss tests which demonstrate the robustness of the selection algorithm, and quantify the distribution of variable sources in the SDSS color-color diagrams. The Milky Way halo structure traced by selected candidate RR Lyrae stars is discussed in Section 4, and in Section 5 we analyze the variability of quasars. Our main results are summarized in Section 6. ",
        "conclusions": "We present a direct comparison of photometric measurements available in public POSS catalogs. The most accurate photometry is provided by the GSC2.2 catalog. The results of photometric recalibration based on a dense grid of calibration stars measured by the SDSS demonstrate that errors in POSS photometry can be reduced by about a factor of two. POSS I magnitudes can be brought to $\\sim$0.15 mag accuracy, and POSS II magnitudes to $\\sim$0.10 mag accuracy. While these apparently irreducible errors are considerably larger than those delivered by modern CCD data ($\\sim0.02$), the POSS catalogs are, nevertheless, invaluable for studying sources variable on long time scales. A particular success of the recalibration method is that the resulting error distribution for POSS photometry is nearly Gaussian, which greatly helps in the design of robust algorithms for selecting candidate variable sources.  We designed and tested algorithms for selecting candidate variable sources using POSS and SDSS photometric measurements. The algorithm's decision to tag a source as a candidate variable is correct in more than 95\\% of cases. The selection criteria can be tuned up to result in samples that contain no more than 30\\% of spurious candidates.  This is the first study that examined the distribution of sources variable on long time scales in SDSS color-color diagrams. Even with the fairly large cutoffs for selecting candidate variables (0.20-0.35 mag), we find that at least 1\\% of faint optical sources appear variable. About 10\\% of the variable population are quasars, although they represent only 0.25\\% of all point sources (for $g<19$).  Using a sample of $\\sim$17,000 spectroscopically confirmed quasars, we demonstrate that the power-law increase of the quasar variability (structure function, rms) with time lag observed for short time lags cannot be extrapolated beyond a few years in the rest frame -- such extrapolation predicts a variability level significantly larger than measured (0.35 vs. 0.60 mag for SDSS-POSS I, and 0.24 vs. 0.35 mag for SDSS-POSS II). The implied turn-over in structure function indicates that the characteristic time scale for optical variability of quasars is of the order 1 year in the rest frame. The long-term ($\\ga$1 year) quasar variability decreases with luminosity and rest-frame wavelength similarly to the short-term ($\\la$1 year) behavior.  A particularly valuable result of comparing POSS and SDSS catalogs is the selection of candidate RR Lyrae stars, which are excellent probes of the Milky Way's halo structure. We demonstrated that the known halo substructures are recovered by the selected candidates, and discovered several new features. This method will eventually yield several thousand RR Lyrae candidates.  Our study also revealed some limitations of the POSS catalogs. In particular, we had to limit our search for variable sources to only isolated point sources detected by both POSS and SDSS. Attempts to find sources detected by only one survey, or variable sources that are blended with another nearby source, were unsuccessful due to overwhelmingly large numbers of false positives.  Despite these shortcomings, the assembled catalogs of candidate variable sources offer a good starting point for further analysis and follow-up observations. For example, light curves and spectra for selected candidates could be obtained even with telescopes of modest size. Such additional data would help improve candidates' classification beyond information provided by SDSS colors. Another potentially interesting research direction is positional cross-correlation with catalogs obtained at other wavelengths (e.g. ROSAT, 2MASS, IRAS). For example, long-period variables such as Mira and other AGB stars are typically strong infrared emitters, and thus could be efficiently separated from the rest of candidate variables.  This study once again demonstrates the importance of maintaining a careful archive of astronomical observations; the data may be valuable long after the acquisition technology becomes obsolete.   \\vskip 0.4in \\leftline"
    },
    "0403/cond-mat0403097_arXiv.txt": {
        "abstract": "Kinetics of collapse and  explosion transitions in microcanonical self-gravitating ensembles is analyzed. A system of point particles interacting via an attractive soft Coulomb potential and confined to a spherical container is considered. We observed that for 100--200 particles  collapse takes $10^3$ -- $10^4$ particle crossing times to complete, i.\\,e., it is  by 2-3 orders of magnitude slower than velocity relaxation. In addition, it is found that the collapse time decreases rapidly with the increase of the softcore radius. We found that such an anomalously long collapse time is caused by the slow energy exchange between a higher-temperature compact core and relatively cold diluted halo. The rate of energy exchange between the faster modes of the core particles and slower-moving particles of the halo is exponentially small in the ratio of the frequencies of these modes. As the softcore radius increases, and the typical core modes become slower,  the ratio of core and halo frequencies decreases and the collapse accelerates. Implications to astrophysical systems and phase transition kinetics are discussed. ",
        "introduction": "\\label{sec_intro}  Many groups of stellar systems have highly universal structures despite the apparent differences in their history and environment~\\cite{bt, astro}. These universal structures are thought to have arisen as a result of relaxation towards equilibrium or to otherwise long-lived states. A comparison between the age of the stellar systems with universal features and corresponding collisional relaxation times reveals, however, that several types of stellar systems, such as elliptical galaxies, have not existed long enough to be collisionally relaxed~\\cite{bt}. Other than collisional types of relaxation, such as ''violent relaxation'', or phase-space mixing caused by strong gravitational field fluctuations~\\cite{lb},  have been suggested to explain this apparent contradiction between the timescales. Yet a full understanding of the kinetics of relaxation in naturally occurring self-gravitating systems systems is still lacking.  A number of fairly idealized models have been analyzed to understand the nature of equilibrium  and transitory states of  stellar systems. A well-studied example  is an ensemble of self-gravitating particles particles with a sufficiently short-range small distance regularization confined in a container. That system exhibits a gravitational phase transition between a relatively uniform high-energy state and a low-energy state with a core-halo structure~\\cite{pr,ki2,ch1,bm2,usg,dv,ch,chs,chi,grv,pet,jm}. During such a transition in a microcanonical ensemble the system undergoes a discontinuous jump from a state that just ceases to be a local entropy maximum to a global entropy maximum state with the same energy but different temperature. A transition from a high-energy uniform state to a lower-energy core-halo state is usually called collapse. The reverse transition during which the core disappears is often referred to as explosion.  It has been recently observed in molecular dynamics (MD) simulations \\cite{jm} that a typical timescale for such gravitational transitions is paradoxically large, for a system of 125 -- 250 particles being of the order of $10^3$  relaxation times for a collapse and $10^2$ relaxtion times for an explosion. The relaxation time,  $t_r={R^{3/2} N^{1/2}}/{\\ln N}$, is the timescale of  typical particle velocity thermalization which proceeds mostly via soft Coulomb collisions~\\cite{bt}. It was also observed in Ref.~\\cite{jm} that the density relaxation, such as the formation of the core, advances relatively fast, while the evolution of the kinetic energy or temperature  proceeds noticeably slower~\\cite{jm}.  In this paper we undertake a more detailed study of kinetics of collapses in self-gravitating systems, in some way completing the investigation initiated in~\\cite{jm}. The structure of the paper is following: After this introduction we briefly outline the simulation setup and present the results for collapse kinetics. A section analyzing a slow core-halo energy transfer as the bottleneck of system relaxation follows. Conclusions and discussion of the results completes the paper. ",
        "conclusions": "In the previous two sections we obtained the following results for the kinetics of collapse from the uniform  to the core-halo state in self-gravitating systems: \\newcounter{ref} \\begin{list}{\\arabic{ref}.}{\\usecounter{ref} \\itemsep=0pt \\leftmargin=1cm} \\item The molecular dynamics simulations revealed that in a system of 100-200 particles the collapse time is by two-three order of magnitude longer than the velocity relaxation and strongly depends on the short-range part of the interaction potential. \\item It was found that the non-equilibrium feature with the slowest relaxation time is the temperature difference between the core and the halo. In contrast, such parameters as number of core particles and core kinetic energy evolve relatively fast. \\item A mechanism similar to the vibrational--translational relaxation in plasmas was suggested for the core -- halo energy exchange. For this mechanism we show that the core-halo thermalization is exponentially slow in the ratio of typical frequencies of the motion of core and halo particles. Despite several rather strong approximations used in our analysis, a theoretical estimate for the relaxation time is in a good agreement (less than by a factor of 2 off) with the simulation results. \\end{list}  So far nothing has been said about the reverse to collapse transition, i.\\,e., explosion, illustrated in Fig.~\\ref{fig_ec2}. \\begin{figure} \\includegraphics[width=.45\\textwidth]{fig5.eps} \\caption{\\label{fig_ec2} Plots of the relative values of (from top to the bottom) number of core particles $N'_{c}(\\t)$ (blue), total temperature  $\\T'(\\t)$ (red), and core core temperature $\\T_{c}(\\t)$ (black) vs $\\t$ for the explosion in system with $\\e=0.5$, $N=125$, and $x_0=0.005$} \\end{figure} It follows from this Figure that similarly to the collapse, the kinetic energy is the slowest-evolving quantity while the number of core particles and core kinetic energy lead the explosion. Since the fastest particles leave the core first, the core is always colder than the rest of the system. At the end of explosion the core becomes just a single cold particle  indistinguishable from any other particle of the system. This explains the sudden jump of the core temperature at the very end of explosion. Using Eq.~(\\ref{tr}) it is straightforward to explain why an explosion is faster than a collapse: Since the explosion starts at higher energy than collapse, the corresponding 'explosion core-halo state' has noticeably fewer core particles (twice for the considered here case) than the final collapse state. This reduces the exponential term in  Eq.~(\\ref{tr}). In addition, the total amount of energy which needs to be transferred between core and halo is smaller in case of explosion than in the case of collapse. It follows from Fig.~\\ref{fig_mf} where temperature jumps at collapse and explosion points are proportional to the energy exchanged between core and halo.    Are these results applicable to self-gravitating systems with other types of short-range regularization? For systems with continuous potentials at low-energies the particle motion near the equilibrium position are harmonic oscillations with a frequency roughly given by Eq.~(\\ref{fc}) where $r_0$ is of order of the core radius. Examples include ensembles with Fourier-truncated Coulomb potentials~\\cite{pet} or exchange interaction in systems with phase-space exclusion~\\cite {ch1}. For potentials with singular short-range part such as hard-core repulsion, the motion of core particles is discontinuous, yet a typical inverse timescale of such motion is given by  Eq.~(\\ref{fc}) as well. This follows from the fact that the expression Eq.~(\\ref{fc}) can be obtained by dividing a typical core particle velocity $\\sqrt{N_c/r_0}$ by the core radius $r_0$. The analysis and conclusion made in the Section III are based on a general principle that for a perturbation of a fast system by a slow one, the rate of energy transfer between these system is exponentially small in the ratio of their frequencies. And since the motion of core particles is always faster than the motion of the halo ones, we conclude that the results obtained here are qualitatively applicable to all collapsing self-gravitating systems independent of the form of the short-range cutoff.  An important implication of the demonstrated here essential non-isothermicity of collapsing system is about the applicability of the Smoluchowski equation description to these systems~\\cite{chs}. Compared to the more complete Boltzmann equation with Landau collision term (often called in case of self-gravitating systems Fokker-Plank-Vlasov equation~\\cite{ki3}), Smoluchowski equation offers a significant simplification: In the physically relevant case of spherically symmetric systems the solution depends only on one radial coordinate $r$. To make such description more realistic, one needs to take into account the non-uniform temperature field. This can be achieved by coupling the Smoluchowski and heat conduction equations. The non-isothermicity of evolution also indicates that the collapse in canonical ensemble must be radically different from its  microcanonical counterpart.  Finally few words about the astrophysical relevance of the obtained results. The structure and mere existence of the equilibrium core is a consequence of short-range regularization of the gravitational potential and the confining container which are the artifacts of the model. Hence the true equilibrium core-halo states never occurs in stellar systems. However, the observation that the collapse progress in hindered by the core  -- halo energy exchange which is exponentially slow in the ratio of the typical frequencies of motion of the core and halo particles remains applicable. Hence it is possible to interpret the states of the systems such as globular clusters whose age is noticeably larger than corresponding velocity relaxation (and therefore, initial core formation) time~\\cite{bt} as the transitory long-living states similar to core--halo thermalization states observed in this study."
    },
    "0403/astro-ph0403155_arXiv.txt": {
        "abstract": "We present the results of a long-term monitoring campaign of the \\has line in a sample of bright OB-supergiants (O7.5--B9) that aims at detecting rotationally modulated changes potentially related to the existence of large-scale wind structures. A total of 22 objects were monitored during 36 nights spread over 6 months in 2001--2002. Coordinated broad-band photometric observations were also obtained for some targets. Conspicuous evidence for variability in \\has is found for the stars displaying a feature contaminated by wind emission.  Most changes take place on a daily time-scale, although hourly variations are also occasionally detected. Convincing evidence for a cyclical pattern of variability in \\has has been found in 2 stars: HD 14134 and HD 42087 (periodic signals are also detected in other stars, but independent confirmation is required). Rotational modulation is suggested from the similarity between the observed recurrence time-scales (in the range 13--25 days) and estimated periods of stellar rotation. We call attention to the atypical case of HD 14134 which exhibits a clear 12.8-d periodicity both in the photometric and in the spectroscopic data sets. This places this object among a handful of early-type stars where one may observe a clear link between extended wind structures and photospheric disturbances. Further modelling may test the hypothesis that azimuthally-extended wind streams are responsible for the patterns of spectral variability in our target stars. ",
        "introduction": "One of the major legacies of the {\\it International Ultraviolet Explorer (IUE)} was to establish the ubiquity of UV line-profile variability in OB stars.  The pattern of variability predominantly consists of extra absorption features (the so-called 'discrete absorption components'; hereafter DACs) moving blueward across the (unsaturated) absorption troughs of P Cygni lines (e.g., Kaper \\etal 1996). The repetitive nature of DACs and the fact that the recurrence time-scale is commensurate with the (otherwise ill-defined) stellar rotational period suggest that rotation plays a key role in inducing the patterns of variability (e.g., Kaper \\etal 1999). UV time-series with a long temporal baseline (i.e., sampling several rotational periods) have revealed that OB stars exhibit a great diversity of behaviours that cannot always be strictly related to the stellar rotation period (e.g., Howarth, Prinja \\& Massa 1995; Prinja, Massa \\& Fullerton 1995; Prinja, Massa \\& Fullerton 2002). Nevertheless, this global phenomenon has been successfully interpreted in several cases in the framework of a model in which the changes are caused by the rotational modulation of large-scale, azimuthally-extended wind streams (e.g., Fullerton \\etal 1997).  The hydrodynamical simulations of Cranmer \\& Owocki (1996) suggest that large-scale photospheric perturbations (e.g., bright 'spots' causing a local enhancement of the radiative driving force) may lead to the development of these structures. Although the underlying physical mechanism remains elusive, pulsational instability can be invoked as a trigger of the wind variability (e.g., Fullerton, Gies \\& Bolton 1996). Large-scale or 'surface' magnetic field structures could be a viable alternative in view of the growing evidence for  (weak) magnetic fields in OB stars (e.g., Henrichs, Neiner \\& Geers 2003). Photospheric perturbations arising from these two phenomena are likely to have a significant impact on the global wind morphology (e.g., ud-Doula \\& Owocki 2002).  With the demise of {\\it IUE}, further progress on this issue is likely to stem in the foreseeable future from time-resolved studies of line variability in the optical. Firstly, intensive UV observations have been carried out over several rotation periods for few stars (Massa \\etal 1995; Prinja \\etal 1998; de Jong \\etal 2001). For practical reasons, an optical survey is much more amenable at collecting high-quality observations covering several cycles for a large, representative sample of OB stars. Secondly, in contrast to the UV resonance lines, optical recombination lines are (at least partly) formed in OB supergiants in the dense, strongly accelerating part of the outflow close to the stellar core. In virtue of their density-squared dependence, they are good probes of the large density gradients expected to prevail in the co-rotating interaction regions (Cranmer \\& Owocki 1996). Revealing a cyclical pattern of variability in these lines would give further credence to the idea that the origin of anisotropic outflows in OB stars is deep-seated and causally linked to processes taking place at the stellar photosphere. Here we discuss the temporal behaviour of \\ha, which is typically formed at $r$$\\sim$ 1.5 $R_{\\star}$ in O-type stars (e.g., Prinja, Fullerton \\& Crowther 1996). Although the line-formation region is likely to be more extended in B-type supergiants, this transition still probes the few inner radii of the stellar wind.  Previous surveys have been very useful in documenting \\has line-profile variability in OB stars (e.g., Ebbets 1982), but frequently suffered from a poor temporal sampling  hampering the study of the line-profile variations on a rotational time-scale (up to several weeks in the case of B supergiants). Despite the persistent and numerous searches for optical line-profile variability in OB stars, evidence for rotational modulation has been found in only few cases (e.g., Moffat \\& Michaud 1981; Kaper \\etal 1997, 1998). In the UV and optical domains the time-scales associated with the changes are often identical, within the uncertainties (e.g., Howarth \\etal 1995; Kaper \\etal 1999; de Jong \\etal 2001; Kaufer, Prinja \\& Stahl 2002). This suggests that \\has variability and DACs diagnose the same underlying physical phenomenon.  The long-term monitoring of the \\has line presented in this paper aims at better establishing the incidence of rotationaly modulated variability in OB supergiants. To potentially link the photospheric and wind activity, coordinated broad-band photometic observations were also carried out for a subset of this sample. Our survey primarily targets early B-type stars and can thus be regarded as complementary to others that concentrate on O-type (Kaper 1998; Kaper \\etal 1998) or late B- to early A-type supergiants (Kaufer \\etal 1996). The results presented in the following supersede previous preliminary reports (Morel \\etal 2004a, b). ",
        "conclusions": "Our long-term monitoring has confirmed the ubiquity of \\has line-profile variations in OB supergiants. The changes result from the combined effect of a varying amount of incipient wind emission and an intrinsically variable underlying photospheric profile, with the relative contribution of these two phenomena varying from one star to another. As a consequence of the large amplitude of the \\has variations observed, a precise determination of the stellar atmospheric parameters will not be generally possible from profile fitting of the Balmer lines in snapshot spectra. Taken at face value, for instance, the EW variations shown in Fig.3 would translate in the most extreme cases into a $\\pm$0.8 dex uncertainty in the derived mass-loss rate (Puls \\etal 1996; equation [43]). Although it remains to be seen whether a sophisticated modelling would result in drastically different figures, the large \\has changes intrinsic to the early-type supergiants may be one of the most important factors limiting the accuracy with which these objects can be used as extragalactic distance indicators (see discussion in Kudritzki 1999).  This survey has revealed the existence of cyclical \\has line-profile variations in HD 14134 and HD 42087. The low detection rate (2 cases out of 17) may have several explanations, including observational limitations or a loss of coherency in the patterns of variability. HD 14134 is of particular interest, as it is one of the very few early-type stars clearly exhibiting the same pattern of variability in the photometric and spectroscopic modes. The existence of azimuthally structured outflows in our program stars might explain the dramatic \\has changes, but firm conclusions must await detailed modelling (see Harries 2000 for the first effort in this direction). Sensitive searches for close companions (e.g., via long-baseline optical interferometric observations) are also needed to firmly rule out a variability arising from wind interaction.  We do not find evidence for a direct coupling between the wind activity and base perturbations arising from pulsations, but we caution that our data are not well suited to examine this aspect in detail. It is conceivable that in several instances the globally anisotropic nature of the outflow primarily results from the putative existence of small-scale magnetic loops anchored at the stellar photosphere (or alternatively a presumably fossil, large-scale dipole field). A major observational effort is currently undertaken to detect such weak magnetic fields in early-type stars (e.g., Henrichs \\etal 2003). Firm detections can be expected in the forseeable future. They would place the above conclusions on a quantitative footing and directly address the relevance of this scenario."
    },
    "0403/astro-ph0403363_arXiv.txt": {
        "abstract": "This paper presents and discusses the application of blind source separation to astrophysical data obtained with the WMAP satellite. Blind separation permits to identify and isolate a component compatible with CMB emission, and to measure both its spatial power spectrum and its emission law. Both are found to be compatible with the present concordance cosmological model. This application demonstrates the usefulness of blind ICA for cosmological applications. ",
        "introduction": "The detection of Cosmic Microwave Background (CMB) fluctuations, and the measurement of their spatial power spectrum, has been over the past three decades subject of intense activity in the cosmology community.  The CMB, discovered in 1965 by Penzias and Wilson, is a relic radiation emitted some 13 billion years ago, when the Universe was about 370.000 years old. This radiation emits, as a function of wavelength, as an almost perfect blackbody at a temperature of 2.726 Kelvin. Small fluctuations of this temperature tracing the seeds of the primordial homogeneities which gave rise to present large scale structures as galaxies and clusters of galaxies, are expected to exist, at the level of a few tens of micro Kelvin.  The importance of measuring these fluctuations to constrain cosmological scenarios describing the history and properties of our Universe is now well established.  The measurement of their statistical properties (and in particular of their spatial power spectrum) permits to drastically constrain the cosmological parameters describing the matter content, the geometry, and the evolution of our Universe \\cite{jungman96}.  The accuracy required for precision tests of the cosmological scenarios, however, is such that it is necessary to disentangle in the data the contribution of several distinct astrophysical sources, all of which emit radiation in the frequency range used for CMB observations \\cite{fb-rg99}. In addition, all measurements being noisy, it is necessary to identify, characterize, and remove as well as possible contributions due to instrumental noise.  CMB emission peaks in the millimeter-wave domain of the electromagnetic spectrum. In this wavelength range however, several astrophysical sources of radiation contribute to the total emission. Some of them are quite well known (or modelled), and others not so well. It is of utmost importance then, for the interpretation of the observations, to separate the various emissions. This can be achieved (to some extent) by the joint processing of various observations performed at various wavelengths.  To first order, the observation at a given wavelength $\\lambda$ can be modeled as a linear superposition of a number of sources. A sequence $\\{\\bX(x,y)\\}$ of observations  --- $J$ maps obtained at several wavelengths --- is modeled as an instantaneous mixture of $K$ independent sources contaminated by independent noise $\\bN(x,y)$: \\begin{displaymath} \\bX(x,y) = \\bA  \\bS(x,y) + \\bN(x,y) \\end{displaymath} where $\\bA$ is an unknown $J\\times K$ matrix.  As a strong prediction of the standard big bang model of cosmology, CMB anisotropies are expected to have a spectral emission law given, to first order, by the derivative (with respect to temperature) of the blackbody law, taken at the temperature of the CMB: \\begin{equation} \\Delta I = \\Delta T \\times \\left[ \\frac{\\partial B_\\nu(T)}{\\partial T} \\right]_{T=2.726\\,{\\rm K}} \\end{equation} where $B_\\nu(T)$ is the Planck law for the emission of a blackbody at temperature $T$. Blind component separation, which permits to estimate the emission law of sky components, provides a unique tool for checking this prediction. In the following, we apply a blind component separation method on WMAP data. ",
        "conclusions": "The SMICA blind component separation has been applied successfully to WMAP data. The analysis is performed on $\\ell$ varying from about 10 to 400, which encompasses all of the first acoustic peak of the CMB. We find that all 10 WMAP maps comprise a common astrophysical component identified here as CMB. The emission law of this first component is compatible with CMB fluctuations to within about one per cent, and its spatial power spectrum in the range $10 \\leq \\ell \\leq 400$ compatible with the measurement obtained by the WMAP team. The second component has an emission law proportional to $\\nu^{-2.7}$ except in the W (90 GHz) channel for which a significant excess is observed. This excess is thought to be due to dust emission contribution to the component. These results demonstrate the usefulness of blind component separation for astrophysical data analysis."
    },
    "0403/astro-ph0403680_arXiv.txt": {
        "abstract": "The current sample of high-redshift Supernova Type Ia, which combines results from two teams, High-$z$ Supernova Search Team and Supernova Cosmology Project, is analyzed for the effects of weak lensing. After correcting SNe magnitudes for cosmological distances, assuming recently published, homogeneous distance and error estimates, we find that brighter SNe are preferentially found behind regions (5-15 arcmin radius) which are overdense in foreground, $z\\sim 0.1$ galaxies. This is consistent with the interpretation that SNe fluxes are magnified by foreground galaxy excess, and demagnified by foreground galaxy deficit, compared to a smooth Universe case. The difference between most magnified and most demagnified SNe is about 0.3-0.4 mag. The effect is significant at $>99\\%$ level. Simple modeling reveals that the slope of the relation between SN magnitude and foreground galaxy density depends on the amount and distribution of matter along the line of sight to the sources, but does not depend on the specifics of the galaxy biasing scheme. ",
        "introduction": "The effects of weak gravitational lensing by the large-scale structure have been detected in several samples of high redshift QSOs, intermediate redshift galaxies, and BATSE GRBs. In the case of point sources, QSOs and GRBs, weak lensing manifests itself as angular (anti-)correlations between these sources and foreground inhomogeneously distributed mass \\citep{wf03,bsm01,wi98,bmg97}, while in the case of galaxies weak lensing is detected through its coherent shear effect (see \\cite{ref03} for a recent review). In principle, there is another, more direct way of detecting weak lensing, which uses fluxes of standard candles. If the observed magnitudes of standard candles are corrected for cosmological distances then the effect of lensing can be seen: brighter sources will lie behind regions of mass density excess, while fainter ones will have mass deficits in their foregrounds.  The best example of cosmological standard candle, Supernovae Type Ia (SNIa) have been extensively observed with the purpose of determining the global geometry of the Universe \\citep{tonry03,perlm99,gar98,riess98}.  Nuisance effects like evolution, variations in individual SN, and gray dust extinction have been studied theoretically and observationally, and have either been corrected for or shown to be small. Weak lensing, another nuisance effect has been addressed theoretically by several authors \\citep{amg03,met99,holz98,wcxo97} and found to be unimportant given the current uncertainties. For example, ~\\citet{wcxo97} used ray tracing through cosmological simulations and found that the lensing induced dispersions on truly standard candles are $0.04$ and $0.02$ mag at redshift $z=1$ and $z=0.5$, respectively, in a COBE-normalized cold dark matter universe with $\\Omega_m=0.4$, $\\Omega_\\Lambda=0.6$, $H_0=65$km/s/Mpc and $\\sigma_8=0.79$. These are small variations compared to the current errors which are $\\simgt 0.2$ mag.  Even though weak lensing effects are estimated to be small for $z_s <1$, they are predicted to be non-negligible for higher redshift sources, so it is not surprising that the highest redshift SNIa, SN1997ff at $z_s=1.755$ has been examined by several authors \\citep{ben02,mgg01,li01} for the effects of weak lensing due to galaxies along the line of sight.  Present day high-$z$ SNIa samples are dominated by lower redshift SNe, and so have not been examined for the effects of lensing. The main goal of this work is to determine if the observed fluxes of the cosmologically distant SNIa have suffered significantly from lensing induced (de-) amplifications.  \\vfill ",
        "conclusions": "\\label{disc}  We detect the signature of weak lensing in the current sample of 50 high redshift SNIa taken from two teams: High-$z$ Supernova Search Team and Supernova Cosmology Project. After correcting SNe magnitudes for cosmological distances (assuming \\citet{tonry03} values), we find that brighter SNe are preferentially found behind regions overdense in foreground galaxies. This $M_{SN}$--$N_>/N$ relation has a slope $\\beta=-0.3$ to $-0.4$, when the angular radii of foreground regions are 5-15 arcmin. The statistical significance is $>99\\%$ (see Fig.~\\ref{figthetatest} and ~\\ref{figmagtest}). The angular radii of 5-15 arcmin imply that the lensing structures are $1-3\\;{h_{0.7}}^{-1}$ Mpc across if they are located at a redshift of 0.1, and so correspond to non-linearly evolved intermediate-scale structure.  Aside from the possibility that the observed $M_{SN}$--$N_>/N$ relation is a fluke, there is one other possible, non-lensing interpretation: Galactic dust obscuration, which would make SNIa sources brighter and APM galaxies more numerous in the directions devoid of Galactic dust. If this interpretation is correct, then (1) Galactic extinction corrections applied to the SNe \\citep{tonry03} are much too small, and (2) Galactic dust is able to change projected APM galaxy number density by $50\\%$, since the observed rms in $\\delta n_{gal}$ is $\\approx 0.5$. We cannot comment on (1), but can estimate the effect of (2). How much excess/deficit in galaxy density can Galactic dust create? Using \\cite{sfd98} data we calculate the average $A_B$ Galactic extinction for the 50 SNe sample to be 0.18 mag. The slope of the R-band APM galaxy number counts is $d\\log N/dm=0.35$, therefore typical ``dust-induced'' $(1+\\delta n_{gal})=10^{A_B\\cdot d\\log N/dm}\\approx 1.16$. This is an upper bound because extinction in the $R$ band will be smaller than in $A_B$, and because extinction averaged over 5-15 arcmin radius patches will be smaller than estimates for individual SNe, given that \\cite{sfd98} extinction map resolution is 6$^\\prime$.1 FWHM.  If Galactic dust is indeed the cause of the $M_{SN}$--$N_>/N$ correlation, then $M_{SN}$ and $N_>/N$ should separately correlate with SNe extinction.  In reality, \\cite{sfd98} extinction estimates do not show a correlation with either $M_{SN}$ or $N_>/N$. Overall, Galactic dust interpretation of the observed $M_{SN}$--$N_>/N$ relation seems unlikely, but cannot be ruled out.  Note that dust {\\it intrinsic} to the $z_l\\sim 0.1$ structures probed by the APM galaxies cannot be invoked to explain the observations, because such dust would produce an effect opposite to the one detected here, i.e. $\\beta$ would be positive. If dust is present in groups and clusters traced by the APM, it will diminish the amplitude of the effect we detect. Presence of dust in groups was suggested by \\citet{bfs88} who found that faint QSO candidates are anti-correlated with foreground groups. However, such an anti-correlation can also be explained by weak lensing \\citep{rwh94,cs99}. Current observations indicate that groups and clusters do not contain significant amounts of dust \\citep{f93,nwm03}.  If the lensing interpretation in correct, then the standard models of mass distribution have some difficulty in reproducing $\\beta$; the observed value would be detected only in $\\simlt 3\\%$ of the cases. We investigate how $\\beta$ is affected by the amount and distribution of mass along the light of sight to the sources, and galaxy biasing schemes.  We find that larger optical depths and broader $\\kappa$ pdf result in steeper $\\beta$ slopes (Fig.~\\ref{figoptdepth}). Optical depth is a function of global geometry, and no realistic cosmological model can give $\\kappa_0\\simgt 0.05$, which is what would be required to comfortably explain the results. Broader $\\kappa$ pdfs mean that mass fluctuations are more extreme than the standard cosmological models allow, a scenario which is in apparant conflict with other means of determining mass fluctuations, like cosmic velocity flows and weak shear lensing.  We also find that biasing has little effect on $\\beta$ (Fig.~\\ref{grid}). This insensitivity to biasing is in contrast to weak lensing induced QSO-galaxy and GRB-galaxy (anti-)correlations, where biasing could, at least partly explain the higher than expected amplitude of the effect \\citep{jai03,wf03}. We conclude that weak lensing of standard candles provides a cleaner probe of the mass distribution at $z\\simlt 0.1$, on $\\sim$few Mpc scales, then lensing induced angular correlations.  In fact, SNIa can provide the perfect means of measuring mass inhomogeneities using gravitational lensing. A set of standard candles at known redshifts can be analyzed using the complementary techniques of weak magnification and weak shear lensing. The advantage of magnification lensing over the more commonly used shear lensing is that with the former one can chose the redshift range of the lenses, whereas the latter yields the cumulative effect of lensing along the entire line of sight to the source. Therefore a large uniform set of intermediate redshift SNIa, such as the ones that would result from the SNAP mission\\footnote{{\\tt http://snap.lbl.gov}} and the LSST mission\\footnote{{\\tt http://www.lsst.org/lsst$\\_$home.html}} would be invaluable for the studies of mass clustering in the nearby Universe. In addition to detecting intermediate and high-redshift SNe for the purposes of estimating the global cosmological parameters and the equation of state of the dark energy, SNAP will also measure weak shear lensing signature due to large scale structure. Combining magnification data (of the type considered in this paper) and shear information \\citep{rho03} for a large set of SNe will allow the study of mass distribution at $z\\simlt 1.5$ with unprecedented accuracy and 3-dimensional spatial resolution.  The final issue we address is the impact of weak lensing on the determination of global cosmological parameters, $\\Omega_m$ and $\\Omega_\\Lambda$ using SNIa standard candles. In principle, weak lensing can affect the derived values of $\\Omega_m$ and $\\Omega_\\Lambda$, if (1) demagnified SNe are preferentially lost from the sample due to faint flux cutoff, and/or (2) the $\\kappa$ pdf is asymmetric and the SN sample size is small (see also \\citet{wcxo97}). If demagnified SNe were lost from the sample then the distribution of SNe in $N_>/N$ would be skewed in the direction of larger $N_>/N$ values. If $\\kappa$ pdf is asymmetric, (and it is, on scales considered here), then most SNe in a small sample will be slightly demagnified compared to average, and will have their $N_>/N$ skewed in the direction of smaller values. So both (1) and (2) would make the distribution of SNe in $N_>/N$ uneven, and would to some extent cancel each other. In the present sample of 50, the distribution of SNe in $N_>/N$ is indistinguishable from uniform, so the average $M_{SN}$ corresponds to the average $N_>/N$, and hence lensing effects by $\\sim$Mpc-size structures probably did not bias the derived values of $\\Omega_m$ and $\\Omega_\\Lambda$."
    },
    "0403/astro-ph0403539_arXiv.txt": {
        "abstract": "The Energetic X-ray Timing Survey Telescope (EXIST) mission concept is under study as the Black Hole Finder Probe (BHFP), one of the three Einstein Probe missions in the Beyond Einstein Program in the current NASA Strategic Plan. EXIST would conduct an all-sky imaging hard X-ray ($\\sim$10-600 keV) survey with unprecedented sensitivity: about 5 $\\times 10^{-13}$ cgs over any factor of 2 bandwidth, or comparable to that achieved at soft X-rays in the ROSAT survey. The proposed angular resolution of 5arcmin, temporal resolution of 10microsec, energy resolution of 1-4 keV over the broad band, and duty cycle of 0.2-0.5 for continuous coverage of any source provide an unprecedented phase space for timing and spectral studies of black holes --from stellar to supermassive, as well as neutron stars and accreting white dwarfs. The large sky coverage allows intrinsically rare events to be studied. One particularly exciting example is the possible detection of tidal disruption of stars near quiescent AGN. Super flares from SGRs could be detected out to the Virgo cluster. The large duty cycle and all sky monitor nature of the mission will enable QPOs from luminous AGN and BH X-ray binaries to be studied on timescales not possible before. I provide an overview of the mission concept and Reference Design, the X-ray timing science prospects for EXIST, and how these might be further optimized in the current Study for EXIST as the BHFP so that EXIST might include many of the desirable features of a next-generation timing mission. ",
        "introduction": "At energies above $\\sim$10 keV, the sky becomes relatively dark as the thermal emission from the brightest X-ray binaries becomes less prominent and diffuse emission from hot gas can not be maintained. The number of bright sources decreases dramatically, but those that remain are likely to be among the most interesting (at least for high energy astrophysics) as they are either accretion powered or non-thermal. In both cases, variability often becomes increasingly prominent with increasing energy, at least up to 20-30 keV beyond which it has usually been difficult to measure (on short timescales) for persistent sources. These hard X-ray sources are, in general, accreting black holes: at fluxes \\Fx(10-20 keV) \\la1 \\X 10$^{-11}$ \\fcgs, the extragalactic AGN dominate the sky above 10 keV. But even near the galactic plane, the logN-logS relations derived from the RXTE/ASM data for \\Fx(2-12keV) \\ga 6 \\X \\X 10$^{-11}$ \\fcgs by Grimm, Gilfanov and Sunyaev (2002) (and discussed by Vrtilek in this volume) would suggest that black holes are \\ga10\\% of the total for source fluxes at these limits, and probably more so if instead measured at 10-20 keV. This inference requires including transients, and the simple fact that of the \\about180 galactic sources detected above the limiting \\Fx by the ASM, all 16 of the currently known black hole candidates (BHCs) in the Galaxy (cf. McClintock and Remillard 2003) would have been detected.  It could then be argued that a follow-on mission to RXTE, with its highly successful timing studies of accreting black hole (BH) and neutron star (NS) systems, would gain most from greatly increased sensitivity above 10 keV. Obviously, this is countered by the greater source fluxes and thus count rates at lower energies: for a simple power law (with photon index 2) approximation to the spectrum, the counts are 10\\X larger in a decade lower energy band. However the relatively non-variable source emission components (e.g. blackbody emission from NS surfaces or the thermal soft continuum from accretion disks around BHs) provide additional `signal backgrounds' at low energies, and indeed the rms amplitude of quasi-periodic oscillations (QPOs) in both BH and NS accretors increases sharply with energy up to at least \\about20 keV (cf. Kaaret, this volume). Thus on balance, the net gains with increased energy band are dependent on the timescale and spectrum of the variable source component(s) to be measured. QPOs can be detected in time T with detection significance n$_{\\sigma}$  (van der Klis 1998)  \\begin{equation} T = 4\\frac{n_\\sigma^2\\Delta\\nu}{r^4s^4}(b+s)^2 \\end{equation}  \\ni where s is the source count rate, b is the background, $\\Delta\\nu$ is the signal bandwidth, and r is the rms amplitude.  We shall use this expression to estimate the approximate sensitivity  for QPOs that could be detected at hard X-ray energies with the Energetic X-ray Imaging Survey Telescope (EXIST), now under Study as the proposed Black Hole Finder Probe (one of the 3 Einstein Probe missions) in the Beyond Einstein Program under the current NASA Strategic Plan. In making these estimates, and others for hard X-ray timing science possible with EXIST, we make use of the so-called ``Reference Design'' for the mission. We first introduce the mission concept and Reference Design and then discuss the principal areas of (hard) X-ray Timing that could be advanced with EXIST. These include: QPOs for bright galactic sources and their long-term trends, QPOs for AGN, stellar encounter events (both stellar disruption and star-disk interactions) with dormant AGN, and super-flare events from soft gamma-ray repeaters (SGRs) in galaxies out to the Virgo cluster. We then describe some of the modifications to the Reference Design which are being investigated in the current BHFP Study and which could further optimize EXIST as a next generation timing mission. ",
        "conclusions": "EXIST in its current Reference Design configuration already incorporates several key attributes of an optimized hard X-ray timing mission. The X-ray Timing mission or observatory concepts discussed as the needed successor for RXTE are based in large part on the requirements to resolve individual pulse trains in QPOs from galactic BH and NS accretors. EXIST would not do this in its Reference Design but could, in principle, with some of the options being studied now by the EXIST team for the BHFP Study. However EXIST, in its primary survey mode configuration is already a superb ASM and would have sensitivity to conduct pathfinder science on the (much) longer timescales needed for variability studies of galactic BHs and AGN. Break frequencies in the power spectra of AGN, and QPOs in AGN, could be measured by EXIST over a range of interesting BH masses.  Of perhaps the most interest for novel timing studies of SMBHs is the suitability of EXIST to provide an unbiased (by low energy extinction) survey for stellar disruption near SMBHs that almost certainly lurk in bulges of all galaxies, without requiring these be known or detected as persistent sources by virtue of being AGN. Dormant AGN, or long since extinct AGN, are available for this study in sufficient numbers that surprisingly large rates of detection are predicted. More detailed study is needed to assess whether a disruption event will indeed produce an accompanying hard X-ray spectral component (as with galactic BH transients), but the recent XMM detection (Komossa et al 2004) of such a component in the faded ``afterglow'' of a 1992 likely stellar disruption discovered with ROSAT makes the hard X-ray component even more likely -- particularly if this XMM spectrum is later observed to decay with the $\\sim\\tau^{-1.1}$ possibly expected.    \\begin{theacknowledgments} I thank the EXIST Science Working Group for discussions and Roger Blandford and Martin Elvis for discussions of AGN variability and stellar disruption. I thank Bill Craig, Neil Gehrels, Fiona Harrison, Jaesub Hong, Ron Ticker and GSFC engineers for their work on developing the EXIST Reference Design. This work was partly supported by NASA grant NAG5 5-5396.   \\end{theacknowledgments}"
    },
    "0403/astro-ph0403013_arXiv.txt": {
        "abstract": "{The variability and rotation of ultra cool dwarfs (UCDs) provide important information on the atmospheres and evolution of these very low mass stars and brown dwarfs.  As part of an ongoing program to investigate this, the projected rotation periods, \\vsini, derived from high resolution VLT/UVES spectroscopy via cross correlation are presented for 16 field UCDs (M9V--L7.5V). This doubles the number of L dwarfs for which \\vsini\\ has been measured.  All targets are found to have \\vsini\\ between 10 and 40\\,\\kms\\ confirming that L dwarfs are rapid rotators.  Radial velocities have also been measured to a precision of 1--2\\,\\kms.  From the random distribution of the rotation axes, $i$, and theoretically predicted radii, one-sided confidence intervals are placed on the rotation periods of individual objects. These are compared with published period data obtained from photometric monitoring programs.  From this, the period of 31\\,hrs for the L0.5 dwarf \\object{2M0746+2000} published by Gelino et al.\\ (\\cite{gelino02}) may be ruled out as the rotation period. The period of 11.2\\,$\\pm$\\,0.8\\,hrs for the L1.5 dwarf \\object{2M1145+2317} obtained by Bailer-Jones \\& Mundt (\\cite{bjm01}) is consistent with the present \\vsini\\ results so is plausibly the true rotation period.  The inclination of the rotation axis is constrained to be $i$=62\\deg--\\,90\\deg\\ with an expectation value of 76\\deg.  Alternatively the data set a lower limit on the radius of 0.1\\rsol, which is within the range of radii predicted by models for brown dwarfs older than 0.5\\,Gyr.  Similarly, the period of 2.7\\,$\\pm$\\,0.1\\,hrs detected by the same authors for \\object{2M1334+1940} is also confirmed as the likely rotation period; the inclination is $i$=27\\deg--\\,44\\deg ($<i>=34$\\deg).  Where no variability or period was detected by the monitoring programs the likely reason is low contrast modulating surface features.  However, in three cases variability but no period was detected, even though the likely rotation period range inferred from \\vsini\\ lies within the timescale to which the monitoring was sensitive.  This reinforces the `masking hypothesis' of Bailer-Jones \\& Mundt (\\cite{bjm01}), the idea that the evolution of photospheric features on timescales shorter than the rotation period obscure the regular modulation of the light curve. As has been previously discussed, a likely candidate for such features is inhomogeneous dust clouds. ",
        "introduction": "Understanding the nature of low mass stars and brown dwarfs requires a detailed characterization of their observable atmospheres. Specifically, the determination of luminosities, chemical compositions and ages ultimately depends on adequate knowledge of the radiative and convective transport mechanisms.  While our understanding of these mechanisms is reasonable (if incomplete) for some types of stars, it is comparatively poor for ultra cool dwarfs (UCDs), late type M, L and T dwarfs. This is not least because at the low effective temperatures involved ($<$3000\\,K), solid dust particles form in significant numbers and varieties, and these have a major impact on the structure of the atmosphere and hence on the observable properties of the star (e.g.\\ Allard et al.\\ \\cite{allard01}; Burrows \\& Sharp \\cite{burrows99}). In particular, global dust structures (`clouds') may play a significant role.  Significant information on UCDs can be obtained from their observable temporal changes,  and several dedicated monitoring programs have now uncovered good evidence for optical and infrared variability (Bailer-Jones \\& Mundt \\cite{bjm99}; Terndrup et al.\\ \\cite{terndrup99}; Tinney \\& Tolley \\cite{tinney99}; Nakajima et al.\\ \\cite{nakajima00}; Bailer-Jones \\& Mundt \\cite{bjm01}; Martin et al.\\ \\cite{martin01}; Bailer-Jones \\cite{bj02}; Burgasser et al.\\ \\cite{burgasser02}; Clarke et al.\\ \\cite{clarke02a}, \\cite{clarke02b}; Gelino \\cite{gelino02}; Gelino et al.\\ \\cite{gelinoetal02}; Bailer-Jones \\& Lamm \\cite{bjl}; Clarke et al.\\ \\cite{clarke03}; Enoch et al.\\ \\cite{enoch03}; Joergens et al.\\ \\cite{joergens03}; Koen \\cite{koen03}; Zapatero Osorio et al.\\ \\cite{zapatero03}) and H$\\alpha$ variability (Hall \\cite{hall02}; Mochnacki et al.\\ \\cite{mochnacki02}; Liebert et al.\\ \\cite{liebert03}).  Of the 80 or so UCDs monitored by these groups, about 30 show evidence for photometric variability. While this figure depends on what one takes as sufficient evidence for variability, at least half are convincingly variable on the scale of a few tens of millimagnitudes. There are indications that this photometric variability is not caused by a simple rotational modulation of a non-uniform photosphere. One candidate is an inhomogeneous dust layer evolving under the influence of convection on time scales shorter than the rotation period (see the introduction to Bailer-Jones \\cite{bj02} for a discussion).  Understanding this phenomenon is important, not only for what it tells us about atmospheric process under these interesting physical conditions, but also for the implications it has concerning conclusions drawn from single epoch observations of UCDs.  In this article I address the specific issue of how the photometric variability and period data relate to UCD rotation, as measured by line broadening from high resolution spectroscopy. This provides the {\\it projected} rotation speed, \\vsini, so using simple statistical arguments and theoretical predictions for UCD radii, constraints can be placed on likely rotation periods. These (plus other published \\vsini\\ data) are compared with the variability data to confirm or refute possible rotation periods and investigate the `masking hypothesis'.  Previously published work has established that UCDs are fast rotators: of the 17 hitherto observed L dwarfs, 16 have \\vsini\\ in the range 10--40\\,\\kms\\ and one (Kelu-1) has 60\\,\\kms\\ (Basri et al.\\ \\cite{basri00}; Mohanty \\& Basri \\cite{mohanty03}; Schweitzer et al.\\ \\cite{schweitzer01}).  In the present work, \\vsini\\ values are determined for 16 UCDs (15 L dwarfs and 1 M dwarf), 14 of which have no previously published measurement.  Moreover, conservative ranges for \\vsini\\ are established which reflect the dominant uncertainties in measuring \\vsini\\ by cross correlation. For the majority of UCDs, the complete \\vsini\\ range lies between 10 and 30\\,\\kms, and there is no measurement below about 10\\,\\kms, thus confirming the rapid rotation of UCDs.  For a discussion of UCD rotation and its relation to chromospheric activity and possible dynamo mechanisms, the reader is referred to Mohanty \\& Basri (\\cite{mohanty03}). ",
        "conclusions": "I have presented \\vsini\\ measurements for 16 ultra cool field dwarfs with spectral types between M9 and L7.5. 14 of these are new, doubling the number of L dwarfs with measured \\vsini.  14 UCDs have \\vsini\\ between 10 and 30\\,\\kms, the other two between 30 and 40\\,\\kms. This confirms a previously published result that, compared to the M dwarfs, L dwarfs are fast rotators, with implications for their angular momentum evolution.  In addition to establishing mean \\vsini\\ values I determined a conservative \\vsini\\ range for each UCD which encompasses the asymmetric uncertainties in fitting the cross correlation peak: this full range is typically 7\\,\\kms.  With theoretically predicted radii of 0.1\\,\\rsol, the expected rotation periods of these UCDs (i.e.\\ using $<\\sin i> =\\pi/4$) is 3--10\\,hrs.  The measured \\vsini\\ determines the maximum rotation period. Using the statistical distribution of $i$, a one-sided 90\\% confidence range on the period of each UCD can be established: this is 5.3--12.1\\,hrs and 1.8--4.0\\,hrs for \\vsini=10\\,\\kms\\ and 30\\,\\kms\\ respectively. A larger radius for a given \\vsini\\ would imply longer periods.  From this analysis, suggested periods for UCDs from published monitoring programs were assessed. The following conclusions were drawn.  The period of 11.2\\,$\\pm$\\,0.8\\,hrs for 2M1145+2317 detected by Bailer-Jones \\& Mundt (\\cite{bjm01}) is likely to be the rotation period for this L1.5 dwarf, and the inclination angle is constrained to be $i$=62\\deg--90\\deg, i.e.\\ a near edge-on rotator. Likewise, the period of 2.7\\,$\\pm$\\,0.1\\,hrs detected for the L1.5 dwarf 2M1334+1940 by the same authors is confirmed as the likely rotation period, and the inclination is $i$=27\\deg--44\\deg. The period of 31\\,hrs detected by Gelino et al.\\ (\\cite{gelino02}) for the brighter component of the binary 2M0746+2000 can be strongly ruled out as a rotation period.  In three cases (2M0913+1841, 2M1145+2317 and SD1203+0015), variability but no period was detected by monitoring programs, even though the present work derives likely periods which should have been detectable by the programs. Although all of these objects are faint, this seems to support the `masking hypothesis', the idea that surface features evolving in distribution and/or brightness faster than the rotation period could mask such a period from being detected by monitoring. The challenge to ongoing work remains to characterise this evolution."
    },
    "0403/astro-ph0403108.txt": {
        "abstract": "{   We report the results of a high-resolution spectroscopic survey of all the stars more luminous than $M_V = 6.5$ mag within 14.5 pc from the Sun. The {\\it Hipparcos} catalog's completeness limits guarantee that our survey is comprehensive and free from some of the selection effects in other samples of nearby stars. The resulting spectroscopic database, which we have made publicly available, includes spectra for 118 stars obtained with a resolving power of $R \\simeq 50,000$, continuous spectral coverage between $\\sim 362-921$ nm, and typical signal-to-noise ratios in the range $150-600$. We derive stellar parameters and perform a preliminary abundance and kinematic analysis of the F-G-K stars in the sample. The inferred metallicity ([Fe/H]) distribution is  centered at about $-0.1$ dex, and shows a standard deviation of $0.2$ dex. A comparison with larger samples of {\\it Hipparcos} stars, some of which have been part of previous abundance studies, suggests that our limited sample is representative of a larger volume of the local thin disk. We identify a number of metal-rich K-type stars which appear to be very old, confirming the claims for the existence of such stars in the solar neighborhood. With atmospheric effective temperatures and gravities derived independently of the spectra, we find that our classical LTE model-atmosphere analysis of metal-rich  (and mainly K-type) stars provides discrepant abundances from neutral and ionized lines of several metals. This ionization imbalance could be a sign of departures from LTE or inhomogeneous structure, which are ignored in the interpretation of the spectra. Alternatively, but seemingly unlikely, the mismatch could be explained by systematic errors in the scale of effective temperatures. Based on transitions of majority species, we discuss abundances of 16 chemical elements. In agreement with earlier studies we find that the abundance ratios to iron of Si, Sc, Ti, Co, and Zn  become smaller as the iron abundance increases until approaching the solar values, but the trends reverse for higher iron abundances. At any given metallicity, stars with a {\\it low} galactic rotational velocity tend to have high abundances of Mg, Si, Ca, Sc, Ti, Co, Zn, and Eu, but low abundances of Ba, Ce, and Nd. The Sun appears deficient by roughly 0.1 dex in O, Si, Ca, Sc, Ti, Y, Ce, Nd, and Eu, compared to its immediate neighbors with similar iron abundances. ",
        "introduction": "The study of the stellar population of the solar neighborhood is of interest for  a wide range of investigations.  Nearby stars are  the only ones for which we can obtain crucial measurements,  such as parallaxes or angular diameters. They are typically the best candidates for very precise spectroscopic followups, such  as searches for lower-mass companions (stellar,  brown dwarf, or planetary), circumstellar disks, or solar-like cycles. Their proximity involves obvious advantages that should be exploited.  From a wider perspective, the stellar population of the solar neighborhood may be representative  of larger parts of the Galactic disk. Early studies revealed that the  simplest models of Galactic chemical evolution could not match the observed metallicity distribution -- the so-called G-dwarf problem (van den Bergh 1962). Recent investigations have shown that more complex scenarios, such as those including inhomogeneous evolution (Malinie et al. 1993) or infall (Chiappini et al. 1997), more closely resemble the observations. However, almost all the derived metallicity distributions of samples approximately limited in volume, have been obtained from photometry (Twarog 1980; Rocha-Pinto \\& Maciel 1996, 1998a; Flynn \\& Morell 1997; Kotoneva et al. 2002). This is worrying, as it has been recognized that some of the employed photometric indices may be seriously affected by chromospheric activity (Rocha-Pinto \\& Maciel 1998b; Favata et al. 1997a) and other systematic effects (Haywood 2002; Reid 2002; Twarog, Anthony-Twarog \\& Tanner 2002). These uncertainties can be avoided by using spectra.  Favata et al. (1997b)  carried out the only existing spectroscopic study of a volume-limited sample of nearby stars, and found a metallicity distribution for K dwarfs which is remarkably more metal-rich than those obtained through photometric analyses. %Their result %has been very controversial, and it has been suggested % that the discrepancy might be traceable to % poor statistics (Rocha-Pinto \\& Maciel 1998a).   We have conducted a complete $M_V$-limited volume-limited survey of nearby stars in an attempt to settle these and other issues. In addition to the Catalogue of Nearby Stars (Gliese \\& Jahrei\\ss\\ 1991, hereafter CNS3) -- used by most of the previous studies-- we have taken advantage of the  {\\it Hipparcos} catalog (ESA 1997). {\\it Hipparcos} is complete down to $V=7.3$\\footnote{This is a lower limit; the actual values depend on galactic latitude and spectral type.}. Restricting the study to stars brighter than $M_V = 6.5$ ($\\simeq$ K2V), {\\it Hipparcos}' completeness reaches out to 14.5 pc. Figure \\ref{cns3} shows that  nearly 25 \\% of the stars included in CNS3 within 14.5 pc from the Sun (parallaxes larger than 69.18 mas) are in fact farther away according to the accurate parallaxes measured by {\\it Hipparcos}.  \\begin{figure}[ht!] \\centering {\\includegraphics[width=5.5cm,angle=90]{./cns3.ps}} \\caption{Comparison of the parallaxes larger than 69.18 mas included in  CNS3, against {\\it Hipparcos}. {\\it Hipparcos} has  revealed that a significant fraction of these stars are indeed further away from the Sun. The solid straight line has a slope of one.} \\label{cns3} \\end{figure}  \\begin{figure}[ht!] \\centering {\\includegraphics[width=5.5cm,angle=90]{./f3n.ps}} \\caption{Distribution of absolute ($M_V$; solid line) and visual ($V$; dashed line) magnitudes of the sample. Approximate spectral types for dwarfs are assigned to the different ranges of $M_V$ as reference.} \\label{dist} \\end{figure}   High-dispersion high signal-to-noise optical spectra were acquired for 118 stars: all the objects  more luminous than $M_V=6.5$ within 14.5 pc from the Sun. Figure \\ref{dist} shows the sample's distribution of absolute and visual $V$ magnitudes: $M_V$ (solid line) and $V$ (dashed line). As a reference, spectral types for dwarfs are assigned to  different ranges in $M_V$. This sample, as a result of the dependence of stellar lifetime on mass and the shape of the initial mass function, is dominated by the coolest stars within our range of spectral types. This should be borne in mind when comparing to most other high-resolution spectroscopic studies of disk stars, which tend to select warmer targets.  We provide a high-quality spectroscopic archive of these {\\it nearest} stars. The data acquisition, reduction and archival, are the subject of Section \\ref{observations}. Many properties of these stars have been compiled: radial velocities, parallaxes, and proper motions. Other properties have been determined afresh: stellar atmospheric parameters, kinematics, ages, and chemical abundances. The resulting catalog is described in \\S \\ref{catalog}. Section \\ref{balance} examines the consistency of the inferred abundances. Section \\ref{discussion} discusses our findings. Section \\ref{extended} compares the kinematics and metallicities of the nearest stars with larger samples of thin-disk stars, and \\S \\ref{wrapping} presents a short summary and provides suggestions for future work. ",
        "conclusions": "\\label{wrapping}   Our spectroscopic analysis of the closest stellar neighbors reveals several interesting features. The derived [Fe/H] distribution for our sample is centered at [Fe/H] $\\simeq -0.1$, a value slightly  higher  than in most previous studies, which concentrated mainly on earlier spectral types. For a subsample of stars, comparison with evolutionary models can provide age estimates with moderate or small uncertainties. This reveals a number of metal-rich stars that appear to be as old as the most metal-poor stars in the sample ($\\ga 10$ Gyr), and older than most solar-metallicity thin-disk stars. These old metal-rich stars show thin-disk kinematics, and disrupt a mild correlation between age and metallicity suggested by the stars with [Fe/H]$<0$.  We find that our LTE abundance analysis does not satisfy the ionization balance for several elements. The discrepancies grow with metallicity, and appear to decrease with effective temperature, reaching extreme levels (nearly 0.5 dex) for the coolest and most metal-rich stars. Similar discrepancies were previously reported by Feltzing \\& Gustafsson (1998) for several K-type metal-rich field dwarfs as well as Schuler et al. (2003) and Yong et al. (2004) for K-type members of the open clusters M34 and the Hyades, respectively. Systematic errors in the trigonometric gravities can be ruled out, as satisfying the ionization balance by modifying the gravity would require lower values by as much as 0.8 dex.  Systematic errors in the $T_{\\rm eff}$ appear unlikely, as they would need upward revisions of up to $\\sim 300$ K. We should emphasize that i) the sample of stars used by Alonso et al. (1996) to build the photometric calibrations employed here reach  metallicities as high as [Fe/H]$=+0.5$ dex, covering adequately the range of (Fe I) abundances of our stars, ii)    there is good agreement between the $T_{\\rm eff}$s derived from Str\\\"omgrem photometry and the $(B-V)$ color index, and iii)  the temperatures derived from the IRFM calibrations and from fitting the wings of Balmer lines are also in good agreement. The only requirement of the IRFM from model atmospheres is to correctly predict the IR continuum flux. The fact that fluxes for the stars used in the Alonso et al. calibrations were obtained with the Kurucz (1993) model atmospheres, which have been tested here and show %as well as the MARCS models the ionization imbalance, indicates an inconsistency between observations and the predictions of such models. We recall that adopting other families of LTE line-blanketed homogeneous model atmospheres is of no help. The problem is probably connected to some of  the assumptions involved in the modeling (maybe the {\\it usual suspects}: LTE  and surface convection). A direct comparison, however, of high-accuracy spectrophotometry with model fluxes should be carried out, before it is possible to definitely state that the inconsistencies are not related to uncertainties in the $T_{\\rm eff}$ scale. Further study of open clusters, double, and multiple systems will help to clarify the source of the  problem described.   Despite the the ionization problem, we use spectral lines of species that are dominant in K-type atmospheres to investigate abundance ratios to iron for a number of chemical elements. Our study suggests that the $\\alpha$ elements and Sc are enhanced relative to iron for metal-rich stars. In an analysis strictly differential to the Sun, we find that some elements' ratios to iron are higher than solar at [Fe/H] $= 0$. This effect is detected, for instance, for most  $\\alpha$ elements and amounts roughly to $0.1$ dex. This fact could be connected to, if not explained by, the Sun being somewhat older than most other dwarfs with similar spectral types in the thin disk -- a conclusion that we reach from examining an extended sample of {\\it Hipparcos} stars with metallicities from the literature.   Our efforts to obtain high-quality spectra for a distance- and $M_V$-limited sample of stars, have resulted in the identification of some serious problems in classical abundance analyses, rather than the very accurate relative abundances we hoped for. The spectra used in this work have been placed in a public archive. We propose this dataset as a standard test for modeling techniques to pass before the retrieved abundances are to be fully trusted."
    },
    "0403/astro-ph0403225_arXiv.txt": {
        "abstract": "The first massive astrophysical black holes likely formed at high redshifts ($z\\gsim 10$) at the centers of low mass ($\\sim 10^6~\\msun$) dark matter concentrations. These black holes grow by mergers and gas accretion, evolve into the population of bright quasars observed at lower redshifts, and eventually leave the supermassive black hole remnants that are ubiquitous at the centers of galaxies in the nearby universe.  The astrophysical processes responsible for the formation of the earliest seed black holes are poorly understood.  The purpose of this review is threefold: (1) to describe theoretical expectations for the formation and growth of the earliest black holes within the general paradigm of hierarchical cold dark matter cosmologies, (2) to summarize several relevant recent observations that have implications for the formation of the earliest black holes, and (3) to look into the future and assess the power of forthcoming observations to probe the physics of the first active galactic nuclei. ",
        "introduction": "It seems established beyond reasonable doubt that some supermassive black holes (SMBHs) were fully assembled early in the history of the universe. The handful of bright quasars at $z\\gsim 6$ are likely powered by holes as massive as $\\sim10^9~\\msun$, and the spectra and metallicity of these objects appear remarkably similar to their counterparts at moderate redshifts (Fan et al.\\ 2003).  Indeed, if one selects individual quasars with the same luminosity, their properties show little evolution with cosmic epoch.\\footnote{A possibly important exception is tentative evidence for increasing Eddington ratios towards higher redshifts, as discussed in \\S~\\ref{zoltansubsec:sdss}.}  This implies that the behavior of individual quasars is probably determined by local physics near the SMBH and is not directly coupled to the cosmological context in which the SMBH is embedded.  However, it is clear that the quasar population as a whole does evolve over cosmic timescales.  Observations from $0\\lsim z \\lsim 6$ in the optical (e.g., the Anglo-Australian Telescope's Two Degree Field, or 2dF, and the Sloan Digital Sky Survey, or SDSS) and radio bands (Shaver et al.\\ 1994) show a pronounced peak in the abundance of bright quasars at $z\\approx 2.5$. Recent X-ray observations confirm the rapid rise from $z=0$ towards $z\\approx 2$ for X-ray luminous sources ($L_X>10^{44}$~ergs~s$^{-1}$; Barger et al.\\ 2003) but have not shown evidence for a decline at still higher redshifts (Miyaji et al.\\ 2000).  The cosmic evolution of quasar black holes between $0 \\lsim z \\lsim 6$ is likely driven by a mechanism other than local physics near the hole.  This is reinforced by the fact that the timescale of activity of individual quasars is significantly shorter than cosmic timescales at $z\\lsim 6$, both on theoretical grounds ($\\sim 4\\times10^7$ yr, the e-folding time for the growth of mass in a SMBH, whose accretion converts mass to radiation with an efficiency of $\\epsilon=\\dot M c^2/L_{\\rm Edd}\\sim 10\\%$) and is limited by its own [Eddington] luminosity), and using the duty cycle of quasar activity inferred from various observations ($\\sim 10^7$ yr, e.g., Martini 2004 and references therein; see also Haiman, Ciotti \\& Ostriker 2004).  In the cosmological context, it is tempting to link the evolution of quasars with that of dark matter halos condensing in a Cold Dark Matter (CDM) dominated universe, as the halo population naturally evolves on cosmic timescales (Efstathiou \\& Rees 1988).  Indeed, this connection has proven enormously fruitful and has resulted in the following broad picture: the first massive astrophysical black holes appear at high redshifts ($z\\gsim 10$) in the shallow potential wells of low mass ($\\lsim 10^8~\\msun$) dark matter concentrations. These black holes grow by mergers and gas accretion, evolve into the population of bright quasars observed at lower redshifts, and eventually leave the SMBH remnants that are ubiquitous at the centers of galaxies in the nearby universe.  Nevertheless, many uncertainties about this scenario remain. Most importantly, the astrophysical process(es) responsible for the formation of the earliest seed black holes (and indeed for the presence of SMBHs at all redshifts) remain poorly understood.  In this review, we focus on the emergence of the first generation of black holes, though many of the important questions are quite general and apply equally to subsequent generations of black holes.  This review is organized as follows. In \\S~\\ref{sec:observations}, we summarize several relevant recent observations that have implications for early black holes.  In \\S~\\ref{sec:first}, we describe theoretical expectations for the formation and growth of these black holes within the paradigm of hierarchical CDM cosmologies and also discuss early black holes in the context of cosmological reionization.  In \\S~\\ref{sec:theory}, we ``zoom in'' and consider the local physics of black hole formation.  In \\S~{\\ref{sec:future}, we look into the future, with the goal to assess the power of forthcoming observations to probe the physics of the first Active Galactic Nuclei (AGN). We offer our conclusions in \\S~\\ref{sec:conclude}. ",
        "conclusions": "\\label{sec:conclude}  In this review, we have summarized theoretical ideas and observational constraints on how massive black holes form at the centers of galaxies, and how such black holes grow via accretion and mergers to give rise to the observed population of black holes in the local and moderate redshift universe.  This remains a poorly understood but important problem.  In addition to being of intrinsic interest for understanding the AGN phenomena, sources of gravity waves, etc., there is strong evidence that the formation and evolution of black holes is coupled to the formation and evolution of the host galaxy in which the black hole resides (e.g., the $\\mh-\\sigma$ relation), and thus to the cosmological formation of nonlinear dark matter structures (i.e., the dark halos surrounding these galaxies).  We anticipate that this will remain a growth area of research in the coming years, with continued rapid progress on both the observational and theoretical fronts."
    },
    "0403/astro-ph0403007_arXiv.txt": {
        "abstract": "Stochastic acceleration of $^3$He and $^4$He from a thermal background by parallel propagating turbulent plasma waves with a single power-law spectrum of the wavenumber is studied. In the model, both ions interact with several resonant modes. When one of these modes dominates, the acceleration rate is reduced considerably. At low energies, this happens for $^4$He, but not for $^3$He where contributions from the two stronger modes are comparable so that acceleration of $^3$He is very efficient. As a result, the acceleration of $^4$He is suppressed by a barrier below $\\sim 100$ keV nucleon$^{-1}$ and there is a prominent quasi-thermal component in the $^4$He spectra, while almost all the injected $^3$He ions are accelerated to high energies. This accounts for the large enrichment of $^3$He at high energies observed in impulsive solar energetic particle events. With reasonable plasma parameters this also provides a good fit to the spectra of both ions. Beyond $\\sim 1$ MeV nucleon$^{-1}$, the spectrum of $^3$He is softer than that of $^4$He, which is consistent with the observed decrease of the $^3$He to $^4$He ratio with energy. This study also indicates that the acceleration, Coulomb losses and diffusive escape of the particles from the acceleration site {\\it all} play important roles in shaping the ion spectra. This can explain the varied spectral shapes observed recently by the {\\it Advanced Composition Explorer}. ",
        "introduction": "The abundance and spectra of most ions observed in impulsive solar energetic particle events (SEPs) qualitatively agree with predictions of stochastic acceleration (SA) by plasma waves or turbulence (PWT) (e.g. M\\\"{o}bius et al. 1980; 1982). However, the observed enrichment of $^3$He, sometimes by more than four orders of magnitude over its coronal abundance (Hsieh \\& Simpson 1970; Serlemitsos \\& Balasubrahmanyan 1975; Mason et al. 2002), remains a theoretical challenge. The aim of this paper is to provide an explanation for this observation in the frame of SA models.  Solar $^3$He rich events are enriched also in other heavy elements (Hurford et al. 1975; Mason et al. 1986) and are closely related to impulsive X-ray flares, scatter-free electron events, and type III radio bursts (Reames, von Rosenvinge \\& Lin 1985; Reames, Meyer \\& von Rosenvinge 1994; Mazur, Mason \\& Klecker 1995; Mason et al. 2002), indicating that a common particle acceleration process may be responsible for all these features. Among the acceleration mechanisms applied to solar flares, the theory of SA, a second order Fermi acceleration process by PWT, has achieved several significant successes. The theory was first introduced to account for the acceleration of protons and other ions responsible for the observed nuclear gamma ray line emissions (Ramaty 1979; Hua, Ramaty \\& Lingenfelter 1989). Later it was shown that it is also applicable to the acceleration of electrons, which produce the impulsive hard X-ray emission (Miller \\& Ramaty 1987; Hamilton \\& Petrosian 1992; Park, Petrosian \\& Schwartz 1997; Petrosian \\& Donaghy 1999). A recent study of SA by cascading Alfv\\'{e}n turbulence indicated that it may also explain heavy ion acceleration as well (Miller 2003).  It is generally accepted that the enrichment of $^3$He is related to its peculiar charge to mass ratio so that certain plasma waves resonantly accelerate or heat it to high energies, producing the observed abundance ratio at a few MeV nucleon$^{-1}$ (M\\\"{o}bius et al. 1980; Zhang 1999). Previous work either discusses the feasibility of selective accelerating $^3$He from very low energies (Temerin \\& Roth 1992; Miller \\& Vi\\~{n}as 1993), or considers a two-stage mechanism; a selective preheating process followed by an acceleration applicable to all ions (Fisk 1978). The former cannot be compared with observation directly because the acceleration of other ions has not been addressed accordingly. For the latter, the observed abundance pattern at high energies is determined by the preheating processes (Paesold, Kartavykh \\& Benz 2003) and the predicted spectrum depends on several parameters with the corresponding physical processes unspecified (Sakurai 1974; Zhang 1995).  Given its simple dispersion relation, Alfv\\'{e}n turbulence is used in most of the SA models (Ramaty 1979; M\\\"{o}bius et al. 1982), which cannot explain the varied ion spectral forms and underestimates the production rate of higher energy particles in disagreement with the recent observations by the {\\it Advanced Composition Explorer} ({\\it ACE}; Mason et al. 2002). However, low energy particles mostly interact with cyclotron and/or whistler waves. To understand the acceleration of particles from a thermal background, it is necessary to use the exact dispersion relation (Steinacker et al. 1997). We have carried out such an investigation recently, which addresses the acceleration of electrons vs protons by waves propagating parallel to magnetic field lines (Petrosian \\& Liu 04, PL04 for short). Besides its achievements in explaining many features of solar flares, we found that the acceleration of low energy protons can be suppressed significantly relative to the electron acceleration due to the dominance of the wave-proton interaction by the resonant Alfv\\'{e}n mode.  In this letter, we show that a similar mechanism makes the acceleration of low energy $^3$He much more favorable relative to that of $^4$He. Because their loss rates are comparable, such a mechanism basically depletes $^3$He from the thermal background to $\\sim$ MeV energies while most of the $^4$He ions are ``trapped'' at low energies, resulting a prominent quasi-thermal $^4$He component. The model not only explains the enrichment of $^3$He at a few MeV nucleon$^{-1}$, but also accounts for the flattening of the $^3$He spectrum at lower energies (Mason, Dwyer \\& Mazur 2000; Mason et al. 2002). It also predicts that beyond a few MeV nucleon$^{-1}$ the spectrum of $^4$He is always harder than the $^3$He spectrum. In the next section, we present our SA model. Its application to SEPs is discussed in \\S\\ \\ref{seps}. In \\S\\ \\ref{sum}, we summarize our results and discuss their implications. ",
        "conclusions": "\\label{sum}  In the paper, we have shown that the observed $^3$He and $^4$He spectra in some impulsive SEPs can be produced via SA by PWT propagating parallel to magnetic field lines. In the model, the particles are injected from a thermal background. Because the interaction of low energy $^4$He is dominated by one of the resonant wave modes, an acceleration barrier forms, giving rise to a prominent quasi-thermal component in the steady state particle spectrum. This is not true for $^3$He ions, however, which can be efficiently accelerated by two waves in the PC branch, from very low energies to $\\sim$ MeV nucleon$^{-1}$, where they escape quickly from the acceleration site, producing a sharp spectral cutoff.  This explains the $^3$He enrichment and its varied spectrum seen in class 2 SEPs, which have relatively short rise-time and are more likely associated with the impulsive phase particle acceleration directly. Class 1 events have less $^3$He enrichment and similar spectra for most ions, indicating a different acceleration process (Mason et al. 2000). The feasibility of producing such events by shocks or turbulence with a different spectrum than that of the present model will be investigated in the future.  Because the $^4$He acceleration time is shorter than that of $^3$He at higher energies and their loss and escape times are comparable, the model predicts a decrease of the $^3$He to $^4$He ratio with energy beyond $\\sim 1$ MeV nucleon$^{-1}$, which is consistent with observations (Mason et al. 2002; Reames et al. 1997; M\\\"{o}bius et al.  1982; M\\\"{o}bius et al. 1980; Hsieh \\& Simpson 1970). Spectral information at higher energies is critical in testing this model prediction. {\\it We emphasize here that a special injection mechanism is not required to explain the observations.} However, a second-phase shock acceleration (Van Hollebeke, McDonald \\& Meyer 1990) can modify the ion spectra and may explain events where the high energy spectral indices of $^3$He and $^4$He are similar (Serlemitsos \\& Balasubrahmanyan 1975; Reames et al. 1997).  The model can also be applied to acceleration of heavy-ions. In this case, the acceleration is dominated by resonant interaction with waves in the HeC branch and ions with low charge to mass ratio are accelerated more efficiently, which may explain the increase of ion enrichment with the increase of their charge to mass ratio (Mason et al. 1986). Because $^3$He acceleration is dominated by interaction with the PC branch, which may be decoupled from the HeC branch, the model may also explain the lack of quantitative correlation of the enrichment of $^3$He and heavy-ions (Mason et al. 1986). We will investigate these aspects in future work.  The primary uncertainties of the model are associated with the turbulence generation mechanism and the subsequent cascade and dissipation processes. We start the investigation with a power law spectrum of the wavenumber for the turbulence, which can be characterized by the turbulence energy density and the spectral index. The acceleration of particles by obliquely propagating waves and turbulence with complex spectral shapes may give quantitatively different results. However, the idea of an efficient depletion of low energy $^3$He from the source region to high energies to produce the observed enhancement seems well founded."
    },
    "0403/gr-qc0403040_arXiv.txt": {
        "abstract": "We study the asymptotic behaviour of the Bianchi type VI$_0$ universes with a tilted $\\gamma$-law perfect fluid. The late-time attractors are found for the full 7-dimensional state space and for several interesting invariant subspaces. In particular, it is found that for the particular value of the equation of state parameter, $\\gamma=6/5$, there exists a bifurcation line which signals a transition of stability between a non-tilted equilibrium point to an extremely tilted equilibrium point. The initial singular regime is also discussed and we argue that the initial behaviour is chaotic for $\\gamma<2$. ",
        "introduction": "In a recent paper we analysed how tilted fluids affect the general late-time behaviour of cosmological models of Bianchi type \\cite{BHtilted}. As the asymptotic behaviour of all Bianchi models with a non-tilted $\\gamma$-law perfect fluid are determined \\cite{DynSys,HHTW,HHWVIII}, the sensitivity of the non-tilted equilibrium points was checked with regards to inclusion of tilted fluids. In this paper, we will consider one of the models in full generality; namely the Bianchi type VI$_0$ model containing a tilted $\\gamma$-law perfect fluid.  The last decades, there have been some investigations of Bianchi models with a tilted fluid \\cite{KingEllis,BN,BS}; in particular, type II \\cite{HBWII} and type V \\cite{Shikin,Collins,HWV}. Apart from these special Bianchi types, the general behaviour of tilted Bianchi universes seems to be poorly understood. The work \\cite{BHtilted} gave us some hints  where interesting behaviour may occur, but a more elaborate analysis is needed in order to understand the more general behaviour of universes with tilted fluids. More specifically, in \\cite{BHtilted} it was shown that for the Bianchi type VI$_0$ model there are no stable non-tilted equilibrium points for equation of state parameter obeying $\\gamma>6/5$. The value $\\gamma=6/5$ signals the onset of an instability with respect to tilt implying that at late times the peculiar velocities of the fluid is a major contributor to the cosmological shear. The Bianchi type VI$_0$  is not the most general model but is sufficiently general to account for many interesting phenomena. Thus we will assume that the universe contains a $\\gamma$-law perfect fluid ($0<\\gamma<2$) which in general can be tilted. The state space is thus 7-dimensional \\cite{DynSys} (compared to the 8-dimensional state space for the most general models).  The dynamical systems approach -- which is the method adopted here -- requires that we find all equilibrium points for the system of equations. These equilibrium points play a special role in the evolution of the system. Not only do they correspond to exact self-similar solutions to Einstein's equations, but  some of them may even serve as attractors for more general solutions. However, not many exact solutions with tilted fluids are known for the Bianchi models. Some type II solutions have been found \\cite{HewittII}, and Rosquist and Jantzen found some type VI$_0$ solutions \\cite{Rosquist,RJ}. Here, we will show that the solutions (for $6/5<\\gamma<3/2$) found by Rosquist and Jantzen are attractors in a four-dimensional invariant subspace of the full state space. However, in the full state space, there always exists an unstable mode. In the full state space, the future attractor for $6/5<\\gamma$ is an extremely tilted model which is connected via a line bifurcation at $\\gamma=6/5$ to a non-tilted equilibrium point. This $\\gamma=6/5$ bifurcation, which correspond to a one-parameter family of exact tilted solutions, seems to have been first discussed in a recent paper by Apostolopoulos \\cite{Apo}. However, not much discussion is devoted to this solution as the solution itself is implicitly given in terms of a cubic. Here, we have introduced a different parameterisation which makes it possible to explicitly write down the solution in terms of the expansion-normalised variables.  We also discuss the initial singular regime, which is a lot more complicated. In fact, there are indications that the initial singular regime possesses an oscillatory, very likely even a chaotic behaviour similar to that found in other Bianchi models \\cite{Spokoiny:1981fs,Belinsky:1970ew,jb2,jb3,cb,dem,DH2000,dh2,hbc}.  The paper is organised as follows. Next, in section \\ref{sect2}, we write down the equations of motion using the orthonormal frame formalism. In section \\ref{sect3} we discuss the equilibrium points which are important for the late-time asymptotic behaviour. Then, in section \\ref{sect4}, we discuss their stability and present some general results regarding the asymptotic behaviour of the tilted type VI$_0$ model. The initial singular regime is discussed in section \\ref{sect5} before we finally summarise our results  in  section \\ref{sect6}. ",
        "conclusions": "\\label{sect6} Here, for the first time, we have analysed the late-time behaviour of general tilted Bianchi type VI$_0$ universes. Our results are summarised in Table \\ref{stable}. We performed a local analysis and found all the future stable equilibrium points for various subclasses of tilted type VI$_0$ models as well as for the general tilted type VI$_0$. In particular, we confirmed the observation in \\cite{BHtilted} on the existence of new self-similar solutions for $\\gamma\\geq 6/5$. These solutions also proved to be important for the late-time behaviour; the Rosquist-Jantzen solutions are late-time attractors for a certain class of models with a two-component tilted fluid; and the line bifurcation at $\\gamma=6/5$ is, in fact, the late-time attractor for general tilted type VI$_0$ solutions. For  $\\gamma>6/5$, the late-time attractor is an extremely tilted model.\\footnote{It should be noted that the $\\gamma=6/5$ line bifurcation looks remarkably similar to the Wainwright $\\gamma=10/9$ line bifurcation of the exceptional model \\cite{Wainwright10-9}. This makes us wonder if there are unknown line bifurcations along the entire line given by $\\gamma=2(3+\\sqrt{-h})/(5+3\\sqrt{-h})$ for the more general type VI$_h$ models \\cite{BHtilted}.}  It is interesting to note that the general late-time attactors lie in the fixed-point-set of $\\phi_1$. This means that in general the solutions are asymptotically  $\\phi_1$-symmetric. More specifically, $\\phi_1$-symmetric implies $v_1=0$, and $v_2^2=v_3^2$.  At early times the analysis suggests chaotic behaviour as $\\tau\\rightarrow -\\infty$. This behaviour seems to be a generic property of tilted Bianchi models \\cite{UEWE}. So this behaviour was not very surprising, in particular considering the fact that the tilted type II model -- which is part of the boundary of type VI$_0$ -- was known  to be chaotic \\cite{HBWII}.  Let us finish off with some comments about future research and how this work may relate to other Bianchi types. Firstly, the work \\cite{BHtilted} indicates that there may be some similarities between the more general type VI$_h$ models and the type VI$_0$. Hence, some of the features of the late-time behaviour found in this work may also appear in the type VI$_h$ models. As goes for the class VII$_0$ model, which is given by $\\bar{N}^2>N^2$, there is one obvious difference. The general non-tilted VII$_0$ model is not asymptotically self-similar \\cite{WHU,NHW}. This happens because the type VII$_0$ state space is not compact; there is no upper bound on $\\bar{N}$. In the terminology of \\cite{BH}, the type VII$_0$ model is \\emph{extremely Weyl dominant} while the type VI$_0$ model is \\emph{Weyl-Ricci balanced} at late times. Hence, we would expect a fairly different behaviour for the tilted type VII$_0$ model at late times than that found here. Moreover, in this work the extremely tilted invariant sets have not been emphasised. For example, a two-fluid model where one fluid is extremely tilted is if particular interest. Further investigations in this case is required.  Nonetheless, there are  quite a few unanswered questions regarding Bianchi models with a perfect fluid. This work has answered some of them. Hopefully, future work will answer more of them."
    },
    "0403/astro-ph0403694_arXiv.txt": {
        "abstract": "The modified newtonian dynamics (MOND)  paradigm of Milgrom can boast of a number of successful predictions regarding galactic dynamics; these are made without the assumption that dark matter plays a significant role.  MOND requires gravitation to depart from Newtonian theory in the extragalactic regime where dynamical accelerations are small.  So far relativistic gravitation theories proposed to underpin MOND have either clashed with the post-Newtonian tests of general relativity, or failed to provide significant gravitational lensing, or violated hallowed principles by exhibiting superluminal scalar waves or an \\textit{a priori} vector field.  We develop a relativistic MOND inspired theory which resolves these problems.  In it gravitation is mediated by metric, a scalar field and a 4-vector field, all three dynamical.  For a simple choice of its free function,  the theory has a Newtonian limit for nonrelativistic dynamics with significant acceleration, but a MOND limit when accelerations are small. We calculate the $\\beta$ and $\\gamma$ PPN coefficients showing them to agree with solar system measurements.   The gravitational light deflection by nonrelativistic systems is governed by the same potential responsible for dynamics of particles.  To the extent that MOND successfully describes dynamics of a system, the new theory's predictions for lensing by that system's visible matter will agree as well with observations as general relativity's predictions made with a dynamically successful dark halo model.  Cosmological models based on the theory are quite similar to those based on general relativity; they predict slow evolution of the scalar field.  For a range of initial conditions, this last result makes it easy to rule out superluminal propagation of metric, scalar and vector waves. ",
        "introduction": "In the extragalactic regime, where Newtonian gravitational theory would have been  expected to be an excellent description, accelerations of stars and gas, as estimated from  Doppler velocities and geometric considerations, are as a rule much larger than those  due to the Newtonian field generated by the visible matter in  the system~\\cite{Oort,Zwicky}.  This is the  ``missing mass'' problem~\\cite{DMU} or ``acceleration discrepancy''~\\cite{Can2}.  It is fashionable to infer from it the existence of much dark matter in systems ranging from dwarf spheroidal galaxies with masses $\\sim 10^6 M_\\odot$ to great clusters of galaxies in the $10^{13} M_\\odot$ regime~\\cite{DMU,Sweden}.  And again, galaxies and clusters of galaxies are found to gravitationally lense  background sources. When interpreted within general relativity (GR), this lensing is anomalously large unless one assumes the presence of dark matter in quantities and with distribution similar to those required to explain the accelerations of stars and gas.  Thus extragalactic lensing has naturally been regarded as confirming the presence of the dark matter suggested by the dynamics.  But the putative dark matter has never been  identified despite much experimental and observational effort~\\cite{TurnerEllis}.  This raises the possibility that the acceleration discrepancy as well as the gravitational lensing anomaly may reflect departures from Newtonian gravity and GR on galactic and larger scales.  Now alternatives to GR are traditionally required to  possess a Newtonian limit for small velocities and potentials; thus the acceleration discrepancy also raises the possibility that the correct relativistic gravitational theory may be of a kind not generally considered hitherto.  In the last two decades Milgrom's modified Newtonian dynamics (MOND) paradigm~\\cite{M1,M2,M3} has gained recognition as a successful scheme for unifying much of extragalactic dynamics phenomenology without invoking ``dark matter''.  In contrast with earlier suggested modifications of Newton's law of universal gravitation~\\cite{Berendzen,Jeans,Morpho,Finzi}, MOND is characterized by an acceleration scale $\\mathfrak{a}_0$, not a distance scale, and its departure from Newtonian predictions is acceleration dependent: \\begin{equation} \\tilde\\mu(|\\mathbf{a}|/\\mathfrak{a}_0)\\mathbf{a} = -\\bm{\\nabla}\\Phi_{\\rm N}. \\label{MOND} \\end{equation} Here $\\Phi_{\\rm N}$ is the usual Newtonian potential of the visible matter, while $\\tilde\\mu(x)\\approx x$ for $x\\ll 1$ and $\\tilde\\mu(x)\\rightarrow 1$ for $x\\gg 1$.  Milgrom estimated  $\\mathfrak{a}_0\\approx 1\\times 10^{-8}$ cm s$^{-2}$ from the empirical data. In the laboratory and the solar system where accelerations are strong compared to $\\mathfrak{a}_0$, formula (\\ref{MOND}) goes over to the Newtonian law $\\mathbf{a} =-\\bm{\\nabla}\\Phi_{\\rm N}  $.  Milgrom constructed formula (\\ref{MOND}) to agree with the fact that  rotation curves of disk galaxies become flat outside their central parts.  That far out a galaxy of mass $M$ exhibits an approximately spherical Newtonian potential.  The scales are such that $|\\bm{\\nabla}\\Phi_{\\rm N}|\\approx GM r^{-2} \\ll \\mathfrak{a}_0$ in this region, and so Eq.~(\\ref{MOND}) with $\\tilde\\mu(x)\\approx x$ gives $|\\mathbf{a}|\\approx (GM\\mathfrak{a}_0)^{1/2} r^{-1}$ which has the $r$ dependence appropriate for the centripetal acceleration $v_c^2/r$ of a radius independent rotational velocity $v_c$---an asymptotically  flat rotation curve.   In fact one obtains the relation $M=(G\\mathfrak{a}_0)^{-1}v_c{}^4$ which leads to the \\textit{prediction} that for any class of galaxies with a constant mass to luminosity ratio $\\Upsilon$ in a specified spectral band, the luminosity in that band should scale as $v_c{}^4$.    And indeed, there exists an empirical law of just this form: the Tully--Fisher law~\\cite{TullyFisher} (TFL) relating near infrared (H--band) luminosity $L_H$ of a spiral disk galaxy to its rotation velocity, $L_H\\propto v_c{}^4$, with the proportionality factor being constant within each galactic morphology  class.  This version of the TFL was established only after MOND was enunciated~\\cite{Aaronson}.  It is in harmony with the MOND prediction in two ways.   First, the infrared light of a galaxy comes mostly from cool dwarf  stars which make up most of its mass (hence giving a tight correlation between $M$ of the predicted relation and $L_H$ of the empirical law). Second, the proportionality coefficient varies from class to class as would be expected from the observed correlation between $\\Upsilon$ of a galaxy and its morphology.  In the alternative dark matter paradigm (which casts no doubt on standard gravity theory), flat rotation curves are explained by assuming  that every disk galaxy is nested inside a roundish spherical halo of dark matter~\\cite{OstrikerPeeblesYahil} whose mass density drops approximately like $r^{-2}$.  The halo is supposed to dominate the gravitational field in the outer parts of the galaxy.  This makes the Newtonian potential approximately logarithmic with radius in those  regions, thereby leading to an asymptotically flat rotation curve.   In practice the dark halo resolution works  only after fine tuning.   It is an observational fact that for bright spiral galaxies the rotation curve in the optically bright region is well explained in Newtonian gravity  by the observed matter~\\cite{Kalnajs}.  But, as mentioned, in the outer regions the visible matter's contribution must be dwarfed by the halo's.  So fine tuning is needed between the dark halo parameters (velocity  dispersion and core radius) and the visible disk ones~\\cite{BahcallCassertano,Sancisi}.  This fine  tuning problem is exacerbated by the TFL $L_H\\propto v_c{}^4$.   Because the infrared luminosity comes from the visible matter in the  galaxy, but  the rotation velocity is mostly set by the halo, the TFL  also requires fine tuning between halo and disk parameters. The standard dark matter explanation of the $r^{-2}$ profile of an halo is that it arises naturally from primordial cosmological perturbations~\\cite{QuinnSalmonZurek}.  The visible galaxy is regarded as forming by dissipational collapse of gas into the potential well of  the halo.  The fine tuning mentioned is then viewed as resulting  from the adjustment of the halo to the gravitation of the incipient disk~\\cite{Sancisi,Barnes}.  But the TFL is observationally a  very sharp correlation; in fact, it is the basis for one of the most reliable methods for gauging distances to spiral galaxies.  Such sharpness is hardly to be expected from statistical  processes of the kind envisaged in galaxy formation, a point emphasized by Sanders~\\cite{SReview}. So in the dark matter picture the TFL is something of a mystery.  There are other MOND successes.  Milgrom predicted early that in galaxies with surface mass density well below $\\mathfrak{a}_0 G^{-1}$, the acceleration discrepancy should be especially large~\\cite{M2}.   In dwarf spirals this property was established empirically years later~\\cite{JobinCarignan}, and it is now known to be exhibited by a large number of low surface brightness galaxies~\\cite{McGaugh}.  Another example: MOND successfully predicts the detailed shape of a rotation curve from the observed matter (stars and gas) distribution on the basis of a single free parameter, $\\Upsilon$, down to correlating  features in the velocity field with those seen in the light distribution~\\cite{Kent,McKent,BBS,Gentile}.    This is especially true in the case of low surface mass density disk galaxies for which MOND's predictions are independent of the specific choice of $\\tilde\\mu(x)$~\\cite{MLSD}, and these MOND theoretical rotation curves fit the observed curves of a number of low surface brightness dwarf galaxies~\\cite{BBS,MCGdB,Begum} very well.   By contrast, the dark halo paradigm requires one or two free parameters apart from $\\Upsilon$  to approximate the success of the MOND predictions~\\cite{SandersMcGaugh}.   In fact, even when the empirical data is analyzed within the dark halo paradigm, it displays the preferred acceleration scale $\\mathfrak{a}_0$ of MOND~\\cite{McGaugh2}.  Occasionally doubt has been cast on MOND's ability to describe clusters of galaxies properly~\\cite{Aguirre}.  Many of these exhibit accelerations not small on scale $\\mathfrak{a}_0$, yet conventional analysis suggests they contain much dark matter in opposition to what MOND would suggest.  Sanders has recently reanalyzed the problem~\\cite{Sanders} with the conclusion that these clusters may contain much as yet undiscovered baryonic matter in the core which should be classed as ``visible'' in connection with MOND. Other MOND successes, outside the province of disk galaxies, have been reviewed elsewhere~\\cite{SReview,SandersMcGaugh,more}.  So the simple MOND formula~(\\ref{MOND}) is very successful.  But it is not a theory.  Literally taken the MOND recipe for acceleration violates the conservation of momentum (and of energy and of angular momentum)~\\cite{M1}.   And MOND entails a paradox: why does the center of mass of a star orbit in its galaxy with anomalously large acceleration given by Eq.~(\\ref{MOND}) with $\\tilde\\mu\\ll 1$, while each parcel of gas composing it is subject to such high acceleration that is should, by the same formula, be accelerated Newtonially ?~\\cite{M1}.   In short, the MOND formula is not a consistent theoretical scheme.  Neither is MOND, as initially stated, complete.  For example, it does not specify how to calculate gravitational lensing by galaxies and clusters of galaxies.  As is well known, in standard gravity theory light deflection is well described only by relativistic theory (GR).  And whereas Newtonian cosmological models work well for part of the cosmological evolution, MOND cosmological models built in analogy with their Newtonian counterparts, though sometimes agreeing with phenomenology~\\cite{Aguirre},  can  yield peculiar predictions~\\cite{Felten} (but see Ref.~\\onlinecite{SandersC}).  In short, a complete, consistent theoretical underpinning of the MOND paradigm which accords with observed facts, and is also relativistic, has been lacking.  This lack is being resolved in measured steps.   A first step was the lagrangian reformulation of MOND~\\cite{BekMilg}  called AQUAL (see Sec.~\\ref{sec:nonrelat}).   AQUAL cures the nonconservation problems and resolves the paradox of the galactic motion of an object whose parts accelerate strongly relative to one another; it does so in accordance with a conjecture of Milgrom~\\cite{M1}.  And for systems with high symmetry AQUAL reduces exactly to the MOND formula (\\ref{MOND}).  A relativistic generalization of AQUAL is easy to construct with help of a scalar field which together with the metric describes gravity~\\cite{BekMilg} (see Sec.~\\ref{sec:rAQUAL} below).  It reduces to MOND approximately in the weak acceleration regime, to Newtonian gravity for strong accelerations, and can be made consistent with the post-Newtonian solar system tests for GR.  But relativistic AQUAL is acausal: waves of the scalar field can propagate superluminally in the MOND regime (see the appendix of Ref.~\\onlinecite{BekMilg} or Appendix~\\ref{sec:A} here).  The problem can be traced to the aquadratic kinetic part of the lagrangian of the theory which mimics that in the original AQUAL.  A theory involving a second scalar field, PCG, was thus developed to bypass the problem~\\cite{Can2,phase,SPCG} (see Sec.~\\ref{PCG} below).   PCG may be better behaved causally  than relativistic AQUAL~\\cite{Rosen}, but it brings woes of its own.  It is marginally in conflict with the observed perihelion precession of Mercury~\\cite{Can2}, and in common with relativistic AQUAL, PCG predicts extragalactic gravitational lensing which is too weak if there is indeed no dark matter.   This last problem is  traceable to a feature common to PCG and relativistic AQUAL: the physical metric is conformal to the metric appearing in the Einstein-Hilbert action~\\cite{Kyoto}.  One way to sidestep this problem without discarding the MOND features is to exploit the direction defined by the gradient of the first scalar field to relate the physical metric to the Einstein metric by a disformal transformation  (see Ref.~\\cite{Kyoto} or Sec.~\\ref{disf} below).  But  it turns out that with this relation the requirement of causal propagation acts to depress gravitational lensing~\\cite{BekSan}, rather than enhancing it as is observationally required.  The persistence of the lensing problem in modified gravitational theories has engendered a folk theorem to the effect that it is impossible for a relativistic theory to simultaneously incorporate the MOND dynamics, observed gravitational lensing and correct post-Newtonian behavior without calling on dark matter~\\cite{Nester,Edery,Soussa1,Soussa2}.  Needless to say, this theorem cannot be proved~\\cite{Soussa3}. Indeed, by the simple device of relating the physical and Einstein metrics via a disformal transformation based on a \\textit{constant} time directed 4-vector, Sanders~\\cite{Sandersstratified} has constructed an AQUAL like ``stratified''  relativistic theory which gives the correct lensing while ostensibly retaining the MOND phenomenology and consistency with the post-Newtonian tests. Admittedly Sanders' stratified theory is a preferred frame theory, and as such outside the traditional framework for gravitational theories.  But it does point out a trail to further progress.  The present paper introduces T$e$V$e\\,$S, a new relativistic gravitational theory devoid of \\textit{a priori} fields, whose nonrelativistic weak acceleration limit accords with MOND while its nonrelativistic strong acceleration regime is Newtonian. T$e$V$e\\,$S is based on a metric, and dynamic scalar and 4-vector fields (one each); it naturally involves one free function, a length scale, and two positive dimensionless parameters, $k$ and $K$.   T$e$V$e\\,$S passes the usual solar system tests of GR, predicts gravitational lensing in agreement with the observations (without requiring dark matter), does not exhibit superluminal propagation, and provides a specific formalism for constructing cosmological models.  In Sec.~\\ref{sec:Found} we summarize the foundations on which a workable relativistic formulation of MOND must stand.  We follow this with a brief critical review of relativistic AQUAL, PCG and disformal metric theories, some of whose elements we borrow. Sec.~\\ref{sec:actions} builds the action for  T$e$V$e\\,$S while Sec.~\\ref{sec:eqns} derives the equations for the metric, scalar and vector fields.  In Sec.~\\ref{sec:GRlimit} we demonstrate that T$e$V$e\\,$S has a GR limit for a range of small $k$ and $K$.  This is shown explicitly for cosmology (Sec.~\\ref{sec:cosmology}) and for quasistatic situations like galaxies (Sec.~\\ref{sec:quasistationary}).   All the above applies for any choice of the free function; in Sec.~\\ref{sec:choice} we make a simple choice for it which facilitates further elaboration.   For spherically symmetric systems the nonrelativistic MOND limit is derived in Sec.~\\ref{sec:Mondish}, while the Newtonian limit is recovered for modestly small $k$ in Sec.~\\ref{sec:Newtonian}.  The above conclusions are extended to nonspherical systems in Sec.~\\ref{sec:nonspherical}.  Sec.~\\ref{sec:postN} shows that the theory passes the usual post-Newtonian solar system tests. Sec.~\\ref{sec:lensing} demonstrates that for given dynamics, T$e$V$e\\,$S gives the same gravitational lensing as does a dynamically successful dark halo model within GR. In Sec.~\\ref{sec:models} we discuss T$e$V$e\\,$S cosmological models with flat spaces showing that they are very similar to the corresponding GR models (apart from the question of cosmological dark matter which is left open), and demonstrating that the scalar field evolves little, and so can be taken to be small and positive.   As discussed next in Sec.~\\ref{sec:causality}, this last conclusion serves to rule out superluminal propagation in T$e$V$e\\,$S. ",
        "conclusions": ""
    },
    "0403/astro-ph0403377_arXiv.txt": {
        "abstract": " ",
        "introduction": "Chemodynamical simulations are characterized by a multi-phase treatment of the interstellar medium (ISM). In most of these calculations the ISM is split into two dynamically different components, the clumpy star-forming molecular clouds and the diffuse warm or hot gas (Theis et al.\\ 1992, Samland et al.\\ 1997). More refined decompositions of the ISM are only done for one-zone models (e.g.\\ Ikeuchi et al.\\ 1984), but not for gravitationally coupled, dynamically evolving systems. Especially, the dust component in galaxies is usually just seen as a tracer population which has no dynamical influence on galactic dynamics. On the other hand, it is well known that even a small admixture of a dynamically cold component can substantially destabilize galactic structures as Jog \\& Solomon (1984) demonstrated for a small amount of cold gas embedded in stellar disks. The aim of our paper is to study, whether a dynamically cold dust phase can also destabilize galactic disks.  Good places for looking on the influence of dust are galactic central regions. During the last years high resolution observations revealed that the circumnuclear regions contain stellar-gaseous mini-disks with sizes of a few hundred pc (e.g.\\ Carollo et al.\\ 1997). These disks have an astonishing rich structure described as mini-bars, spiral-like dust lanes, star-forming rings or spiral arms. Many nuclear regions are also well described as patchy or multi-armed. Investigations by Regan \\& Mulchaey (1999) showed that these spirals are the most common morphological structures in the central regions of galaxies.  Since the nature of the mini-spirals is of great importance for our understanding of a variety of astrophysical processes like the mass accretion into galactic nuclei, several explanations for have been invoked. E.g.\\ Athanassoula (1992) studied the gas flows in and around bars by (stellar-dynamical) orbital analysis. She found that the existence of different periodic orbit families (or equivalently the existence of inner Lindblad resonances (ILR)) is a key criterion for the existence and the shape of dust lanes.  Different to the stellar-dynamical interpretation Engl\\-maier \\& Shlosman (2000) suggested that mini-spirals in central regions of galactic disks are related to the formation of grand-design spiral patterns in galaxies. They argued that gas density waves -- different to stellar density waves -- are not completely damped or absorbed at the ILR and, thus, they may generate spiral structures at all radii including the nuclear regions. Such a model might explain the continuity of some spiral features at small and large radii as well as the low arm-interarm contrast observed in galactic centers. A similar idea of induced structure formation is to invoke secondary bars or small nuclear bars located inside the ILR (Wada \\& Koda 2001).  A common property of all the mentioned mechanisms is that they result in structures dominated by two or a few arms. On the other hand, Elmegreen \\etal (2002) stressed that nuclear dust spirals differ from main-disk spirals in several respects: the nuclear spirals are very irregular with both trailing and leading components that often cross.  All discussed mechanisms have some weak points and can not fully account for all the observational data. We consider here a new approach based on the observation that circumnuclear disks of galaxies are dusty. By means of a stability analysis Noh \\etal (1991) showed that a dust component can strongly destabilize proto-planetary disks. The admixture of only 2\\% of dust enhances the growth rates of the dominant gaseous phase significantly. A conservative estimate of the dust-to-gas ratio $r$ in the solar neighborhood gives an average value of 0.6\\% ranging from 0.2\\% up to 4\\% in $H_2$ regions (Spitzer 1978). These are lower limits because large, massive dust grains do not have a detectable extinction. Direct satellite-based measurements give $r$-values up to 2\\% (Frisch \\etal 1999).  Little is yet known about the true dust-to-gas ratio in external galaxies, but some galaxies show large values. E.g.\\ M51 has a dust-to-gas ratio of 2\\%. Taking the observed relation between the dust-to-gas ratio and the metallicity (Issa et al.\\ 1990) into account, it is reasonable to assume larger dust mass fractions in the metal-enriched central regions of galaxies than on average.  In this paper we study the dynamical influence of a dust component on circumnuclear gaseous disks. We restrict our analysis to systems without a nuclear or a large scale bar and we do not consider any formation or destruction processes of the dust. This investigation is done by 2D multi-component hydrodynamical simulations. ",
        "conclusions": ""
    },
    "0403/gr-qc0403106_arXiv.txt": {
        "abstract": "In loop quantum gravity, modifications to the geometrical density cause a self-interacting scalar field to accelerate away from a minimum of its potential. In principle, this mechanism can generate the conditions that subsequently lead to slow-roll inflation. The consequences for this mechanism of various quantization ambiguities arising within loop quantum cosmology are considered. For the case of a quadratic potential, it is found that some quantization procedures are more likely to generate a phase of slow--roll inflation. In general, however, loop quantum cosmology is robust to ambiguities in the quantization and extends the range of initial conditions for inflation. ",
        "introduction": "\\setcounter{equation}{0}  Loop quantum gravity (LQG) or quantum geometry  is at present the main background independent and non--perturbative candidate for a quantum theory of gravity (see for example \\cite{rovelli98,thiemann02}). Key successes of this approach have been the prediction of discrete spectra for geometrical operators \\cite{geometrical_op}, the existence of well defined operators for the matter Hamiltonians which provides a cure for the ultraviolet divergences \\cite{thiemann_matter}, and the derivation of the Bekenstein--Hawking entropy formula \\cite{bek_hawking}.  Given that LQG effects are likely to have important consequences in high energy and high curvature regimes, early universe cosmology provides a natural environment to test these new features. In recent years, considerable interest has focused on applying LQG to minisuperspace cosmological models (see \\cite{martin_hugo,martin_kevin_2,ICGC} for reviews) and this has resolved various difficulties encountered in  conventional Wheeler--de Witt quantization for isotropic models \\cite{martin1,martin2} and anisotropic models \\cite{martin_date}.  {}From the loop quantum cosmology (LQC) perspective, the evolution of the universe is comprised of the three distinct phases. Initially, there is a truly discrete quantum phase which is described by a difference equation \\cite{martin1,martin2}. A key consequence of this discretization is the removal of the initial singularity \\cite{martin1}. As its volume increases, the universe enters an intermediate semi--classical phase in which the evolution equations take a continuous form but include modifications due to non--perturbative quantization effects \\cite{martin3}. Finally, there is the classical phase in which the usual continuous ODE/PDE cosmological equations are recovered and quantum effects vanish.  The above intermediate phase (which is demarcated by two scales, as discussed in detail in \\S 2) is the most important as far as phenomenological studies of LQC are concerned and will be subject of our study here. Key results arising from this phase include the setting up of suitable initial conditions for inflation \\cite{martin3,martin_kevin_1,tsm03}, possible signatures on the CMB spectrum such as the loss of power at largest scales and the running of the scalar spectral index \\cite{tsm03} and the avoidance of a big crunch in closed models \\cite{st03}.  Despite these successes, however, it is known that the quantization procedure itself is not unique. In general, quantization ambiguities arise when composite operators are constructed from the corresponding classical expressions. Some of these ambiguities are more fundamental from a theoretical point of view than others. Nevertheless, it is possible that the range of different choices that lead to a viable cosmological model might be significantly narrowed by employing not only theoretical considerations but also by confronting phenomenological predictions with observations. Two immediate questions arise in this connection: (i) do some/all ambiguities of this kind lead to observable predictions/consequences and (ii) are some predictions of LQC (such as establishing the correct initial conditions for inflation) robust with respect to these ambiguities? For example, if a particular ambiguity were to result in observational effects that were subsequently detected, this could provide an observational basis for invoking that particular choice. On the other hand, the inability of ambiguities to produce observationally distinguishable consequences would make it difficult to discriminate between them without theoretical input. Whatever the outcome, a study of these questions would demonstrate the range of phenomenological possibilities allowed by LQC and would thus enable a clearer comparison to be made between the theoretical predictions and observations.  A study of the effects of quantization ambiguities on the LQG's capacity to remove singularities was recently made \\cite{bojowaldCQG2002}, where it was found that the removal of the big-bang singularity is robust with respect to a number of different ambiguities.  Our aim here is to study the effects of quantization ambiguities on the evolution of the very early universe. In particular, we shall study the way both kinematical and dynamical ambiguities can affect the ability of LQC to set up the appropriate conditions for inflation.  We focus primarily on the possibility that in LQC inflation may arise naturally even when the inflaton is initially located near to a minimum of its potential. As we shall see, there are two mechanisms that affect the setting of the initial conditions for conventional slow--roll inflation. Firstly, there is an anti-friction effect induced by the semi-classical modification of the scalar field equation of motion which can, in an expanding universe, accelerate the field away from its initial value. Secondly, there are effects due to the quantum mechanical uncertainty relations. Since these mechanisms are in principle unrelated, each is treated separately in order to maintain a clear distinction between the two. Finally, in phenomenological studies of this nature, there are constraints, in addition to the observational bounds, that arise from the requirement that the region of parameter space under consideration should be consistent with the approximations that apply in the semi--classical and classical epochs.  We shall see that these also provide strong constraints on the set of ambiguities that lead to appropriate initial conditions for the onset of inflation.  The plan of the paper is as follows. In \\S 2, the ambiguities that arise in LQC are discussed, including a comparison of such ambiguities from the perspective of the full theory. Section 3 considers issues relevant to the setting of initial conditions in inflation. Section 4 presents an approximate analytical scheme for following the evolution of the field during the semi--classical phase and \\S 5 presents the numerical results. We conclude with a discussion and consider future directions in \\S 6. ",
        "conclusions": "In this paper we have considered some of the cosmological consequences of ambiguities that arise in loop quantum gravity. We have focused primarily on the importance of these ambiguities in realizing the conditions that lead to a phase of slow--roll inflation when the inflaton is initially located in a minimum of its potential. We have invoked a semi--classical treatment, where the dynamics of the universe is governed by non--linear ODEs. The key requirement for successful inflation within this framework is that sufficient inflation is possible in the region of parameter space where this approximation remains valid. In particular, this implies that the initial kinetic energy of the inflaton must be sufficiently small.  Our main conclusion is that the initial conditions for slow--roll inflation can be realized in a wide region of parameter space in loop quantum cosmology. In this sense, therefore, LQC is robust to ambiguities in the quantization and enhances the allowed range of initial conditions for inflation. In particular, kinetic energies many orders of magnitude below the Planck scale can lead to inflation. Moreover, parameters can be chosen such that the turning point of the inflaton is near to $ 3 l_{\\rm Pl}^{-1}$, corresponding to the largest scales accessible to observations. Equation (\\ref{Bphimax}) shows that $\\phi_{\\max}$ depends only logarithmically on $j$ and $\\dot{\\phi}_i$ and thus is sensitive only to the order of magnitude of those values. (Indeed, by using the Lambert function $W(x)$ defined such that $W(x)e^{W(x)}=x$ and behaving like a logarithm for $x>1$, we have $l_{\\rm Pl}\\phi_{\\rm max}=W(140\\sqrt{12\\pi} j^6| \\dot{\\phi}_i l_{\\rm Pl}^2|e^{\\sqrt{12\\pi}l_{\\rm Pl}\\phi_i})$.) This verifies what can already be observed from Fig.\\ 2 of Ref. \\cite{tsm03}.  For the choice $l=3/4$, conditions for successful inflation can be achieved in the {\\sc Fried} quantization scheme if $j$ is sufficiently high even when the field is located at the potential minimum. The field tends to move further up the potential in the {\\sc Fried} quantization than in {\\sc Ham} quantization. In this sense, {\\sc Fried} quantization might be favored from a phenomenological point of view at the semi--classical level, as it results in a larger region of parameter space for successful inflation (this has already been indicated in \\cite{golam}). This is interesting given that the {\\sc Ham} quantization scheme is directly based on a Hamiltonian whereas {\\sc Fried} involves additional multiplication with eigenvalues of the geometrical density operator and thus appears less natural conceptually.  In {\\sc Ham} quantization, the field needs to be displaced from its minimum, for example by quantum uncertainty effects, if it is to move sufficiently far up the potential without violating the bounds imposed by the semi--classical approximation. However, care should be taken in interpreting the kinetic bound (\\ref{kineticlimit}) -- the solution (\\ref{kineticsol}) overestimates the value of the field's kinetic energy at the transition since the form of the eigenvalue function, $D_l$, given in Eq. (\\ref{approxD}) represents its asymptotic form in the limit $a \\ll a_*$. Moreover, the estimate for $a_S$, the scale factor at the transition epoch, given in Eq. (\\ref{approxaS}), is not precise, since there is no exact definition for this parameter. If the kinetic bound (\\ref{kineticlimit}) were to be relaxed slightly, then it would become marginally consistent with the horizon problem constraint (\\ref{Bhorizon}). This indicates that {\\em if} successful inflation could be realized within this semi--classical framework, the conditions would be such that just enough e--folds of accelerated expansion would arise for the horizon problem to be solved. As a result, astrophysically observable scales today would correspond to the turning point in the field's dynamics.  It should be emphasized that failure to satisfy the kinetic bounds does not necessarily rule out certain parameter choices or quantization schemes, but only limits the allowed range where the approximation to ODEs can be employed. When the Hubble length becomes smaller than the fundamental discreteness scale, $a_i$, the effective ODEs describing the cosmic dynamics become invalid and must be replaced by the full quantum equations, which are a difference equation for the wave function in the case of loop quantum cosmology \\cite{bojowald02a}. It is possible that in this framework both the {\\sc Ham} and {\\sc Fried} schemes would become equally viable at a phenomenological level.   Considering values $l < 3/4$ does not alter the qualitative behavior of the dynamics. In general, for a given initial kinetic energy, this leads to an increase in the lowest value of $j$ consistent with successful inflation in {\\sc Fried} quantization, since the semi--classical phase must last longer. Conversely, when $l >3/4$, successful inflation is possible for lower initial kinetic energies and lower values of $j$.  For {\\sc Ham} quantization, reducing $l$ makes it harder to satisfy the horizon and kinetic bounds simultaneously. The kinetic energy of the field increases less rapidly for $l \\rightarrow 0$ and it might therefore be expected that it would be easier to satisfy the kinetic bound. However, in this case the field lacks the kinetic energy needed to reach a sufficiently high value after the semi-classical era. Thus, the phenomenological indications for $l$ are in agreement with the expectations from the full theory which lead to $l\\geq 1/2$.  In fact, the results allow us to draw further lessons for the full theory. As discussed in \\S 2.2, {\\sc Ham} is analogous to the full quantization procedure whose dynamics is governed by a constraint, while {\\sc Fried} does not have a full analog. From the phenomenological point of view we can separate the range of parameters for different quantizations into three distinct domains: (i) values which do not lead to sufficient slow-roll inflation; (iia) values which lead to sufficient inflation in such a way that the scalar field reaches a turning point around $3l_{\\rm Pl}^{-1}$; and (iib) values leading to sufficient inflation with a turning point much higher than $3l_{\\rm Pl}^{-1}$. As we have seen, {\\sc Fried} is phenomenologically more robust to ambiguities and most cases fall into class (iib). On the other hand, {\\sc Ham} is less robust but has most realisations in class (iia) which are more likely to lead to observable effects by putting the maximal inflaton value just at the borderline for sufficient inflation. This may indicate that observations of full loop quantum gravity are indeed within reach.  As for $j$, the lower bounds are larger than unity, as expected, but not unreasonably large. Thus, the scenario appears realistic and does not require fine tuning, which is also a consequence of the weak logarithmic dependence of $\\phi_{\\rm max}$ on $j$.  A key open question in LQC concerns the evolution of scalar (density) and tensor (gravitational wave) perturbations generated during the semi--classical epoch. Although standard perturbation theory is unstable in this regime \\cite{tsm03}, it is possible that these perturbations may imprint characteristic signatures of this phase on the CMB power spectrum. Indeed, there have recently been a number of suggestions for explaining the apparent suppression of the CMB power spectrum on large angular scales through quantum gravity effects, including LQC \\cite{tsm03}, non--commutative geometry \\cite{roy1} and higher--order curvature terms in string--inspired models \\cite{shinji}. It is important, therefore, to investigate whether observations will be able to discriminate between these different possibilities. Finally, it would also be interesting to extend the analysis of this work to the regime of the quantum difference equations of the full theory and to investigate whether any observable effects for the different quantization schemes can be uncovered.  \\vspace{.3in} \\centerline{\\bf Acknowledgments} \\vspace{.3in} We are grateful to Roy Maartens for discussions. DM is supported by a PPARC studentship. PS thanks Queen Mary, University of London for warm hospitality during the initial stages of this work. He is supported by a CSIR grant.  \\begin{appendix}"
    },
    "0403/astro-ph0403188_arXiv.txt": {
        "abstract": "The collapse of massive stars not only produces observable outbursts across the entire electromagnetic spectrum but, for Galactic (or near-Galactic) supernovae, detectable signals for ground-based neutrino and gravitational wave detectors.  Gravitational waves and neutrinos provide the only means to study the actual engine behind the optical outbursts: the collapsed stellar core.  While the neutrinos are most sensitive to details of the equation of state, gravitational waves provide a means to study the mass asymmetries in this central core.  We present gravitational wave signals from a series of 3-dimensional core-collapse simulations with asymmetries derived from initial perturbations caused by pre-collapse convection, core rotation, and low-mode convection in the explosion engine itself.  A Galactic supernovae will allow us to differentiate these different sources of asymmetry.  Combining this signal with other observations of the supernova, from neutrinos to gamma-rays to the compact remnant, dramatically increases the predictive power of the gravitational wave signal.  We conclude with a discussion of the gravitational wave signal arising from collapsars, the leading engine for long-duration gamma-ray bursts. ",
        "introduction": "The mechanism behind core-collapse supernovae (SNe) remains one of the longest outstanding problems in astrophysics.  Whether or not simulations of stellar collapse produce explosions depends sensitively on the numerical implementations of the neutrino transport and the equation of state physics (Herant et al. 1994; Burrows, Hayes, Fryxell 1995; Janka \\& M\\\"uller 1996; Mezzacappa et al. 1998; Fryer \\& Warren 2002; Buras et al. 2003).  10\\% differences in this physics can make the difference between the success or failure of an explosion.  But it may be that the sensitivity of the current simulations occurs only because theorists are missing essential aspects of physics in the core.  For example, rotation in the core (Fryer \\& Heger 1998; Akiyama et al. 2003; Kotake, Yamada, \\& Sato 2003a; Fryer \\& Warren 2004), asymmetric collapse (Burrows \\& Hayes 1996; Fryer 2004a), one-sided convection due to oscillations (Blondin, Mezzacappa \\& DeMarino 2003), or the merger of modes (Scheck et al. 2003) may help or hinder the explosion.  The reason we know so little about the core and the resultant supernova explosion is because we have very few ways of directly observing the engine behind these explosions.  The collapsed core is enshrouded by the outer layers of the star and it is not until the supernova shock reaches the edge of the star and becomes optically thin to photons that we can finally see the supernova.  Then begins the laborious task of working back from the optical display to study the workings of the inner core.  Although these observations (see Akiyama et al. 2003 and references therein), and observations of the neutron star remnant (see Lai, Chernoff, \\& Cordes 2001; Brisken et al. 2003 for reviews), have shown that the supernova explosion is almost certainly asymmetric, the level of asymmetry in the engine is very difficult to determine from these indirect methods. Gravitational waves and neutrinos provide the only means to directly probe the supernova engine.  Neutrinos have only been detected in SN 1987A (Hirata et al. 1987; Bionta et al. 1987), and gravitational waves (GWs) have yet to be conclusively detected in supernovae.  Unfortunately, the signals of both are sufficiently weak to limit their detection to Galactic or near-Galactic supernovae (Neutrinos: Burrows, Klein \\& Gandhi 1998; GWs: Fryer, Holz \\& Hughes: FHH 2002; Dimmelmeier, Font \\& M\\\"uller 2002; Kotake, Yamada \\& Sato 2003b; Fryer \\& Warren 2004; M\\\"uller et al. 2004).  Although predictions for the GW signal from stellar collapse have been much higher in the past (see FHH 2002; New 2003 for reviews), these predictions were based on stellar cores rotating much faster than even the fastest rotating supernova progenitors produced today (Heger, Langer \\& Woosley 2000).  Modern progenitors of stellar collapse are only marginally, if at all, unstable to bar instabilities, and strong bar modes (the source of strong gravitational wave signals) do not occur in collapse calculations using these progenitors (Fryer \\& Warren 2004).  Even if bars did form, they form at low densities, and at most produce a marginally detectable signal at 10\\,Mpc (Rampp, M\\\"uller \\& Ruffert 1998; FHH 2002).  GW observations, like neutrinos, are limited to Galactic or near-Galactic supernovae.  Even though the rate of such nearby SNe are low (0.02 year$^{-1}$), the insight gained from these supernovae will provide essential information into stellar collapse.  In this paper, we will focus on 3 phases of gravitational waves: \\begin{itemize} \\item[I)] {\\bf Stellar Bounce} - The signal produced during the bounce caused when the stellar core reaches nuclear densities and its runaway free-fall is abruptly halted.  For initial progenitors with asymmetries or rotation, the quadrupole moment in the progenitors will vary wildly during bounce, producing a strong signal. \\item[II)] {\\bf Convection} - The signal produced during the convective phase of the supernova engine.  Even if the initial conditions are symmetric, the convection modes may merge to form low-mode convection with a significant time-variable quadrupole moment. \\item[III)] {\\bf Neutrinos} - The signal produced by asymmetries in the neutrino emission.  Asymmetries in the collapse or in the convection will lead to asymmetries in the neutrino emission.  These asymmetries lead to a time-variable quadrupole moment. \\end{itemize} The GW signal can not only be used to distinguish these phases, but also the asymmetry that causes the signal in each phase.  In this paper, we will use the results from recent 3-dimensional collapse simulations (Fryer \\& Warren 2002, 2004; Fryer 2004a) to study the different characteristics of these signals, focusing especially on what these signals can tell us about stellar collapse and supernova explosions.  But a lot of uncertainty remains in such simulations and we use comparisons with other work and analytic estimates to gauge the uncertainties in these calculations.  The models, and the numerical methods used to calculate the GW signal, are discussed in \\S 2.  \\S 3 shows the results of these simulations, along with comparisons to other work and analytic estimates.  To make GWs a powerful probe of the inner core, we must study a bigger picture and correlate the GW signal with other supernova observables: neutrino signal, nucleosynthetic yields and other explosion asymmetry characteristics in the observations (\\S 4).  We conclude with a discussion of the observability of gamma-ray bursts under the collapsar model and a review of the key observations and the constraints they can place on the supernova engine. ",
        "conclusions": "A number of asymmetries in the stellar collapse can produce a gravitational wave signal in a Galactic supernova that should be detectable by advanced LIGO.  These asymmetries are caused by rotation in the star, asymmetries induced by explosive nuclear burning just prior to collapse, or low mode convection in the supernova engine. Each has a distinctive gravitational wave signal which can be used to study the mass motions in the supernova engine.  But the real power of GWs arises when combined and corroborated with other observations.  The correlation between GW and these observations is summarized in Tables 2 and 3.  Fast rotating progenitors can produce a strong signal at bounce, mild asymmetries in the ejecta and fast-spinning neutron star remnants.  The timing of the GW signal is within 5\\,ms (maybe even closer) of the neutrino signal and such an event could be used to constrain the neutrino mass.  Asymmetries in the collapsing star caused by explosive burning can produce strong signals at late times through asymmetric neutrino emission and mild velocities on the neutron star, but the explosion is roughly symmetric.  Low mode convection will produce a gravitational wave signal during the convective engine phase, asymmetric explosions and possibly strong neutron star kicks.  The GW signal will help determine the mechanism behind the SN and our understanding of these additional phenomena.  Supernovae are not the only explosions produced by stellar collapse. The favored mechanism for long-duration gamma-ray bursts, the collapsar model (Woosley 1993; MacFadyen \\& Woosley 1999), invokes the collapse of a massive star.  Collapsars are massive stars that do not produce strong supernova explosions, but instead collapse to form black holes.  If the star is rotating fast enough, the high angular momentum stellar material will form an accretion disk around the black hole.  Energy liberated from the disk drives a strong explosion (even stronger than supernovae) in a jet along the rotation axis.  Because rotation is required to produce a collapsar explosion, the collapsing stars will necessarily be fast rotators.  Thus they will have a strong bounce GW signal.  However, this signal is unlikely to be much stronger than our fast rotating collapse models (too much rotation will also prevent the collapsar engine from working - MacFadyen \\& Woosley 1999).  It is likely that collapsars arise from extremely massive stars (Heger et al. 2003), but, because the cores of massive stars are essentially all the same, this does not alter the GW signal significantly from our extreme rotating case (Dupuis, Fryer, \\& Heger 2004).  The bounce of collapsar driven gamma-ray bursts will only be detectable when they occur in the galaxy.  The accretion onto the black hole will cause ringing, but this signal too will only be detectable in a Galactic gamma-ray burst.  Finally, for the low-mass disks produced in collapsars (Popham, Fryer, \\& Woosley 1998; MacFadyen \\& Woosley 1999), no disk instabilities will form and the accretion phase of collapsars are unlikely to produce strong GW signals.  The narrowly collimated, low-baryon jet produced in collapsars will produce a signal akin to neutrino driven signals. However, for a bipolar jet, there will be no signal along the jet axis.  Perpendicular to the jet axis, the signal (using equation 8) can be as high as $5 \\times 10^{-24}$ for a $10^{51}$\\,erg jet, on par with the signal from the disk, but still only detectable within the Galaxy and only if the jet is not directed toward us (so not associated with GRBs, but maybe with hypernovae).  What this all boils down to is that collapsar driven gamma-ray bursts, like supernovae, will not be detectable beyond the Galaxy.  Given their low rate ($\\sim 10^{-5} {\\rm year^{-1}}$ in the Galaxy), collapsars will not be a strong GW source.  However, other gamma-ray burst models will produce strong GW signals (e.g. neutron star mergers) and GWs can easily distinguish these two burst engines."
    },
    "0403/nucl-th0403019_arXiv.txt": {
        "abstract": "We analyse the temperature dependence of pairing correlations in the inner crust matter of neutron stars. The study is done in a finite-temperature HFB approach and by using a zero range pairing force adjusted to the pairing  properties of infinite neutron matter. Within the same approach we investigate how the specific heat of the inner crust depends on temperature,  matter inhomogeneity, and the assumption used for the pairing force. It is shown that in a physical relevant range of densities the pairing properties of inner crust matter depend significantly on temperature. The finite-temperature HFB calculations  show also that the specific heat is rather sensitive to the presence of nuclear clusters inside the inner crust. However, the most dramatic change of the specific  heat is determined  by the scenario used for the neutron matter  superfluidity. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403627_arXiv.txt": {
        "abstract": "BL~Lacertae was the target of an extensive multiwavelength monitoring campaign in the second half of 2000. The campaign had revealed optical and X-ray intraday variability on time scales of $\\sim 1.5$~hours and evidence for significant spectral variability both at optical and X-ray frequencies. During the campaign, BL~Lacertae was observed in two different activity states: A quiescent state with relatively low levels of optical and X-ray fluxes and a synchrotron cutoff at energies below the X-ray regime, and a flaring state with high levels of optical and X-ray emission and a synchrotron cutoff around or even beyond $\\sim 10$~keV. In this paper, we are using both leptonic and hadronic jet models to fit the broadband spectra and spectral variability patterns observed in both activity states in 2000. We start out with global spectral models of both activity states. Subsequently, we investigate various flaring scenarios for comparison with the observed short-term variability of BL Lacertae in 2000. For our leptonic jet model, we find that the short-term variability, in particular the optical and X-ray spectral variability, can be best represented with a flaring scenario dominated by a spectral-index change of the spectrum of ultrarelativistic electrons injected into the jet. Based on this result, a detailed model simulation of such a flaring scenario, reproducing the observed optical and X-ray spectral variability and broadband SED of BL Lacertae during the BeppoSAX pointing around Nov. 1, 2000, simultaneously, is presented. Our leptonic modeling results are compared to fits using the hadronic synchrotron-proton blazar (SPB) model. That model can reproduce the observed SEDs of BL~Lacertae in a scenario with $\\mu$-synchrotron dominated high-energy emission. It requires a significantly higher magnetic field than the leptonic model ($\\sim 40$~G vs. $\\sim 2$~G in the leptonic model) and a lower Doppler factor associated with the bulk motion of the emission region ($D \\sim 8$ vs. $D \\sim 18$ in the leptonic model). The hadronic model predicts a significantly larger $\\gtrsim 100$~GeV flux than the leptonic models, well within the anticipated capabilities of VERITAS and MAGIC. ",
        "introduction": "Introduction}  BL Lacertae (= 1ES~2200+420; z = 0.069) was historically the prototype of the BL Lac class of active galactic nuclei (AGN). These objects are characterized by continuum properties similar to those of flat-spectrum radio quasars (non-thermal optical continuum, high degree of linear polarization, rapid variability at all wavelengths, radio jets with individual components often exhibiting apparent superluminal motion), but do usually show only weak emission lines (with equivalent width in the rest-frame of the host galaxy of $< 5$~\\AA), if any. In BL~Lacertae itself, however, H$\\alpha$ (and H$\\beta$) emission lines with equivalent widths in excess of 5~\\AA\\ have been detected during a period of several weeks in 1995 \\citep{ver95,cor96}, and in 1997 \\citep{cor00}. Superluminal motion of $\\beta_{\\rm app}$ up to $7.1 \\pm 0.3$ has been observed in this object \\citep{denn00}.  BL~Lacertae is classified as a low-frequency peaked BL~Lac object \\citep[LBL;][]{fossati98}. From an interpolation between the GHz radio spectrum and the IR -- optical spectrum, it can be inferred that its low-frequency spectral component typically peaks at mm to $\\mu$m wavelengths, while the high-frequency component seems to peak in the multi-MeV -- GeV energy range. BL~Lacertae has been studied in detail during various intensive multiwavelength campaigns \\citep[e.g.][]{bloom97,sambruna99,madejski99,ravasio02,villata02,boettcher03}. It is a particularly interesting object for detailed X-ray studies since it is in the X-ray regime where the two broad components of the multiwavelength SEDs of BL~Lacertae (and other LBLs) are overlapping and intersecting. X-ray observations of this source at different epochs show significant flux and spectral variability, indicating that the X-ray emission is at times dominated by the high-energy end of the synchrotron emission, while at other occasions it is dominated by the low-frequency portion of the high-energy bump of the SED. In fact, BL~Lacertae has repeatedly shown a concave shape \\citep[e.g.,][]{madejski99,ravasio02}, with rapid variability mainly restricted to the low-energy excess portion of the spectrum \\citep[e.g.,][]{ravasio02,ravasio03}.  In the framework of relativistic jet models, the low-frequency (radio -- optical/UV) emission from blazars is interpreted as synchrotron emission from nonthermal electrons in a relativistic jet. The high-frequency (X-ray -- $\\gamma$-ray) emission could either be produced via Compton upscattering of low frequency radiation by the same electrons responsible for the synchrotron emission \\citep[leptonic jet models; for a recent review see, e.g.,][]{boettcher02}, or due to hadronic processes initiated by relativistic protons co-accelerated with the electrons \\citep[hadronic models, for a recent discussion see, e.g.,][]{muecke01,muecke03}. The lack of knowledge of the primary jet launching mechanism and the difficulty in constraining the jet composition from general energetics considerations \\citep[for recent discussions see, e.g.,][]{sm00,gc01} are currently leaving both leptonic and hadronic models open as viable possibilities. Also, detailed simulations of particle acceleration at relativistic shocks or shear layers, which may be relevant for the acceleration of ultrarelativistic particles in blazar jets, show that a wide variety of particle injection spectra may result in such scenarios \\citep[e.g.,][]{ob02,so03}, greatly differing from the standard spectral index of 2.2 -- 2.3, which was previously believed to be a universal value in relativistic shock acceleration scenarios \\citep[e.g.,][]{achterberg01,gallant99}. Thus, both the nature of the matter in blazar jets and the energy spectra of ultrarelativistic particles injected into the emission regions in blazar jets are difficult to constrain from first principles. For this reason, we are leaving these aspects as virtually free parameters in our models, and attempt to constrain them through the results of detailed time-dependent modeling of blazar emission.  While simultaneous broadband spectra are very useful to constrain both types of blazar jet models, there still remain severe ambiguities in their interpretation w.r.t. the dominant electron cooling, injection, and acceleration mechanisms, as was recently illustrated for the case of W~Comae by \\cite{bmr02}. Those authors have also demonstrated that a combination of broadband spectra with timing and spectral variability information, in tandem with time-dependent model simulations \\citep{krm98,gm98,kataoka00,ktl00,lk00,sikora01,bc02,krawczynski02,kusunose03} can help to break some of these degeneracies. For this reason, we had organized an intensive multiwavelength campaign to monitor BL~Lacertae in the second half of 2000 at radio, optical, X-ray, and very-high-energy (VHE) $\\gamma$-ray frequencies, putting special emphasis on detailed variability information. The results of this multiwavelength campaign were published in \\cite{boettcher03}; for more details on the optical and X-ray observations, see also \\cite{villata02} and \\cite{ravasio03}, respectively. In \\S \\ref{observations}, we briefly highlight the main results of that campaign, emphasizing those aspects that we will specifically use here to constrain our model calculations.  The purpose of this paper is to use both leptonic and hadronic jet models to fit the spectral energy distributions (SEDs) and spectral variability of BL~Lacertae observed in 2000. Following a brief description of both leptonic and hadronic models in \\S \\ref{model_descriptions}, we will first present spectral fits to the SEDs of BL~Lac in \\S \\ref{SED_modeling}. Our code used to evaluate leptonic models allows us to make detailed simulations of the spectral variability resulting from different plausible flaring scenarios. In \\S \\ref{var_modeling}, we will first do a general parameter study of various scenarios and compare the results qualitatively with the observed spectral variability trends in BL~Lacertae (\\S \\ref{parameter_study}). This will allow us to decide on a preferable flaring scenario, for which we subsequently optimize our choice of model parameters to fit simultaneously the SED and spectral variability of BL~Lacertae as observed in 2000 (\\S \\ref{var_fit}). Possible physical scenarios triggering the inferred variability of the electron injection spectrum will be discussed in \\S \\ref{discussion}. We summarize in \\S \\ref{summary}.  Throughout this paper, we refer to $\\alpha$ as the energy spectral index, $F_{\\nu}$~[Jy]~$\\propto \\nu^{-\\alpha}$. A cosmology with $\\Omega_m = 0.3$ and $\\Omega_{\\Lambda} = 0.7$ and $H_0 = 70$~km~s$^{-1}$~Mpc$^{-1}$ is used. ",
        "conclusions": "Discussion}  Overall, our parameter choice for the leptonic models for BL~Lacertae are in reasonable agreement with those found by other authors based on earlier multiwavelength campaigns on this source \\citep[e.g.,][]{madejski99,bb00,ravasio03}. In agreement with those authors, our best-fit Doppler factor of $D = 18$ is well within the range typically found for blazar modelling, and the magnetic field of $B = 2$~G is intermediate between typical values found for leptonic modeling of flat-spectrum radio quasars (FSRQs) and high-frequency peaked BL~Lac objects (HBLs). While FSRQs are usually successfully modelled with $B \\gtrsim $~a few G \\citep[e.g.,][]{vmontigny97,sambruna97,mukherjee99,ghisellini99,hartman01}, typical values found for HBLs are of the order of $B \\lesssim 0.1$~G \\citep[e.g.][]{tmg98,kataoka99,ca99,kataoka00,petry00,krawczynski02}.  In our analysis of the spectral variability patterns, we have found that those patterns can successfully be modelled with a fluctuation of the electron injection spectral index. Remarkably, our time-dependent fits indicate that an injection index larger than $q \\sim 2.3$, even during the peak of an individual short-term flare, is required. If the injection of ultrarelativistic electrons into the emitting volume is caused by Fermi acceleration at relativistic shocks, detailed numerical studies have shown that with fully developed turbulence in the downstream region, a unique asymptotic index of $q \\sim 2.2$ -- 2.3 should be expected \\citep[e.g.,][]{achterberg01,gallant99}. However, recently \\cite{ob02} have shown that Fermi acceleration might lead to drastically steeper injection spectra if the turbulence is not fully developed. Furthermore, depending on the orientation of the magnetic field at the shock front, an abrupt steepening of the injection spectra may result if the shock transits from a subluminal to a superluminal configuration. In this context, our leptonic fit results may indicate that such predominantly geometric effects, may be the cause of the rapid variability observed in BL~Lacertae.  In their analysis of the {\\it BeppoSAX} + optical continuum spectra of the Oct. 31 -- Nov. 2 observations, \\cite{ravasio03} have noticed that the time averaged optical and LECS + MECS X-ray spectra can not be connected smoothly using a single power-law or a smoothly connected broken power-law. They have suggested and investigated several possibilities how this discrepancy could be remedied, including a variable dust-to-gas ratio, the bulk Compton process \\citep{sikora97}, a multi-component emission model, and a spectral upturn resulting from Klein-Nishina effects on the electron cooling rates. The results of our combined leptonic spectral + variability modeling of BL~Lacertae suggest that the discrepancy ultimately arises artificially as a result of the time averaging involved when producing the high-quality {\\it Beppo}SAX LECS + MECS spectrum. The fact that our time-dependent leptonic fits reproduce the observed ranges of optical and X-ray fluxes and spectral indices simultaneously resolves the issue of this optical -- X-ray spectral discrepancy, and removes the need for any additional assumptions concerning intergalactic absorption and/or electron populations in the jet.  In contrast to leptonic models, the hadronic SPB model required significantly larger magnetic field strengths (of order 20 -- 40~G) on the length scale of the size of the emission region of $\\sim 10^{15}$~cm. The range of Doppler factors appears about a factor of 2 lower than in the leptonic models. The total jet power, which turns out to be below the estimated accretion disk luminosity, remains larger in hadronic models than in leptonic ones owing to the higher particle and field energy content. The state transition from low to higher activity in 2000 is well described by an increase of the co-moving particle energy and the bulk Lorentz factor. In the picture of diffusive shock acceleration (in the test particle limit) the maximum particle energies are related to the magnetic turbulence spectrum \\citep[see, e.g.,][and references therin]{bs87,reimer04}. The required maximum electron and proton energies in the presented models can be understood if the particles gain energy by diffusive shock acceleration in a $\\propto k^{-1.1\\ldots -1.3}$ turbulence spectrum where $k$ is the wave number in the turbulent magnetic field. Spectral-index changes can not be ruled out, but they are not the dominant cause of the spectral variability of BL~Lacertae in the framework of the presented modeling using the SPB model.  An interesting diagnostic for the particle content in the jet -- in addition to any neutrino detections -- might be achieved through high sensitivity observations in the (sub-)TeV energy range by e.g. MAGIC or VERITAS. While both, leptonic and hadronic models, predict a similar flux level in the GeV-energy range, hadronic models predict about an order of magnitude higher flux values than leptonic ones do above 40 GeV for BL Lacerate in 2000. Furthermore, considering the results of our variability study, the predicted VHE $\\gamma$-ray flux from a leptonic jet only reaches the peak level mentioned above during short flares. It depends critically on the duty cycle of such flaring events whether a sufficient time-averaged level of VHE flux can be sustained for low-energy-threshold Cherenkov telescopes like MAGIC to accumulate a measurable signal.  Our leptonic SED + spectral variability fit predicted spectral hysteresis at soft X-ray energies which might serve as a confirmation of our fit results. The limited count statistics of our {\\it Beppo}SAX observations did not firmly establish nor rule out the existence of X-ray spectral hysteresis. More sensitive, dedicated observations by {\\it Chandra} and/or {\\it XMM-Newton} would be extremely helpful to test this prediction. According to our hadronic model fits presented in this paper, flares of BL~Lacertae were primarily caused by increasing particle energies. If this is indeed the dominant flaring mechanism and it is not accompanied by significant changes of the electron spectral index, then short-term X-ray spectral variability might be reasonably well represented by the patterns resulting from our leptonic models with increasing $\\gamma_2$ with at most very moderate spectral-index changes. These did not show significant spectral hysteresis. However, fluctuations of the electron injection spectral index could not be excluded in our hadronic fits. Consequently, the presence of pronounced soft X-ray spectral hysteresis in BL~Lacertae may slightly favour the leptonic models, while its absence would favour hadronic models, but the discriminating power of such a measurement in the case of BL~Lacertae would be rather limited."
    },
    "0403/astro-ph0403411_arXiv.txt": {
        "abstract": "We propose two hypotheses which characterize the quasi-equilibrium state that realizes after a cold collapse of self-gravitating system. The first hypothesis is the linear temperature-mass (TM) relation, which yields a characteristic  non-Gaussian velocity distribution. The second  is the local virial (LV) relation, which, combining the linear TM relation, determines a unique mass density profile as $\\rho(r)={\\rho }% _{0}r^{-4}e^{-r_0/r}$. Although this density profile is unphysical in the central region, the region is just inner a few  percent around the center of a bound region in cumulative mass, which is beyond the resolution of our numerical simulations. Hence posing two hypotheses is compatible to the numerical simulations for almost the whole region of the virialized bound state. Actually, except for this inner part, this density profile fits well to  the data of cold collapse simulations.  Two families of spherical and isotropic models, polytropes  and King models, are examined from a view point of these two hypotheses. We found that the LV relation imposes a strong constraint on these models: only polytropes with index $n \\sim 5$ such as Plummer's model are compatible with the numerical results characterized by the two hypotheses among these models. King models with the concentration parameter $% c \\sim 2$ violate the LV relation while they are consistent with the $R^{1/4} $ law for the surface brightness. Hence the above characteristics can serve as a guideline to build up the models for the bound state after a cold collapse, besides the conventional criteria concerning the asymptotic behavior. ",
        "introduction": "Collisionless self-gravitating systems (SGS) eventually settle down to a quasi-equilibrium state through the phase mixing and the violent relaxation processes under the potential oscillation\\cite{Lynden67}. This quasi-equilibrium state is a prototype of the astronomical objects such as galaxies and clusters of galaxies. These objects often show universal profiles in various aspects.  It would be true that the energy distribution function and the mass density distribution function are quite appropriate to characterize SGS, and actually many studies have been done from this viewpoint. For example, the density profile is found to be $\\rho \\propto r^{-4}$ in spherical and cold collapses, \\textit{i.e.,} collapses with small initial virial ratio \\cite{Henon64,Albada82} and $\\rho \\propto r^{-3}-r^{-4}$ in the cluster-pair merging processes\\cite{Merrall03}.  On the other hand, the velocity distribution function has not fully been examined so far despite its importance \\footnote{% So far, full velocity distributions have been mainly analyzed only for the line-of-sight velocity profiles because they can be directly compared with the observational data \\cite{Dejonghe87,Gerhard93}}; anisotropy of velocity dispersion or lower moments of velocity distributions with respect to the deviations from Gaussian distributions have been mainly examined (e.g. \\cite{Binney82,Gerhard93}). However, the full velocity distribution function plays an important role in identifying the quasi-equilibrium state, which should be properly characterized by the whole phase-space probability functions. Merrall \\textit{et al.} numerically showed that the velocity distributions in the central regions of the bound particles become Gaussian, when the bound particles are well relaxed \\cite{Merrall03}. However, since SGS are manifestly non-additive systems \\cite{Nakamichi04}, the ordinary Gaussian distribution would not be expected for the whole bound particles. Moreover, such non-Gaussian velocity distributions would be firmly connected with the density profile characteristic to SGS % \\cite{Gerhard93,Kazantzidis04}.  Recently, we have simulated self-gravitating N-body systems by utilizing the leap-frog symplectic integrator on GRAPE-5, a special-purpose computer designed to accelerate N-body simulations\\cite{GRAPE}. The simulations were performed mainly in the case with cold collapse initial conditions. It is certain that the processes may be too simplified to deduce the ultimate nature of actual SGSs, in which merging or tidal effects take place. However, one of our main goals is to establish the universal properties of SGS through the violent relaxation process, i.e. the relaxation through the oscillation of gravitational potential. This process is rather common in actual SGSs such as galaxies and dark matters. In this context, the cold collapse is considered to be a typical case which causes the violent relaxation. Hence it is meaningful to examine the cold collapse as a first step.  When the system experiences violent gravitational processes such as a cold collapse and a cluster-pair collision, we obtained a universal velocity distribution profile expressed as the democratic (=equally weighted) superposition of Gaussian distributions of various temperatures(DT distribution hereafter) in which the local temperature $T(r)$ is defined using the local velocity variance $\\langle v^{2}\\rangle $ as $T(r):=m\\langle v^{2}\\rangle (r)/3k_{B}$, where $m$ is the particle mass \\cite{Osamu04}. Moreover, we have found that the locally defined temperature linearly falls down in the intermediate cluster region outside the central part, provided it is described against the cumulative mass $M_{r} $, \\textit{i.e.,} $dT/dM_{r}=const.$ This fact is consistent with the appearance of DT distributions.  Furthermore to the linear TM relation, we have also obtained another peculiar fact that the LV relation between the locally defined potential energy and kinetic energy holds except for the weak fluctuations. More precisely, the local temperature $T(r)$ is proportional to the local potential $\\Phi (r)$, with constant proportionality $6k_{B}T(r)=-m\\Phi (r)$ for a wide class of cold collapse simulations. It was originally pointed out by Eddington that Plummer's model satisfies this condition exactly  \\cite{Eddington16}. Moreover it was also pointed out that Plummer's model can be extended to the anisotropic model of a stellar system satisfing the LV relation \\cite{Evans05}. We found that the bound region, \\textit{i.e.,% } the region composed of the bound particles % keeps this relation quite well at least during the time interval much longer than the initial crossing time. This is a prominent correlation in the phase space distribution function, and also constructs another backbone character of SGS after a cold collapse.  In sec.\\ref{sec:level2} and \\ref{sec:level3}, we will show that the above two hypotheses  for SGS are supported by the results of a variety of cold collapse simulations. Then in sec.\\ref{sec:level4}, combining the two hypotheses, \\textit{i.e.,} the linear TM relation and the LV relation, we obtain a unique density profile $\\rho (r)=\\rho_{0}(r/r_{0})^{-4}e^{-r_{0}/r}$ which asymptotically behaves as $\\rho\\propto r^{-4}$. Justification of this density profile will also be discussed. In sec.\\ref{sec:relaxation}, we will ascertain that the above two hypotheses can arise even in the collisionless quasi-equilibrium state, where the region inside a half-mass radius is not yet under the effect of two-body relaxation. In sec.\\ref{sec:models} the above two hypotheses will be explored in two typical collisionless static models with spherical configuration and isotropic velocity dispersions; polytropes and King models. Sec.\\ref{sec:conclusions} is devoted to the summaries and conclusions of this paper. % ",
        "conclusions": "\\label{sec:conclusions} In this paper, we have proposed (a) the linear TM relation and (b) the LV relation as the hypotheses of self-gravitating structures after a cold collapse. Based on these two hypothetical proposals, we have uniquely determined the mass density profile (\\ref{sotakai}) in the quasi-equilibrium state.  Our density profile behaves as $\\rho \\sim r^{-4}$ in the outer region and fits well both to the data of cold collapse simulations. Although this density profile is unphysical in the inner region, where it decreases with the decreasing radius, we showed that such an unphysical region is just a few percent around the central region of the bound state. Hence, we can safely neglect this region at least for the resolution of our numerical results.  In the central region, we also have to notice the relaxation process in such a central region. In fact, we examined how two-body relaxation is developed in a  bound state after a cold collapse. We estimated the time scale of the  two-body relaxation with $t^{\\ast}_{rel}$, which is defined as the time when the region within a half-mass radius achieves a two-body relaxation on average. Both the linear TM relation and the LV relation appear even when the two-body relaxation is not attained. Hence, we can say that our two hypotheses are not affected through the two-body relaxation but are characteristic even in collisionless system in our estimation with $t^{\\ast}_{rel}$.  From a viewpoint of these two hypotheses, we compared our results with two isotropic static models,i.e., polytropes and King models. The polytrope with $n=5$ is special among those models, since it exactly satisfies the LV relation. Plummer's model, the polytrope with $n = 5$ with a flat boundary condition at the center also satisfies the constant TM relation quite well in the intermediate region $0.2M_{tot}<M_{r}<0.8M_{tot}$. Moreover, its density fits well to the numerical data quite well in this intermediate region. On the other hand, both its TM relation and density profile in the outer region are slightly shifted  from those in cold collapse simulations: the temperature falls off more steeply and the density behaves as $r^{-5}$ in the outer region. As for King models, in contrast, there is no parameter range in which both the TM relation and density profile fit well to the numerical data even in the intermediate region.  We also examined the deviations from the two hypotheses quantitatively for these two models. Only the polytropes with $n \\sim 5$ are acceptable in that their deviations from both of the hypotheses are compatible with the numerical results of cold collapse. The LV relation imposes a strong constraint for these models. Especially for King models, the deviation from LV relation takes the maximum value around the parameter $c\\sim 2$, in which the surface brightness fits well to $R^{1/4}$ law. This suggests that considering only the asymptotic behavior of the bound state is not enough to describe the collisionless static models to express observational or numerical results.  Combining the results for the above two models, we can speculate that Plummer's model follows both of the hypotheses  we proposed (a) LV relation and (b) linear TM relation. However, this model is  lack of the following properties of the bound states in the actual simulations of cold collapse: (i) deviation from a hydrostatic equilibrium state and (ii) anisotropy of velocity dispersion in the outer region.  As for (i), the hydrostatic equilibrium condition is not attained especially when the violent degree of the initial collapse is strong, since the bound state slowly expands outward and is not balanced with a gravitational force after a cold collapse. Hence, to express these conditions we may have to consider not the static solution but the stationary solution admitting the coherent motion of matters.  As for (ii), we derived that for a mild collapsing case, the LV relation fluctuates more largely in the outer region and the slope of the TM relation at this region becomes a little steeper at least in the collisionless stage. We suspect that these results are related to the anisotropy of this system.  Hence for more precise description of the system possessing two hypotheses, it seems necessary to treat the model with the anisotropy of the velocity dispersion. So far  several collisionless static models with the anisotropic velocity dispersions or with the non-spherical symmetry have been proposed \\cite{Henon73,Evans05,Bertin84,Stiavelli85,Dejonghe87,Gerhard91,Bertin03}. Incompleteness of the violent relaxation or the existence of conservative quantities during the phase mixing seem to be a candidate as the model to induce such anisotropy \\cite{Lynden67,Stiavelli85}. The hypotheses we proposed are expected to be used as a guideline to build up the models following the characters of the bound states after a cold collapse by combining with these anisotropic models.  It also seems important how widely our two hypotheses are applicable in more general mixing processes including a merging process which yields a cuspy density profile in the central region \\cite{Navarro96}. If we want to know the detail of such a tiny central region, such as the existence of the cuspy density profile, we have to treat  such a central region more carefully. We speculate that in such a region, at least one of the two hypotheses we proposed will be broken. These points are now under investigation."
    },
    "0403/astro-ph0403282_arXiv.txt": {
        "abstract": "Cosmological ($\\Lambda$CDM) TreeSPH simulations of the formation and evolution of galaxy groups and clusters have been performed. The simulations invoke star formation, chemical evolution with non-instantaneous recycling, metallicity dependent radiative cooling, strong star-burst, and (optionally) AGN, driven galactic super-winds, effects of a meta-galactic UV field and thermal conduction. Results for two clusters, one Virgo-like ($T$$\\simeq$3 keV) and one (sub) Coma-like ($T$$\\simeq$6 keV), are presented. At $z$=0 the stellar contents of both clusters consist of a central dominant (cD) galaxy surrounded by cluster galaxies and intracluster stars. The intracluster (IC) stars are found to contribute 20-40\\% of the total cluster B-band luminosity at $z$=0 and to form at a mean redshift $\\bar{z}_f\\sim$3 being on average about 0.5 Gyr older than the stars in cluster galaxies. UBVRIJHK surface brightness profiles of the IC star populations are presented; the profile of the larger cluster matches the observed V-band profile of the cD in Abell 1413 ($T\\simeq$8 keV) quite well. The typical colour of the IC stellar population is B-R=1.4-1.5, comparable to the colour of sub-$L^*$ E and S0 galaxies. The mean Iron abundance of the IC stars is about solar in the central part of the cluster ($r\\sim$100 kpc) decreasing to about half solar at the virial radius. The IC stars are $\\alpha$-element enhanced with a weak trend of [O/Fe] increasing with $r$ and an overall [O/Fe]$\\sim$0.4~dex, indicative of dominant enrichment from type II supernov\\ae.  The IC stars are kinematically significantly colder than the cluster galaxies: The velocity dispersions of the IC stars are in the inner parts of the clusters ($r\\sim$100-500 kpc) only about half of those of the cluster galaxies increasing slightly outward to about 70\\% at $r$=1-2 Mpc. The typical projected velocity dispersion in the Virgo-like cluster at $R\\ga$50 kpc is 300-600 km/s depending on orientation and projected distance from the cluster center. Rotation is found to be dynamically insignificant for the IC stars. The velocity distributions of IC stars {\\it and} clusters galaxies are in one cluster highly radially anisotropic, in the other close to isotropic. ",
        "introduction": "Current hierarchical or ``bottom-up'' structure formation theories predict that % the field star populations of haloes of galaxies like the Milky Way should consist partly of stars originally born in small proto-galaxies and later tidally stripped from these by the main galaxy or through interaction with other proto-galaxies. The fraction of halo field stars of such origin depends (apart from obvious cosmic variance effects) on the properties of the dark matter (e.g., cold dark matter, CDM, vs. warm dark matter, WDM: for WDM the hierarchy is broken at the free-streaming mass scale, e.g., Sommer-Larsen \\& Dolgov 2001) and the nature of the star-formation process. For example, some halo stars may be born in rapidly expanding, supernova driven proto-galactic supershells and hence be detached from their parent proto-galaxies right from the outset in space, as well as phase-space \\citep[e.g.,][]{M.97,SL.03}. If dark matter is cold and star-formation in expanding supershells is not a common phenomenon the above fraction could be quite large, in principle at least, approaching unity \\citep[e.g.,][]{H.03}. The halo stars resulting from tidal stripping or disruption of a proto-galaxy will stay localized in phase-space for a long period and several such ``streams'' of halo stars have been detected in the haloes of the Milky Way and M31 (e.g., Helmi et al. 1999; Ferguson et al. 2002).  From the point of view of structure formation, galaxy clusters can be seen as scaled up versions of galaxies in hierarchical scenarios. In particular, gravity is expected to strip or disrupt cluster galaxies in a similar way as satellite galaxies in galaxy haloes and a population of intracluster (or ``field'') stars should thus reside between the cluster galaxies. It has been known for decades that cD galaxies often are embedded in extended envelopes presumed to consist of stars tidally stripped off galaxies in the process of being engulfed by the cD \\citep[e.g.,][]{O76,D79}. In recent years it has been possible to perform quantitative studies of these stellar envelopes through ultra-deep surface photometry of the general stellar population \\citep[e.g.,][]{G.00,G.04,F.02,F.04a}, or imaging/spectroscopy of individual planetary nebulae \\citep[e.g.,][]{A.02,A.03,F.04b} or supernovae Type Ia \\citep{G.03}. The potential importance of intracluster stars in relation to the chemical enrichment of the intracluster medium (ICM) was recently discussed by \\cite{Z.04} and Lin \\& Mohr (2004).  On the theoretical front \\cite{N.03} have used an N-body dark matter only fully cosmological simulation of the formation of a Virgo-like cluster to make predictions about intracluster stars. They find that unrelaxed velocity distributions and (bulk) streaming motions of the IC stars should be common due to the large dynamical timescales in clusters. Hence the evolution towards a relaxed coarse-grained phase-space distribution function through phase-mixing has proceeded much less far\\footnote{A visualized example from one of our simulations is given at http://www.tac.dk/\\~~\\hspace{-1.4mm}jslarsen/ICstars} than, e.g., in the Galactic halo, where dynamical timescales are much shorter.  The dark matter only simulations are complemented by various N-body simulations which invoke a more realistic modelling of the stellar properties of galaxies in order to study the fate of galaxies in a cluster environment (e.g., ``galaxy harassment'', Moore et al. 1996, ``tidal stirring'', Mayer et al. 2001 and the formation of the central cluster galaxy, Dubinski 1998). The advantage of this approach is that a very high resolution in the stellar component can be obtained. In relation to the properties of the IC stars, recent progress has been made by \\cite{D.03}, and \\cite{F.04a}.  In general the properties of the system of IC stars are set by two main effects: a) the cool-out of gas and subsequent conversion of cold, high-density gas to stars in individual galaxies and b) the stripping/disruption of the galaxies through interactions with other galaxies and the main cluster potential. Since such interactions will generally affect the star-formation rate (as long as a reservoir of gas is available) the former process is intimately coupled to the latter.  Only fairly recently has it been possible to carry out fully cosmological gas-dynamical/N-body simulations of the formation and evolution of galaxy clusters at a sufficient level of numerical resolution and physical sophistication that the cool-out of gas, star-formation, chemical evolution and gas inflows and outflows related to individual cluster galaxies can be modelled to, at least some, degree of realism \\citep[e.g.,][]{V03,T.04}. Recently \\cite{M.04} analyzed 117 clusters formed in a large cosmological TreeSPH simulation \\citep{B.04} to study the properties of the IC stars. They determined average density profiles of the IC stars as well as stars in galaxies for their large sample of clusters. One of the interesting conclusions reached is that only at fairly large cluster-centric distances ($r\\ga$0.4-0.5$r_{\\rm{vir}}$) does the (spherically averaged) density of stars in galaxies (excluding the cD) exceed that of IC stars. Complementary work has recently been performed by \\cite{W.04}, who simulated a Coma-like cluster at very high numerical resolution. They find an IC star fraction of 10-20\\%, in good agreement with observations, and a tendency for the fraction to increase with time. In relation to on-going observational efforts, they moreover find that line-of-sight velocities of intra-cluster planetary nebulae can be used for fairly accurate cluster mass estimation, even for dynamically quite unrelaxed clusters, if more than 5 fields at a range of projected radii are covered.  We have recently undertaken similar simulations of galaxy groups and clusters. We are building on our TreeSPH code used for simulating galaxy formation \\citep[e.g.,][]{SL.03}, improved to include modelling of non-instantaneous chemical evolution \\citep{L.02}, metallicity-dependent, atomic radiative cooling, strong supernova, and (optionally) AGN, driven galactic winds and thermal conduction. The main results of the simulations concerning the properties of the ICM and cluster scaling relations are described in Paper~I \\citep{R.04}, while the characteristics of the population of the simulated cluster galaxies are presented in Paper~II \\citep{RPSL04} --- in this Paper~III of the series, we discuss our first results on the properties of the IC stars in two simulated clusters, with masses and temperatures similar to the Virgo and Coma clusters, respectively. Our results complement those of \\cite{M.04} and \\cite{W.04}, since we present for the IC stars also results on self-consistently calculated multiband surface brightness profiles, colours, abundances and detailed kinematics (Willman et al. 2004 also present a kinematics study, which we compare to in this paper).  In section 2 we briefly describe the code and the numerical simulations, in section 3 we present and discuss the results obtained, and finally section 4 constitutes our conclusions. ",
        "conclusions": "In this paper we have discussed the properties of the intracluster (IC) stellar populations. Our results are based on cosmological simulations of galaxy clusters including self-consistently metal-dependent atomic radiative cooling, star-formation, supernova driven galactic super-winds, non-instantaneous chemical evolution and the effects of a meta-galactic, redshift dependent UV field. In relation to modelling the properties of the IC stars this is an important step forward with respect to previous theoretical works on the subject which were based on purely N-body (collisionless) simulations, apart from the recent, interesting work by \\cite{M.04}.  The main results from our simulations of a Virgo-like (C1) and a (sub) Coma-like (C2) galaxy cluster regarding the IC stellar populations are as follows:  The intracluster (IC) stars are found to contribute 20-40\\% of the total cluster B-band luminosity at $z$=0 and to form at a mean redshift $\\bar{z}_f\\sim$3, somewhat larger than the mean formation redshift of the stars in cluster galaxies % which is about 2.5. This difference corresponds to a time span of about 0.5 Gyr. \\cite{M.04} find somewhat lower mean formation redshifts of 1.9 and 1.7, respectively --- the reason for this discrepancy between the two works remains to be identified.  We calculate UBVRIJHK surface brightness profiles of the IC star populations and find that the profile of the larger cluster matches the observed V-band profile of Abell 1413 ($T\\simeq$8 keV) quite well. For the Virgo-like cluster we fail, however, to match the flat profile between 200 and 400 kpc projected cluster-centric distance observationally (but indirectly) inferred for the Virgo cluster by \\cite{A.02} from planetary nebulae counts. We note though that the galaxy distribution in Virgo is also unusually flat \\citep{B.87}. The V-band surface brightness is found to reach V$_{lim}\\simeq$28.3 mag/arcsec$^2$ \\citep{F.02} in cluster C2 at $r\\sim$350 kpc and in cluster C1 at $r\\sim$250 kpc. The typical colour of the IC stellar population is B-R=1.4-1.5, comparable to the colour of sub-$L^*$ E and S0 galaxies. The mean Iron abundance of the IC stars is about solar in the central part of the cluster ($r\\sim$100 kpc) gently decreasing to about half solar at the virial radius. The IC stars are $\\alpha$-element enhanced with, e.g, [O/Fe] increasing slightly with $r$ and characterized by a typical [O/Fe]$\\sim$0.4 dex.  The IC stars are kinematically significantly colder than the cluster galaxies: The velocity dispersions of the IC stars are in the inner parts of the cluster ($r\\sim$100-500 kpc) only about half of those of the cluster galaxies increasing slightly to about 70\\% at $r$=1-2 Mpc. The typical projected velocity dispersion in the Virgo-like cluster at $R\\ga$50 kpc is 300-600 km/s depending on orientation and projected distance from the cluster center. Rotation is found to be dynamically insignificant for the IC stars. The velocity distributions of IC stars {\\it and} clusters galaxies are in one cluster highly radially anisotropic, in the other close to isotropic.  A test simulation of a $T\\simeq$1.1 keV group at higher numerical resolution indicates that the results presented are largely robust to resolution changes. Work is in progress to simulate also at higher resolution the two clusters discussed in this paper. Moreover, we are in the process of enlarging significantly our sample of galaxy clusters."
    },
    "0403/astro-ph0403557_arXiv.txt": {
        "abstract": "From a survey of the 3400~\\AA\\ region in the earliest O-type spectra, we have found that two of the four O2 giants observed in the Large Magellanic Cloud have the O~IV lines there stronger than the N~IV, while the other two have the opposite.  A Small Magellanic Cloud counterpart also has the N~IV stronger than O~IV.  Inspection of the blue spectra of these stars shows that the former pair have weaker N lines in all ionization states (III, IV, V) present as well as lines of C~IV $\\lambda$4658, while the latter three have stronger N lines and greater He/H.  Space ultraviolet observations of two of the N-strong stars show N~V wind profiles substantially stronger than those of C~IV, while in the N-weak stars the C~IV features are equal to or stronger than the N~V.  The N-strong stars are now reclassified as ON2~III(f*), newly defining that category.  These characteristics strongly suggest a larger fraction of processed material in the atmospheres of the ON2 stars, which we confirm by modeling the optical spectra.  In the context of current models, it is in turn implied that the ON2 stars are in a more advanced evolutionary state than the others, and/or that they had higher initial rotational velocities.  The recent formulation of the effects of rotation on massive stellar evolution introduces an additional fundamental parameter, which the CNO abundances are in principle able to constrain. We present some illustrative comparisons with current Geneva evolutionary models for rotating massive stars.  It is possible that these very hot, nitrogen-rich objects are products of homogeneous evolution.  Our results will provide motivation for further physical modeling of the atmospheres and evolutionary histories of the most massive hot stars. ",
        "introduction": "Anomalies in the relative intensities of CNO lines in OB spectra have been described for some time (Walborn 1976, 1988, 2003; and references therein).  Anticorrelations of N vs.\\ C, O and correlations with He/H in the expected sense have encouraged interpretations in terms of mixing of CNO-cycled material into the atmospheres and winds of massive stars.  An important point has been the recognition that the morphologically normal majority of OB supergiants actually have some degree of N and He enhancements already, while the OBC minority have normal (i.e., main-sequence) abundances and the OBN group have more extreme mixing (Maeder \\& Conti 1994; Smith \\& Howarth 1994).  The recent development of models for massive stellar evolution with rotation (Maeder \\& Meynet 2000) is an advance toward realism since massive stars often rotate rapidly, but another dimension is added to the phase space that must be explored in order to compare observations with theory.  Fortunately, a new observational diagnostic is added concurrently, namely the CNO abundances, since the degree of mixing also depends on the initial rotation.  Indeed, progress has already been made toward understanding apparent discrepancies among the HRD locations of some mid-O giants and supergiants in the Small Magellanic Cloud (SMC; Walborn et~al.\\ 2000), in terms of their CNO abundances and inferred initial rotations (Crowther et~al.\\ 2002; Hillier et~al.\\ 2003). Related results for main-sequence stars in the SMC giant H~II region NGC~346 are reported by Bouret et~al.\\ (2003).  Spectral classification of the hottest massive stars is relatively difficult, because of their scarcity and the paucity of reliable features in their optical spectra.  Recently, an increased sample of high-quality data has supported classification developments for them, including introduction of the new types O2 and O3.5 (Walborn et~al.\\ 2002).  In this system, reliance is placed on ratios of the selective emission lines (Walborn 2001) of N~IV and N~III, primarily on an empirical basis.  These new spectral types have yet to be modeled systematically and calibrated in terms of physical parameters, including confirmation that the O2 extension corresponds to higher effective temperatures as hypothesized.  The analysis of these spectra is also difficult, because of uncertainties in the parameters and the atmospheric physics.  E.g., the special behavior of the N~IV and N~III emission lines has to be reproduced and interpreted.  Currently, modeling and calibration developments are underway to address these issues; we shall present some relevant results here.  Certainly additional spectral features to constrain the classification and analysis of the hottest spectra would be welcome.  Drissen et~al.\\ (1995) pointed out certain little-observed N~IV and O~IV lines in the 3400~\\AA\\ region of early-O spectra, that warranted further investigation in a larger sample including longer-wavelength classification standards. We have now accomplished such a survey, which will be fully reported by Morrell et~al.\\ (in preparation).  Here we describe our unanticipated discovery of the first CNO anomalies in optical O2 spectra, which are relevant to both the classification and modeling of these extreme objects. ",
        "conclusions": "Curiously, the luminosities, and hence the masses, estimated by Walborn et~al.\\ (2002) for the O2~III(f*) stars HDE~269810 (150~$M_{\\odot}$, average of two methods) and Sk~$-$68$^{\\circ}$~137 (100~$M_{\\odot}$) are considerably higher than those for the ON2~III(f*) stars LH64$-$16 (62~$M_{\\odot}$), LH10$-$3061 (81~$M_{\\odot}$), and NGC~346$-$3 (82~$M_{\\odot}$).  The opposite relationship might have been expected, since both mass-loss rates and mixing increase with mass (Meynet \\& Maeder 2000).  On the other hand, perhaps the longer evolutionary timescales at the lower masses favor the observability of mixing effects.  Until recently, these abundance results would have been interpreted simply in terms of more advanced evolutionary ages for the ON2 than the O2 objects, which remains a possibility.  (Of course, mass transfer in a binary system is an alternative mechanism likely responsible for at least some members of the ON class---see Bolton \\& Rogers 1978; Levato et~al.\\ 1988; Walborn \\& Howarth 2000; Lennon 2003).  In the context of evolutionary models with rotation for massive stars (Maeder \\& Meynet 2000; Meynet \\& Maeder 2000), however, the degree of mixing of processed material to the surface at a given evolutionary age is a function of the initial rotational velocity. Indeed, CNO abundances in massive stars now become critical diagnostics of their initial parameters and evolutionary histories.  For a given metallicity, the observed $T_{\\rm eff}$, $L$, and CNO abundances in principle enable derivation of the corresponding initial mass, current age, and initial rotational velocity.  Complete grids of evolutionary tracks for the most massive stars, at Magellanic Cloud metallicities and with a range of rotational velocities, are currently under development in Geneva. Meanwhile, we display illustrative comparisons in Figures~6--8, courtesy of G.~Meynet and A.~Maeder.  Figure~6 is a theoretical HR Diagram in which the locations of the stars discussed here are compared with Geneva evolutionary tracks for initial masses of 60 and 120~$M_{\\odot}$, at both LMC and SMC metallicities (mass fractions Z~=~0.008 and 0.004, respectively), and with an initial equatorial rotational velocity of 300~km~s$^{-1}$.  All of the stars lie near the zero-age main sequence (ZAMS), which is consistent with an early evolutionary stage for the two objects with normal spectra; the present spectroscopic analysis again indicates higher masses for them than for the ON2 objects. However, such a location is surprising for the latter, highly mixed stars.  Their initial masses still place them above the Humphreys-Davidson Limit, so they cannot be interpreted as post-red supergiants evolving blueward.  A more relevant possibility may be that of ``homogeneous'' evolution, along tracks that rise and curve blueward from the ZAMS in the HRD, which is predicted in the cases of strong mixing and/or very high initial rotational velocities (Langer \\& El~Eid 1986; Maeder 1987; Langer 1992; Meynet \\& Maeder 2000).  The 60~$M_{\\odot}$ homogeneous track at solar metallicity (Z~=~0.020) with an initial rotational velocity of 400~km~s$^{-1}$ from the last reference is also reproduced in Fig.~6; it trends near the locations of the ON2~III(f*) points.  Indeed, Bouret et al. (2003) have concluded that such a track is required to understand NGC~346-3 as coeval with the other early O stars in the cluster.  Figures~7 and 8 compare our derived N/C and N/O ratios, respectively, relative to their corresponding interstellar media, taken at face value from Table~1, with the predictions of the same Geneva models.  The N/C ratios in the ON2~III(f*) stars are the largest reported to date for near-main-sequence objects and are substantially larger than the standard model predictions, suggesting even more rapid mixing and perhaps providing additional support for homogeneous evolution.  In contrast, the N/O ratios are compatible with the standard model predictions and indicate early operation of the ON cycle, as expected in only the most massive stars.  On the other hand, the even larger N enhancements derived for later-type O supergiants in the Magellanic Clouds by Crowther et~al.\\ (2002) and Hillier et~al.\\ (2003) may imply that the ON2~III(f*) stars are simply predecessors of those objects along redward evolutionary tracks.  Again, much stronger mixing effects near the main sequence than predicted by current models, at least for N/C, are then required by our observations. Further analysis of these phenomena, with improved data and relative to the latest models for rotating massive stars, will clarify these issues and provide an improved understanding of their post-main sequence evolution."
    },
    "0403/hep-th0403026_arXiv.txt": {
        "abstract": "For a 4--dimensional spatially-flat Friedmann-Robertson-Walker universe with a scalar field $\\phi(x)$, potential $V(\\phi)$ and constant equation of state $w=p/\\rho$, we show that an expanding solution characterized by $\\epsilon=3(1+w)/2$ produces the same scalar perturbations as a contracting solution with $\\widehat{\\epsilon}=1/\\epsilon$.  The same symmetry applies to both the dominant and subdominant scalar perturbation modes.  This result admits a simple physical interpretation and generalizes to $d$ spacetime dimensions if we define $\\epsilon\\equiv[(2d-5)+(d-1)w]/(d-2)$. ",
        "introduction": "\\label{intro}  In inflationary cosmology \\cite{inflation}, a nearly scale-invariant spectrum of primordial density perturbations is produced as comoving scales leave the Hubble horizon during an early burst of accelerated expansion \\cite{inflation_perts}. In cyclic cosmology \\cite{cyclic_ekpyrotic}, the \\emph{same} perturbation spectrum is produced as comoving scales leave the Hubble horizon during a period of slow decelerated contraction \\cite{cyclic_perts,TTS}.  This agreement between two physically dissimilar models is unexpected, but not coincidental. As we shall show, the relationship between inflation and the cyclic model may be viewed as a special case of a surprisingly simple and general duality between expanding and contracting cosmologies.  A Friedmann-Robertson-Walker (FRW) universe with a single scalar field $\\phi$ and potential $V(\\phi)$ is a simple yet important system.  In particular, it is the canonical 4d effective theory used to model the production of density perturbations in both inflationary and cyclic cosmology.  Recent results hint at a connection between two apparently--unrelated regimes of this model: (\\emph{i}) expanding, in the $w\\rightarrow -1$ limit, and (\\emph{ii}) contracting, in the $w\\rightarrow\\infty$ limit.  Here $w\\equiv p/\\rho$ denotes the ratio of pressure to energy density. Long-wavelength scale-invariant density perturbations are produced in the limits (\\emph{i}) and (\\emph{ii}) \\cite{PHZ_conditions}, and the small deviations from scale invariance near these two limits are related by the simple substitution $\\epsilon\\rightarrow 1/\\epsilon$, where $\\epsilon\\equiv 3(1+w)/2$ \\cite{spectral_tilt}.  The relationship noted in \\cite{PHZ_conditions,spectral_tilt} between expanding $w\\approx -1$ models and contracting $w\\gg 1$ models, turns out to be a special case of a general and exact duality relating expanding and contracting models with identical perturbation spectra.  In this paper, we derive this duality, focusing on the case where $\\epsilon$ (or $w$) is time-independent. In a companion paper \\cite{dual2}, we generalize the discussion to the case where $\\epsilon$ is time-varying.  When $\\epsilon$ is constant, or varies sufficiently slowly, the duality is simple: an expanding universe characterized by $\\epsilon$ produces \\emph{exactly} the same scalar perturbations as a contracting universe characterized by $\\widehat{\\epsilon}=1/\\epsilon$.  This duality applies in arbitrary spacetime dimension (not just 3+1); it applies for all $w$ (not just the $w\\rightarrow -1$ and $w\\rightarrow\\infty$ limits discussed in \\cite{PHZ_conditions,spectral_tilt}); it applies to all wavelengths (not just the long-wavelength limit); it applies to \\emph{both} the dominant scalar perturbation mode \\emph{and} a subdominant remainder (which are related to the growing and decaying modes, respectively).  This duality is of general theoretical interest since it provides a new relationship between expanding and contracting universes, and exposes an unexpected symmetry of cosmological perturbation theory. It is also relevant to cosmological models, like the ekpyrotic and cyclic \\cite{cyclic_ekpyrotic} scenarios, in which perturbations produced during a period of contraction are proposed to propagate through a bounce into a subsequent expanding phase. These models require that the growing-mode long-wavelength perturbation spectrum is preserved across the bounce.  This has been a controversial matter.  At first, some authors argued that growing-mode perturbations produced in a contracting phase must match to pure decaying-mode perturbations as one follows them across a bounce into an expanding phase \\cite{contrascale}.  At heart, this conclusion followed from requiring that the bounce corresponds to a comoving or constant-energy-density slice.  However, recent five dimensional calculations \\cite{TTS,newbrand} indicate that comoving or constant-energy-density slices are inappropriate for matching, since the bounce event (represented in five dimensions by a brane collision) is not synchronous in these slices.  Matching on collision-synchronous slices results in the propagation of growing-mode perturbations across a bounce, with no change in the shape of the long-wavelength spectrum.  Other aspects of the cyclic and ekpyrotic models have been criticized \\cite{kallosh} (see \\cite{ansc} for replies), and some have argued that a bounce is impossible altogether \\cite{contrabounce}. While the consistency of a bounce remains to be proven, recent work has shown that the traditional hazard of chaotic mixmaster behavior is strongly suppressed in the contracting phase of the cyclic model \\cite{chaos}.  The metric perturbations exhibit ultralocal behavior in which anisotropies remain small, right up to a few Planck times before the bounce.  In this situation, causality suggests that the bounce should not disturb correlations on macroscopic scales over which there can be no communication in this finite time interval.  Under these conditions, works by several groups \\cite{TTS,pro,newbrand} suggest that perturbations generated during the contracting phase may pass into the expanding phase.  The subject continues to be an area of active research. If the latter suggestions are made rigorous, it would give added significance to the results presented herein.  Other dualities have been identified in the literature \\cite{Wands, Starobinsky, Brustein, scalefacdual, stringdual} that relate cosmological solutions and perturbations.  See Section \\ref{other} for a comparison.  The duality presented here has the distinctive property that it relates two solutions that are stable under perturbations.  Hence, both solutions can plausibly play a role in realistic cosmological models.  The layout of the paper is as follows.  In section \\ref{background}, we introduce the background (unperturbed) model: a spatially-flat FRW model with scalar field $\\phi$, potential $V(\\phi)$, and constant $\\epsilon$.  In section \\ref{scalars}, we briefly review scalar perturbation theory, before deriving exact solutions for Mukhanov's $u$ and $v$ variables.  We note that $u$ is invariant under $\\epsilon\\rightarrow 1/\\epsilon$.  This invariance is independent of our vacuum choice for the fluctuations, and provides our first glimpse of the duality.  In section \\ref{grow_decay} we use $u$ and $v$ to separate scalar perturbations into pieces that are dominant and subdominant at long wavelengths, and show that each piece is {\\it both independently invariant} under $\\epsilon \\to 1/\\epsilon$. We show how the dominant and subdominant pieces relate to the scalar perturbation growing and decaying modes.  In section \\ref{tensors} we consider tensor perturbations. In section \\ref{other}, we contrast our duality with other kinds of cosmological dualities that have been studied in the literature \\cite{Wands,Starobinsky,Brustein,scalefacdual,stringdual}. In section \\ref{discuss}, we interpret our duality as a relation between the scale factor and the Hubble parameter, discuss its observational significance, and mention some open questions.  In an appendix, we generalize our results to $d$ spacetime dimensions. ",
        "conclusions": "\\label{discuss}  Beyond its inherent theoretical interest, our duality may be observationally relevant if (as discussed in section \\ref{intro}) long-wavelength correlations produced during a contracting phase successfully propagate into a subsequent expanding phase.  This suggests a fundamental degeneracy: an ideal measurement of the ``primordial\" scalar perturbation spectrum may be unable to determine whether the perturbations were generated by an expanding phase or by its contracting dual.  Luckily, as shown in section \\ref{tensors}, tensor perturbations break this degeneracy: an expanding model produces a much redder tensor spectrum than its contracting dual.  In particular, a detection of tensors in the cosmic microwave background would indicate that these perturbations were generated in an expanding phase, since the dual contracting phase would produce an undetectably small tensor spectrum on these cosmological length scales \\cite{gravity_waves}.  Using the background solutions (\\ref{background_soln}), we may think of the duality as relating two different scale factors, $a(\\tau)$ and $\\widehat{a}(\\tau)$, or two different scalar potentials, $V(\\phi)$ and $\\widehat{V}(\\phi)$.  Alternatively, recall that $a\\propto |t|^{1/\\epsilon}$ and $H^{-1}\\equiv a^{2}/a'\\propto t$, where $t$ is the proper time of a comoving observer and $H^{-1}$ is the ``Hubble scale.\"  Thus, two dual universes are related by \\begin{equation} \\label{dln/dln_correspondence} \\textrm{d\\,ln}H/\\textrm{d\\,ln}a =\\textrm{d\\,ln}\\widehat{a}/\\textrm{d\\,ln}\\widehat{H}. \\end{equation}  So the duality effectively swaps the scale factor with the Hubble scale, and simultaneously swaps expansion and contraction.  For example, in the $\\epsilon\\rightarrow 0$ ($w\\rightarrow -1$) limit, the scale factor grows rapidly while the Hubble length grows slowly; whereas in the $\\epsilon\\rightarrow\\infty$ ($w\\rightarrow\\infty$) limit, the Hubble length shrinks rapidly while the scale factor shrinks slowly.  Expanding models in which modes exit the horizon most rapidly ($w\\rightarrow -1$) and most slowly ($w\\rightarrow -1/3$), are associated with contracting models in which modes exit most rapidly ($w\\rightarrow\\infty$) and most slowly ($w\\rightarrow -1/3$), respectively.  Finally, let us mention several open questions.  We have studied the duality using a simple model---spatially-flat FRW spacetime with a single scalar field $\\phi$; what happens in more complicated models? We have used linear perturbation theory; what happens in the nonlinear regime?  We have restricted ourselves to time-independent $\\epsilon$; what if we allow time-varying $\\epsilon$?  This final question is treated in a companion paper \\cite{dual2}."
    },
    "0403/astro-ph0403671_arXiv.txt": {
        "abstract": "In this paper, focusing on the case of single scalar field, we discuss various expanding and contracting phases generating primordial perturbations, and study the relation between the primordial perturbation spectrum from these phases and the parameter $w$ of state equation in details. Furthermore, we offer an interesting classification for the primordial perturbation spectrum from various phases, which may have important implications for building an early universe scenario embedded in possible high energy theories. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403392_arXiv.txt": {
        "abstract": "s#1#2#3{{ \\centering{\\begin{minipage}{4.5in}\\baselineskip=10pt\\footnotesize \\parindent=0pt #1\\par \\parindent=15pt #2\\par \\parindent=15pt #3 \\end{minipage}}\\par}}  \\def\\keywords#1{{ \\centering{\\begin{minipage}{4.5in}\\baselineskip=10pt\\footnotesize {\\footnotesize\\it Keywords}\\/: #1 \\end{minipage}}\\par}}  \\newcommand{\\bibit}{\\nineit} \\newcommand{\\bibbf}{\\ninebf} \\renewenvironment{thebibliography}[1] {\\frenchspacing \\ninerm\\baselineskip=11pt \\begin{list}{\\arabic{enumi}.} {\\usecounter{enumi}\\setlength{\\parsep}{0pt} \\setlength{\\leftmargin 12.7pt}{\\rightmargin 0pt} % \\setlength{\\itemsep}{0pt} \\settowidth {\\labelwidth}{#1.}\\sloppy}}{\\end{list}}  \\newcounter{itemlistc} \\newcounter{romanlistc} \\newcounter{alphlistc} \\newcounter{arabiclistc} \\newenvironment{itemlist} {\\setcounter{itemlistc}{0} \\begin{list}{$\\bullet$} {\\usecounter{itemlistc} \\setlength{\\parsep}{0pt} \\setlength{\\itemsep}{0pt}}}{\\end{list}}  \\newenvironment{romanlist} {\\setcounter{romanlistc}{0} \\begin{list}{$($\\roman{romanlistc}$)$} {\\usecounter{romanlistc} \\setlength{\\parsep}{0pt} \\setlength{\\itemsep}{0pt}}}{\\end{list}}  \\newenvironment{alphlist} {\\setcounter{alphlistc}{0} \\begin{list}{$($\\alph{alphlistc}$)$} {\\usecounter{alphlistc} \\setlength{\\parsep}{0pt} \\setlength{\\itemsep}{0pt}}}{\\end{list}}  \\newenvironment{arabiclist} {\\setcounter{arabiclistc}{0} \\begin{list}{\\arabic{arabiclistc}} {\\usecounter{arabiclistc} \\setlength{\\parsep}{0pt} \\setlength{\\itemsep}{0pt}}}{\\end{list}}  \\newcommand{\\fcaption}[1]{ \\refstepcounter{figure} \\setbox\\@tempboxa = \\hbox{\\footnotesize Fig.~\\thefigure. #1} \\ifdim \\wd\\@tempboxa > 5in {\\begin{center} \\parbox{5in}{\\footnotesize\\smalllineskip Fig.~\\thefigure. #1} \\end{center}} \\else {\\begin{center} {\\footnotesize Fig.~\\thefigure. #1} \\end{center}} \\fi}  \\newcommand{\\tcaption}[1]{ \\refstepcounter{table} \\setbox\\@tempboxa = \\hbox{\\footnotesize Table~\\thetable. #1} \\ifdim \\wd\\@tempboxa > 5in {\\begin{center} \\parbox{5in}{\\footnotesize\\smalllineskip Table~\\thetable. #1} \\end{center}} \\else {\\begin{center} {\\footnotesize Table~\\thetable. #1} \\end{center}} \\fi}  \\def\\@citex[#1]#2{\\if@filesw\\immediate\\write\\@auxout {\\string\\citation{#2}}\\fi \\def\\@citea{}\\@cite{\\@for\\@citeb:=#2\\do {\\@citea\\def\\@citea{,}\\@ifundefined {b@\\@citeb}{{\\bf ?}\\@warning {Citation `\\@citeb' on page \\thepage \\space undefined}} {\\csname b@\\@citeb\\endcsname}}}{#1}}  \\newif\\if@cghi \\newcommand{\\citeup}{\\cite}  \\def\\pmb#1{\\setbox0=\\hbox{#1} \\kern-.025em\\copy0\\kern-\\wd0 \\kern.05em\\copy0\\kern-\\wd0 \\kern-.025em\\raise.0433em\\box0} \\def\\mbi#1{{\\pmb{\\mbox{\\scriptsize ${#1}$}}}} \\def\\mbr#1{{\\pmb{\\mbox{\\scriptsize{#1}}}}}  \\def\\fnm#1{$^{\\mbox{\\scriptsize #1}}$} \\def\\fnt#1#2{\\footnotetext{\\kern-.3em {$^{\\mbox{\\scriptsize #1}}$}{#2}}}  \\def\\fpage#1{\\begingroup \\voffset=.3in \\thispagestyle{empty}\\begin{table}[b]\\centerline{\\footnotesize #1} \\end{table}\\endgroup}  \\def\\runninghead#1#2{\\pagestyle{myheadings} \\markboth{{\\protect\\footnotesize\\it{\\quad #1}}\\hfill} {\\hfill{\\protect\\footnotesize\\it{#2\\quad}}}} \\headsep=15pt  \\font\\tenrm=cmr10 \\font\\tenit=cmti10 \\font\\tenbf=cmbx10 \\font\\bfit=cmbxti10 at 10pt \\font\\ninerm=cmr9 \\font\\nineit=cmti9 \\font\\ninebf=cmbx9 \\font\\eightrm=cmr8 \\font\\eightit=cmti8 \\font\\eightbf=cmbx8 \\font\\sevenrm=cmr7 \\font\\fiverm=cmr5  \\newtheorem{theorem}{\\indent Theorem}  \\newtheorem{lemma}{Lemma}  \\newtheorem{definition}{Definition} \\newtheorem{corollary}{Corollary}  \\newcommand{\\proof}[1]{{\\tenbf Proof.} #1 $\\Box$.}     \\textwidth=5truein \\textheight=7.8truein  \\def\\qed{\\hbox{${\\vcenter{\\vbox{                        % \\hrule height 0.4pt\\hbox{\\vrule width 0.4pt height 6pt \\kern5pt\\vrule width 0.4pt}\\hrule height 0.4pt}}}$}}  \\renewcommand{\\thefootnote}{\\fnsymbol{footnote}}        %  \\begin{document}  \\def\\PsfigVersion{1.10} \\def\\setDriver{\\DvipsDriver} % \\ifx\\undefined\\psfig\\else\\endinput\\fi  \\let\\LaTeXAtSign=\\@ \\let\\@=\\relax \\edef\\psfigRestoreAt{\\catcode`\\@=\\number\\catcode`@\\relax} \\catcode`\\@=11\\relax \\newwrite\\@unused \\def\\ps@typeout#1{{\\let\\protect\\string\\immediate\\write\\@unused{#1}}}  \\def\\DvipsDriver{ \\ps@typeout{psfig/tex \\PsfigVersion -dvips} \\def\\PsfigSpecials{\\DvipsSpecials} \t\\def\\ps@dir{/} \\def\\ps@predir{} } \\def\\OzTeXDriver{ \\ps@typeout{psfig/tex \\PsfigVersion -oztex} \\def\\PsfigSpecials{\\OzTeXSpecials} \\def\\ps@dir{:} \\def\\ps@predir{:} \\catcode`\\^^J=5 }   \\def\\figurepath{./:} \\def\\psfigurepath#1{\\edef\\figurepath{#1:}}  \\def\\DoPaths#1{\\expandafter\\EachPath#1\\stoplist} \\def\\leer{} \\def\\EachPath#1:#2\\stoplist{% \\ExistsFile{#1}{\\SearchedFile} \\ifx#2\\leer \\else \\expandafter\\EachPath#2\\stoplist \\fi} \\def\\ps@dir{/} \\def\\ExistsFile#1#2{% \\openin1=\\ps@predir#1\\ps@dir#2 \\ifeof1 \\closein1 \\else \\closein1 \\ifx\\ps@founddir\\leer \\edef\\ps@founddir{#1} \\fi \\fi} \\def\\get@dir#1{% \\def\\ps@founddir{} \\def\\SearchedFile{#1} \\DoPaths\\figurepath }   \\def\\@nnil{\\@nil} \\def\\@empty{} \\def\\@psdonoop#1\\@@#2#3{} \\def\\@psdo#1:=#2\\do#3{\\edef\\@psdotmp{#2}\\ifx\\@psdotmp\\@empty \\else \\expandafter\\@psdoloop#2,\\@nil,\\@nil\\@@#1{#3}\\fi} \\def\\@psdoloop#1,#2,#3\\@@#4#5{\\def#4{#1}\\ifx #4\\@nnil \\else #5\\def#4{#2}\\ifx #4\\@nnil \\else#5\\@ipsdoloop #3\\@@#4{#5}\\fi\\fi} \\def\\@ipsdoloop#1,#2\\@@#3#4{\\def#3{#1}\\ifx #3\\@nnil \\let\\@nextwhile=\\@psdonoop \\else #4\\relax\\let\\@nextwhile=\\@ipsdoloop\\fi\\@nextwhile#2\\@@#3{#4}} \\def\\@tpsdo#1:=#2\\do#3{\\xdef\\@psdotmp{#2}\\ifx\\@psdotmp\\@empty \\else \\@tpsdoloop#2\\@nil\\@nil\\@@#1{#3}\\fi} \\def\\@tpsdoloop#1#2\\@@#3#4{\\def#3{#1}\\ifx #3\\@nnil \\let\\@nextwhile=\\@psdonoop \\else #4\\relax\\let\\@nextwhile=\\@tpsdoloop\\fi\\@nextwhile#2\\@@#3{#4}} \\ifx\\undefined\\fbox \\newdimen\\fboxrule \\newdimen\\fboxsep \\newdimen\\ps@tempdima \\newbox\\ps@tempboxa \\fboxsep = 3pt \\fboxrule = .4pt \\long\\def\\fbox#1{\\leavevmode\\setbox\\ps@tempboxa\\hbox{#1}\\ps@tempdima\\fboxrule \\advance\\ps@tempdima \\fboxsep \\advance\\ps@tempdima \\dp\\ps@tempboxa \\hbox{\\lower \\ps@tempdima\\hbox {\\vbox{\\hrule height \\fboxrule \\hbox{\\vrule width \\fboxrule \\hskip\\fboxsep \\vbox{\\vskip\\fboxsep \\box\\ps@tempboxa\\vskip\\fboxsep}\\hskip \\fboxsep\\vrule width \\fboxrule} \\hrule height \\fboxrule}}}} \\fi \\newread\\ps@stream \\newif\\ifnot@eof       % \\newif\\if@noisy        % \\newif\\if@atend        % \\newif\\if@psfile       % {\\catcode`\\%=12\\global\\gdef\\epsf@start{% \\def\\epsf@PS{PS} \\def\\epsf@getbb#1{% \\openin\\ps@stream=\\ps@predir#1 \\ifeof\\ps@stream\\ps@typeout{Error, File #1 not found}\\else {\\not@eoftrue \\chardef\\other=12 \\def\\do##1{\\catcode`##1=\\other}\\dospecials \\catcode`\\ =10 \\loop \\if@psfile \\read\\ps@stream to \\epsf@fileline \\else{ \\obeyspaces \\read\\ps@stream to \\epsf@tmp\\global\\let\\epsf@fileline\\epsf@tmp} \\fi \\ifeof\\ps@stream\\not@eoffalse\\else \\if@psfile\\else \\expandafter\\epsf@test\\epsf@fileline:. \\\\% \\fi \\expandafter\\epsf@aux\\epsf@fileline:. \\\\% \\fi \\ifnot@eof\\repeat }\\closein\\ps@stream\\fi}% \\long\\def\\epsf@test#1#2#3:#4\\\\{\\def\\epsf@testit{#1#2} \\ifx\\epsf@testit\\epsf@start\\else \\ps@typeout{Warning! File does not start with `\\epsf@start'.  It may not be a PostScript file.} \\fi \\@psfiletrue} % {\\catcode`\\%=12\\global\\let\\epsf@percent=% \\long\\def\\epsf@aux#1#2:#3\\\\{\\ifx#1\\epsf@percent \\def\\epsf@testit{#2}\\ifx\\epsf@testit\\epsf@bblit \\@atendfalse \\epsf@atend #3 . \\\\% \\if@atend \\if@verbose{ \\ps@typeout{psfig: found `(atend)'; continuing search} }\\fi \\else \\epsf@grab #3 . . . \\\\% \\not@eoffalse \\global\\no@bbfalse \\fi \\fi\\fi}% \\def\\epsf@grab #1 #2 #3 #4 #5\\\\{% \\global\\def\\epsf@llx{#1}\\ifx\\epsf@llx\\empty \\epsf@grab #2 #3 #4 #5 .\\\\\\else \\global\\def\\epsf@lly{#2}% \\global\\def\\epsf@urx{#3}\\global\\def\\epsf@ury{#4}\\fi}% \\def\\epsf@atendlit{(atend)} \\def\\epsf@atend #1 #2 #3\\\\{% \\def\\epsf@tmp{#1}\\ifx\\epsf@tmp\\empty \\epsf@atend #2 #3 .\\\\\\else \\ifx\\epsf@tmp\\epsf@atendlit\\@atendtrue\\fi\\fi}    \\chardef\\psletter = 11 % \\chardef\\other = 12  \\newif \\ifdebug % \\newif\\ifc@mpute % \\c@mputetrue %  \\let\\then = \\relax \\def\\r@dian{pt } \\let\\r@dians = \\r@dian \\let\\dimensionless@nit = \\r@dian \\let\\dimensionless@nits = \\dimensionless@nit \\def\\internal@nit{sp } \\let\\internal@nits = \\internal@nit \\newif\\ifstillc@nverging \\def \\Mess@ge #1{\\ifdebug \\then \\message {#1} \\fi}  { % \\catcode `\\@ = \\psletter \\gdef \\nodimen {\\expandafter \\n@dimen \\the \\dimen} \\gdef \\term #1 #2 #3% {\\edef \\t@ {\\the #1}% \\edef \\t@@ {\\expandafter \\n@dimen \\the #2\\r@dian}% % \\t@rm {\\t@} {\\t@@} {#3}% } \\gdef \\t@rm #1 #2 #3% {{% \\count 0 = 0 \\dimen 0 = 1 \\dimensionless@nit \\dimen 2 = #2\\relax \\Mess@ge {Calculating term #1 of \\nodimen 2}% \\loop \\ifnum\t\\count 0 < #1 \\then\t\\advance \\count 0 by 1 \\Mess@ge {Iteration \\the \\count 0 \\space}% \\Multiply \\dimen 0 by {\\dimen 2}% \\Mess@ge {After multiplication, term = \\nodimen 0}% \\Divide \\dimen 0 by {\\count 0}% \\Mess@ge {After division, term = \\nodimen 0}% \\repeat \\Mess@ge {Final value for term #1 of \\nodimen 2 \\space is \\nodimen 0}% \\xdef \\Term {#3 = \\nodimen 0 \\r@dians}% \\aftergroup \\Term }} \\catcode `\\p = \\other \\catcode `\\t = \\other \\gdef \\n@dimen #1pt{#1} % }  \\def \\Divide #1by #2{\\divide #1 by #2} %  \\def \\Multiply #1by #2% {{% \\count 0 = #1\\relax \\count 2 = #2\\relax \\count 4 = 65536 \\Mess@ge {Before scaling, count 0 = \\the \\count 0 \\space and count 2 = \\the \\count 2}% \\ifnum\t\\count 0 > 32767 % \\then\t\\divide \\count 0 by 4 \\divide \\count 4 by 4 \\else\t\\ifnum\t\\count 0 < -32767 \\then\t\\divide \\count 0 by 4 \\divide \\count 4 by 4 \\else \\fi \\fi \\ifnum\t\\count 2 > 32767 % \\then\t\\divide \\count 2 by 4 \\divide \\count 4 by 4 \\else\t\\ifnum\t\\count 2 < -32767 \\then\t\\divide \\count 2 by 4 \\divide \\count 4 by 4 \\else \\fi \\fi \\multiply \\count 0 by \\count 2 \\divide \\count 0 by \\count 4 \\xdef \\product {#1 = \\the \\count 0 \\internal@nits}% \\aftergroup \\product }}  \\def\\r@duce{\\ifdim\\dimen0 > 90\\r@dian \\then   % \\multiply\\dimen0 by -1 \\advance\\dimen0 by 180\\r@dian \\r@duce \\else \\ifdim\\dimen0 < -90\\r@dian \\then  % \\advance\\dimen0 by 360\\r@dian \\r@duce \\fi \\fi}  \\def\\Sine#1% {{% \\dimen 0 = #1 \\r@dian \\r@duce \\ifdim\\dimen0 = -90\\r@dian \\then \\dimen4 = -1\\r@dian \\c@mputefalse \\fi \\ifdim\\dimen0 = 90\\r@dian \\then \\dimen4 = 1\\r@dian \\c@mputefalse \\fi \\ifdim\\dimen0 = 0\\r@dian \\then \\dimen4 = 0\\r@dian \\c@mputefalse \\fi \\ifc@mpute \\then \\divide\\dimen0 by 180 \\dimen0=3.141592654\\dimen0 \\dimen 2 = 3.1415926535897963\\r@dian % \\divide\\dimen 2 by 2 % \\Mess@ge {Sin: calculating Sin of \\nodimen 0}% \\count 0 = 1 % \\dimen 2 = 1 \\r@dian % \\dimen 4 = 0 \\r@dian % \\loop \\ifnum\t\\dimen 2 = 0 % \\then\t\\stillc@nvergingfalse \\else\t\\stillc@nvergingtrue \\fi \\ifstillc@nverging % \\then\t\\term {\\count 0} {\\dimen 0} {\\dimen 2}% \\advance \\count 0 by 2 \\count 2 = \\count 0 \\divide \\count 2 by 2 \\ifodd\t\\count 2 % \\then\t\\advance \\dimen 4 by \\dimen 2 \\else\t\\advance \\dimen 4 by -\\dimen 2 \\fi \\repeat \\fi \\xdef \\sine {\\nodimen 4}% }}  \\def\\Cosine#1{\\ifx\\sine\\UnDefined\\edef\\Savesine{\\relax}\\else \\edef\\Savesine{\\sine}\\fi {\\dimen0=#1\\r@dian\\advance\\dimen0 by 90\\r@dian \\Sine{\\nodimen 0} \\xdef\\cosine{\\sine} \\xdef\\sine{\\Savesine}}}  \\def\\psdraft{ \\def\\@psdraft{0} } \\def\\psfull{ \\def\\@psdraft{100} }  \\psfull  \\newif\\if@scalefirst \\def\\psscalefirst{\\@scalefirsttrue} \\def\\psrotatefirst{\\@scalefirstfalse} \\psrotatefirst  \\newif\\if@draftbox \\def\\psnodraftbox{ \\@draftboxfalse } \\def\\psdraftbox{ \\@draftboxtrue } \\@draftboxtrue  \\newif\\if@prologfile \\newif\\if@postlogfile \\def\\pssilent{ \\@noisyfalse } \\def\\psnoisy{ \\@noisytrue } \\psnoisy \\newif\\if@bbllx \\newif\\if@bblly \\newif\\if@bburx \\newif\\if@bbury \\newif\\if@height \\newif\\if@width \\newif\\if@rheight \\newif\\if@rwidth \\newif\\if@angle \\newif\\if@clip \\newif\\if@verbose \\def\\@p@@sclip#1{\\@cliptrue} \\newif\\if@decmpr \\def\\@p@@sfigure#1{\\def\\@p@sfile{null}\\def\\@p@sbbfile{null}\\@decmprfalse \\openin1=\\ps@predir#1 \\ifeof1 \\closein1 \\get@dir{#1} \\ifx\\ps@founddir\\leer \\openin1=\\ps@predir#1.bb \\ifeof1 \\closein1 \\get@dir{#1.bb} \\ifx\\ps@founddir\\leer \\ps@typeout{Can't find #1 in \\figurepath} \\else \\@decmprtrue \\def\\@p@sfile{\\ps@founddir\\ps@dir#1} \\def\\@p@sbbfile{\\ps@founddir\\ps@dir#1.bb} \\fi \\else \\closein1 \\@decmprtrue \\def\\@p@sfile{#1} \\def\\@p@sbbfile{#1.bb} \\fi \\else \\def\\@p@sfile{\\ps@founddir\\ps@dir#1} \\def\\@p@sbbfile{\\ps@founddir\\ps@dir#1} \\fi \\else \\closein1 \\def\\@p@sfile{#1} \\def\\@p@sbbfile{#1} \\fi } \\def\\@p@@sfile#1{\\@p@@sfigure{#1}} \\def\\@p@@sbbllx#1{ \\@bbllxtrue \\dimen100=#1 \\edef\\@p@sbbllx{\\number\\dimen100} } \\def\\@p@@sbblly#1{ \\@bbllytrue \\dimen100=#1 \\edef\\@p@sbblly{\\number\\dimen100} } \\def\\@p@@sbburx#1{ \\@bburxtrue \\dimen100=#1 \\edef\\@p@sbburx{\\number\\dimen100} } \\def\\@p@@sbbury#1{ \\@bburytrue \\dimen100=#1 \\edef\\@p@sbbury{\\number\\dimen100} } \\def\\@p@@sheight#1{ \\@heighttrue \\dimen100=#1 \\edef\\@p@sheight{\\number\\dimen100} } \\def\\@p@@swidth#1{ \\@widthtrue \\dimen100=#1 \\edef\\@p@swidth{\\number\\dimen100} } \\def\\@p@@srheight#1{ \\@rheighttrue \\dimen100=#1 \\edef\\@p@srheight{\\number\\dimen100} } \\def\\@p@@srwidth#1{ \\@rwidthtrue \\dimen100=#1 \\edef\\@p@srwidth{\\number\\dimen100} } \\def\\@p@@sangle#1{ \\@angletrue \\edef\\@p@sangle{#1} % } \\def\\@p@@ssilent#1{ \\@verbosefalse } \\def\\@p@@sprolog#1{\\@prologfiletrue\\def\\@prologfileval{#1}} \\def\\@p@@spostlog#1{\\@postlogfiletrue\\def\\@postlogfileval{#1}} \\def\\@cs@name#1{\\csname #1\\endcsname} \\def\\@setparms#1=#2,{\\@cs@name{@p@@s#1}{#2}} \\def\\ps@init@parms{ \\@bbllxfalse \\@bbllyfalse \\@bburxfalse \\@bburyfalse \\@heightfalse \\@widthfalse \\@rheightfalse \\@rwidthfalse \\def\\@p@sbbllx{}\\def\\@p@sbblly{} \\def\\@p@sbburx{}\\def\\@p@sbbury{} \\def\\@p@sheight{}\\def\\@p@swidth{} \\def\\@p@srheight{}\\def\\@p@srwidth{} \\def\\@p@sangle{0} \\def\\@p@sfile{} \\def\\@p@sbbfile{} \\def\\@p@scost{10} \\def\\@sc{} \\@prologfilefalse \\@postlogfilefalse \\@clipfalse \\if@noisy \\@verbosetrue \\else \\@verbosefalse \\fi } \\def\\parse@ps@parms#1{ \\@psdo\\@psfiga:=#1\\do {\\expandafter\\@setparms\\@psfiga,}} \\newif\\ifno@bb \\def\\bb@missing{ \\if@verbose{ \\ps@typeout{psfig: searching \\@p@sbbfile \\space  for bounding box} }\\fi \\no@bbtrue \\epsf@getbb{\\@p@sbbfile} \\ifno@bb \\else \\bb@cull\\epsf@llx\\epsf@lly\\epsf@urx\\epsf@ury\\fi } \\def\\bb@cull#1#2#3#4{ \\dimen100=#1 bp\\edef\\@p@sbbllx{\\number\\dimen100} \\dimen100=#2 bp\\edef\\@p@sbblly{\\number\\dimen100} \\dimen100=#3 bp\\edef\\@p@sbburx{\\number\\dimen100} \\dimen100=#4 bp\\edef\\@p@sbbury{\\number\\dimen100} \\no@bbfalse } \\newdimen\\p@intvaluex \\newdimen\\p@intvaluey \\def\\rotate@#1#2{{\\dimen0=#1 sp\\dimen1=#2 sp \\global\\p@intvaluex=\\cosine\\dimen0 \\dimen3=\\sine\\dimen1 \\global\\advance\\p@intvaluex by -\\dimen3 \\global\\p@intvaluey=\\sine\\dimen0 \\dimen3=\\cosine\\dimen1 \\global\\advance\\p@intvaluey by \\dimen3 }} \\def\\compute@bb{ \\no@bbfalse \\if@bbllx \\else \\no@bbtrue \\fi \\if@bblly \\else \\no@bbtrue \\fi \\if@bburx \\else \\no@bbtrue \\fi \\if@bbury \\else \\no@bbtrue \\fi \\ifno@bb \\bb@missing \\fi \\ifno@bb \\ps@typeout{FATAL ERROR: no bb supplied or found} \\no-bb-error \\fi \\count203=\\@p@sbburx \\count204=\\@p@sbbury \\advance\\count203 by -\\@p@sbbllx \\advance\\count204 by -\\@p@sbblly \\edef\\ps@bbw{\\number\\count203} \\edef\\ps@bbh{\\number\\count204} \\if@angle \\Sine{\\@p@sangle}\\Cosine{\\@p@sangle} {\\dimen100=\\maxdimen\\xdef\\r@p@sbbllx{\\number\\dimen100} \\xdef\\r@p@sbblly{\\number\\dimen100} \\xdef\\r@p@sbburx{-\\number\\dimen100} \\xdef\\r@p@sbbury{-\\number\\dimen100}} \\def\\minmaxtest{ \\ifnum\\number\\p@intvaluex<\\r@p@sbbllx \\xdef\\r@p@sbbllx{\\number\\p@intvaluex}\\fi \\ifnum\\number\\p@intvaluex>\\r@p@sbburx \\xdef\\r@p@sbburx{\\number\\p@intvaluex}\\fi \\ifnum\\number\\p@intvaluey<\\r@p@sbblly \\xdef\\r@p@sbblly{\\number\\p@intvaluey}\\fi \\ifnum\\number\\p@intvaluey>\\r@p@sbbury \\xdef\\r@p@sbbury{\\number\\p@intvaluey}\\fi } \\rotate@{\\@p@sbbllx}{\\@p@sbblly} \\minmaxtest \\rotate@{\\@p@sbbllx}{\\@p@sbbury} \\minmaxtest \\rotate@{\\@p@sbburx}{\\@p@sbblly} \\minmaxtest \\rotate@{\\@p@sbburx}{\\@p@sbbury} \\minmaxtest \\edef\\@p@sbbllx{\\r@p@sbbllx}\\edef\\@p@sbblly{\\r@p@sbblly} \\edef\\@p@sbburx{\\r@p@sbburx}\\edef\\@p@sbbury{\\r@p@sbbury} \\fi \\count203=\\@p@sbburx \\count204=\\@p@sbbury \\advance\\count203 by -\\@p@sbbllx \\advance\\count204 by -\\@p@sbblly \\edef\\@bbw{\\number\\count203} \\edef\\@bbh{\\number\\count204} } \\def\\in@hundreds#1#2#3{\\count240=#2 \\count241=#3 \\count100=\\count240\t% \\divide\\count100 by \\count241 \\count101=\\count100 \\multiply\\count101 by \\count241 \\advance\\count240 by -\\count101 \\multiply\\count240 by 10 \\count101=\\count240\t% \\divide\\count101 by \\count241 \\count102=\\count101 \\multiply\\count102 by \\count241 \\advance\\count240 by -\\count102 \\multiply\\count240 by 10 \\count102=\\count240\t% \\divide\\count102 by \\count241 \\count200=#1\\count205=0 \\count201=\\count200 \\multiply\\count201 by \\count100 \\advance\\count205 by \\count201 \\count201=\\count200 \\divide\\count201 by 10 \\multiply\\count201 by \\count101 \\advance\\count205 by \\count201 \\count201=\\count200 \\divide\\count201 by 100 \\multiply\\count201 by \\count102 \\advance\\count205 by \\count201 \\edef\\@result{\\number\\count205} } \\def\\compute@wfromh{ \\in@hundreds{\\@p@sheight}{\\@bbw}{\\@bbh} \\edef\\@p@swidth{\\@result} } \\def\\compute@hfromw{ \\in@hundreds{\\@p@swidth}{\\@bbh}{\\@bbw} \\edef\\@p@sheight{\\@result} } \\def\\compute@handw{ \\if@height \\if@width \\else \\compute@wfromh \\fi \\else \\if@width \\compute@hfromw \\else \\edef\\@p@sheight{\\@bbh} \\edef\\@p@swidth{\\@bbw} \\fi \\fi } \\def\\compute@resv{ \\if@rheight \\else \\edef\\@p@srheight{\\@p@sheight} \\fi \\if@rwidth \\else \\edef\\@p@srwidth{\\@p@swidth} \\fi } \\def\\compute@sizes{ \\compute@bb \\if@scalefirst\\if@angle \\if@width \\in@hundreds{\\@p@swidth}{\\@bbw}{\\ps@bbw} \\edef\\@p@swidth{\\@result} \\fi \\if@height \\in@hundreds{\\@p@sheight}{\\@bbh}{\\ps@bbh} \\edef\\@p@sheight{\\@result} \\fi \\fi\\fi \\compute@handw \\compute@resv} \\def\\OzTeXSpecials{ \\special{empty.ps /@isp {true} def} \\special{empty.ps \\@p@swidth \\space \\@p@sheight \\space \\@p@sbbllx \\space \\@p@sbblly \\space \\@p@sbburx \\space \\@p@sbbury \\space startTexFig \\space } \\if@clip{ \\if@verbose{ \\ps@typeout{(clip)} }\\fi \\special{empty.ps doclip \\space } }\\fi \\if@angle{ \\if@verbose{ \\ps@typeout{(rotate)} }\\fi \\special {empty.ps \\@p@sangle \\space rotate \\space} }\\fi \\if@prologfile \\special{\\@prologfileval \\space } \\fi \\if@decmpr{ \\if@verbose{ \\ps@typeout{psfig: Compression not available in OzTeX version \\space } }\\fi }\\else{ \\if@verbose{ \\ps@typeout{psfig: including \\@p@sfile \\space } }\\fi \\special{epsf=\\@p@sfile \\space } }\\fi \\if@postlogfile \\special{\\@postlogfileval \\space } \\fi \\special{empty.ps /@isp {false} def} } \\def\\DvipsSpecials{ \\special{ps::[begin] \t\\@p@swidth \\space \\@p@sheight \\space \\@p@sbbllx \\space \\@p@sbblly \\space \\@p@sbburx \\space \\@p@sbbury \\space startTexFig \\space } \\if@clip{ \\if@verbose{ \\ps@typeout{(clip)} }\\fi \\special{ps:: doclip \\space } }\\fi \\if@angle \\if@verbose{ \\ps@typeout{(clip)} }\\fi \\special {ps:: \\@p@sangle \\space rotate \\space} \\fi \\if@prologfile \\special{ps: plotfile \\@prologfileval \\space } \\fi \\if@decmpr{ \\if@verbose{ \\ps@typeout{psfig: including \\@p@sfile.Z \\space } }\\fi \\special{ps: plotfile \"`zcat \\@p@sfile.Z\" \\space } }\\else{ \\if@verbose{ \\ps@typeout{psfig: including \\@p@sfile \\space } }\\fi \\special{ps: plotfile \\@p@sfile \\space } }\\fi \\if@postlogfile \\special{ps: plotfile \\@postlogfileval \\space } \\fi \\special{ps::[end] endTexFig \\space } } \\def\\psfig#1{\\vbox { \\ps@init@parms \\parse@ps@parms{#1} \\compute@sizes \\ifnum\\@p@scost<\\@psdraft{ \\PsfigSpecials \\vbox to \\@p@srheight sp{ \\hbox to \\@p@srwidth sp{ \\hss } \\vss } }\\else{ \\if@draftbox{ \\hbox{\\fbox{\\vbox to \\@p@srheight sp{ \\vss \\hbox to \\@p@srwidth sp{ \\hss % \\hss } \\vss }}} }\\else{ \\vbox to \\@p@srheight sp{ \\vss \\hbox to \\@p@srwidth sp{\\hss} \\vss } }\\fi    }\\fi }} \\psfigRestoreAt \\setDriver \\let\\@=\\LaTeXAtSign      \\def\\VEV#1{{\\left\\langle #1 \\right\\rangle}} \\def\\lesssim{\\mathrel{\\rlap{\\lower4pt\\hbox{\\hskip1pt$\\sim$}} \\raise1pt\\hbox{$<$}}}         % \\def\\gtrsim{\\mathrel{\\rlap{\\lower4pt\\hbox{\\hskip1pt$\\sim$}} \\raise1pt\\hbox{$>$}}}         % \\def\\hatn{{\\bf \\hat n}}   \\runninghead{Theory of Cosmic Microwave Background Polarization}{Cabella and Kamionkowski}  \\normalsize\\textlineskip \\thispagestyle{empty} \\setcounter{page}{1}   \\vspace*{0.88truein}  \\fpage{1} \\centerline{\\bf THEORY OF COSMIC MICROWAVE BACKGROUND POLARIZATION\\footnote{Lectures given at the 2003 Villa Mondragone School of Gravitation and Cosmology: ``The Polarization of the Cosmic Microwave Background,'' Rome, Italy, September 6--11, 2003.}} \\vspace*{0.37truein} \\centerline{\\footnotesize PAOLO CABELLA\\footnote{Paolo.Cabella@roma2.infn.it}} \\vspace*{0.015truein} \\centerline{\\footnotesize\\it Dipartimento di Fisica, Universit\\`a Tor Vergata, Roma I-00133, Italy} \\vspace*{10pt} \\centerline{\\footnotesize MARC KAMIONKOWSKI\\footnote{kamion@tapir.caltech.edu}} \\vspace*{0.015truein} \\centerline{\\footnotesize\\it California Institute of Technology, Mail Code 130-33, Pasadena, CA~~91125, USA}  \\vspace*{0.21truein} \\abstracts{These lectures introduce some of the basic theory of cosmic microwave background (CMB) polarization with the primary aim of developing the theory of CMB polarization from inflationary gravitational waves, as well as some of the related theory of weak gravitational lensing (cosmic shear) of CMB polarization.  We begin with production of polarization by Thomson scattering.  We then discuss tensor-harmonic analysis (the ``grad-curl'' or ``E-B'' decomposition) on a flat and full sky in some detail.  The Boltzmann/Einstein equations required to predict the CMB temperature/polarization pattern due to primordial gravitational waves are derived.  We show that gravitational waves produce a curl component of the CMB polarization while density perturbations (at linear order) do not. We then show how cosmic shear induces a curl component from a curl-free surface of last scattering.  We describe, though in less detail, how higher-order correlations can be used to subtract the cosmic-shear--induced curl.  Several exercises are provided.}{}{}   \\vspace*{1pt}\\textlineskip      % ",
        "introduction": "Just five years ago, the small-scale structure of the CMB was hidden behind a veil of experimental limitation.  That is no longer the case.  We now know empirically that the CMB has a wealth of detailed information in its temperature \\cite{cmbexperiments} and polarization pattern \\cite{polarizationdetections}.  We are now also confident in our theoretical understanding of these fluctuations: they are produced by a nearly scale-invariant spectrum of primordial density perturbations.  With a solid theoretical foundation, we are poised to move considerably further with future more precise CMB experiments.  Prospects for new advances include an even more detailed picture of the physics at the surface of last scattering, the physics of inflation, gravitational-wave backgrounds, the distribution of mass in the intermediate-redshift and current Universe, and perhaps the behavior of the dark energy.  These lectures will {\\it not} discuss all these wonderful possibilities.  Instead, we focus here on detailing some of the basics of the theory of CMB polarization with the aim primarily of understanding the polarization due to inflationary gravitational waves, as well as the related effects of weak gravitational lensing (``cosmic shear'') on CMB polarization. There is now a vast literature on CMB theory, including a number of reviews \\cite{reviews} and even a few books \\cite{books}.  Thus, experts are unlikely to find anything new in here.  However, some students may find that the discussion and level of detail for the specific subjects on which we focus provide a useful complement to these other more comprehensive sources.  In particular, we discuss tensor harmonics (the ``curl-grad'' or ``E-B'' decomposition) on both a flat and a full sky in detail, and we provide a detailed justification, from scratch, of the statement that gravitational waves produce a curl component in the CMB while density perturbations do not.  We show in detail that cosmic shear induces a curl component in the CMB polarization, but our discussion of how this shear may be subtracted with higher-order correlations is a bit sketchier.  We do not deal with any of the many important issues in data analysis.  The plan of the lectures is as follows:  We begin in Section 2 by reviewing how Thomson scattering produces polarization. Section 3 develops tensor harmonic analysis on a flat sky and Section 4 deals briefly with flat-sky correlation functions. Section 5 deals with tensor harmonic analysis on the full sky and Section 6 with the correlation functions.  Section 7 provides some comments, and in Section 8 we discuss some of the effects of cosmic shear on the polarization pattern.  A few exercises are sprinkled throughout to help conscientious students develop their facility with subject matter.  We reference primarily the most important original papers on relevant subjects and/or those that were particularly useful to us in the preparation of these lectures, but our referencing does not provide a full guide to the literature.  For that we refer the reader to the more comprehensive reviews \\cite{reviews}. ",
        "conclusions": ""
    },
    "0403/astro-ph0403501_arXiv.txt": {
        "abstract": "The 6dF Galaxy Survey (6dFGS) aims to measure the redshifts of around \\aimsixdf\\ galaxies, and the peculiar velocities of a 15\\,000-member sub-sample, over almost the entire southern sky. When complete, it will be the largest redshift survey of the nearby universe, reaching out to about $z\\sim0.15$, and more than an order of magnitude larger than any peculiar velocity survey to date.  The targets are all galaxies brighter than $K_{\\rm tot} = 12.75$ in the 2MASS Extended Source Catalog (XSC), supplemented by 2MASS and SuperCOSMOS galaxies that complete the sample to limits of $(H, J, r_F, b_J) = (13.05, 13.75, 15.6, 16.75)$. Central to the survey is the Six-Degree Field (6dF) multi-fibre spectrograph, an instrument able to record 150 simultaneous spectra over the $5.7^\\circ$-field of the UK Schmidt Telescope. An adaptive tiling algorithm has been employed to ensure around 95\\% fibering completeness over the 17046~deg$^2$ of the southern sky with $|\\,b\\,|>10^\\circ$. Spectra are obtained in two observations using separate V and R gratings, that together give $R \\sim 1000$ over at least 4000 -- 7500\\,\\AA\\ and signal-to-noise ratio $\\sim 10$ per pixel. Redshift measurements are obtained semi-automatically, and are assigned a quality value based on visual inspection. The 6dFGS database is available at {\\tt http://www-wfau.roe.ac.uk/6dFGS/}, with public data releases occuring after the completion of each third of the survey. ",
        "introduction": "\\label{sec:introduction}  Wide-scale redshift surveys such as the 2dF Galaxy Redshift Survey and Sloan Digital Sky Survey (2dFGRS, Colless \\etal\\ 2001b; SDSS, York \\etal\\ 2001) have made significant advances in our understanding of the matter and structure of the wider universe.  These include the precise determination of the luminosity function of galaxies (\\eg\\ Folkes \\etal\\ 1999, Cross \\etal\\ 2001, Cole \\etal\\ 2001, Blanton \\etal\\ 2001, Madgwick \\etal\\ 2002, Norberg \\etal\\ 2002, Blanton \\etal\\ 2003), the space density of nearby rich galaxy clusters (De Propris \\etal\\ 2002, Goto \\etal\\ 2003), and large-scale structure formation and mass density (Peacock \\etal\\ 2001, Percival \\etal\\ 2001, Efstathiou \\etal\\ 2002, Verde \\etal\\ 2002, Lahav \\etal\\ 2002, Zehavi \\etal\\ 2002, Hawkins \\etal\\ 2003, Szalay \\etal\\ 2003).  While such surveys have also greatly refined our view of the local universe, better determination of several key parameters can be made where a knowledge of galaxy mass can be combined with redshift. The 2dFGRS and SDSS surveys are both optically-selected, inevitably biasing them in favour of currently star-forming galaxies. The 2dF and Sloan spectrographs also have fields of view too small to allow full sky coverage in realistic timescales, limiting their utility for dynamical and cosmographic studies. The 6dF Galaxy Survey (6dFGS) is a dual redshift/peculiar velocity survey that endeavours to overcome the limitations of the 2dFGRS and SDSS surveys in these areas.  The primary sample for the 6dFGS is selected in the $K_s$ band from the 2MASS survey (Jarrett \\etal\\ 2000). The magnitudes used in the selection are estimated total magnitudes. These features combined mean that the primary sample is as unbiased a picture of the universe, in terms of the old stellar content of galaxies, as is possible at the current time. The near-infrared selection also enables the survey to probe closer to the Galactic equator before extinction becomes an issue; the survey covers the entire southern sky with $|\\,b\\,|>10^\\circ$.  The redshift of a galaxy includes both recessional and peculiar velocity components, so that a redshift survey alone does not furnish a true three-dimensional distribution for the galaxies. However, by measuring these components separately, it is possible to determine the three-dimensional distributions of both the galaxies and the underlying mass.  Observationally, significantly greater effort is required to obtain the distances and peculiar velocities of galaxies than their redshifts. Distance estimators based on the Fundamental Plane of early-type galaxies (FP; Dressler \\etal\\ 1987, Djorgovski \\& Davis 1987), as used in the 6dFGS, require measurements of the galaxies' internal velocity dispersions. Velocity dispersions demand a signal-to-noise ratio (S/N) in the spectrum at least 3 to 5 times higher than redshift measurements. Furthermore, the FP distance estimators need such spectroscopy to be supported by photometry which can be used to determine the galaxies' surface brightness profiles.  We measure peculiar velocities as the discrepancy between the redshift and the estimated distance. For realistic cosmologies, peculiar velocities increase weakly with distance, but remain $<1000 \\kms$. Redshift errors are small, and have little or no dependence on distance. On the other hand, all existing distance estimators have significant, intrinsic and fractional, uncertainties in distance; for the FP estimators this uncertainty is typically about 20\\% for a single galaxy measurement. This linear increase in the errors with distance, compared with more or less fixed peculiar velocities, means that the uncertainty on a single peculiar velocity becomes dominated by the intrinsic uncertainties at redshifts $cz \\sim 5000 \\kms$. Consequently, all previous peculiar velocity surveys have traced the velocity field, and hence the mass distribution, only out to distances of about 5000\\kms\\ (for relatively dense field samples of individual galaxies; \\eg\\ Dressler \\etal\\ 1987, Giovanelli \\etal\\ 1998, da Costa \\etal\\ 2000) or 15000\\kms\\ (for relatively sparse cluster samples, where distances for multiple galaxies are combined; \\eg\\ Lauer and Postman 1994, Hudson \\etal\\ 1999, Colless \\etal\\ 2001a). The smaller volumes are highly subject to cosmic variance, while the larger volumes are too sparsely sampled to reveal much information about the velocity field. In order to differentiate cosmological models, and constrain their parameters, both the survey volume and galaxy sampling need to be significantly increased. Moreover, since peculiar velocities amount to at most a few per cent of the velocity over much of the volume, both the underlying sample and the direct distance estimates have to be extraordinarily homogeneous, preferably involving unchanging telescopes, instrumental set ups and procedures in each case.  The two facets of the 6dF Galaxy Survey are a redshift survey of \\aimsixdf\\ galaxies over the southern sky, and a peculiar velocity survey of a subset of some 15\\,000 of these galaxies.  The aims of the redshift survey are: (1)~To take a near-infrared selected sample of galaxies and determine their luminosity function as a function of environment and galaxy type.  From these, the stellar mass function and mean stellar fraction of the local universe can be ascertained and compared to the integrated stellar mass inferred from measures of cosmic star formation history. (2)~To measure galaxy clustering on both small and large scales and its relation to stellar mass. In doing so, the 6dFGS will provide new insight into the scale-dependence of galaxy biasing and its relationship to dark matter. (3)~To determine the power spectrum of this galaxy clustering on scales similar to those spanned by the 2dF Galaxy Redshift Survey and Sloan Digital Sky Surveys. (4)~To delineate the distribution of galaxies in the nearby universe. The 6dFGS will not cover as many galaxies as either the 2dF Galaxy Redshift Survey or Sloan Digital Sky Survey. However, because it is a large-volume survey of nearby galaxies, it provides the ideal sample on which to base a peculiar velocity survey. In this respect, the complete 6dFGS will be more than ten times larger in number and span twice the volume of the PSCz Survey (Saunders \\etal\\ 2000), previously the largest survey of the local universe. (5)~To furnish a complete redshift catalogue for future studies in this regime. As discussed above, precise distance measures impose stronger demands on signal-to-noise than do redshifts.  With such higher quality spectra, it is also possible to infer properties of the underlying stellar population, such as ages and chemical abundances. Having such measurements along with the galaxy mass and knowledge of its local envrionment, will afford unprecedented opportunities to understand the processes driving galaxy formation and evolution.  This redshift catalogue will provide the basis for a volume-limited sample of early-type galaxies for the peculiar velocity survey. The aims of the peculiar velocity survey are: (1)~A detailed mapping of the density and peculiar velocity fields to around 15000\\kms\\, over half the local volume. (2)~The inference of the ages, metallicities and star-formation histories of E/S0 galaxies from the most massive systems down to dwarf galaxies. The influence of local galaxy density on these parameters will also be of key interest to models of galaxy formation. (3)~To understand the bias of galaxies (number density versus total mass density field) and its variation with galaxy parameters and environment.  One novel feature of the 6dF Galaxy Survey compared to earlier redshift and peculiar velocity surveys is its near-infrared source selection. The main target catalogues are selected from the Two Micron All Sky Survey (2MASS; Jarrett \\etal\\ 2000) using total galaxy magnitudes in $JHK$. There are several advantages of choosing galaxies in these bands. First, the near-infrared spectral energy distributions (SEDs) of galaxies are dominated by the light of their oldest stellar populations, and hence, the bulk of their stellar mass. Traditionally, surveys have selected target galaxies in the optical where galaxy SEDs are dominated by younger, bluer stars. Second, the E/S0 galaxies that will provide the best targets for Fundamental Plane peculiar velocity measures represent the largest fraction by galaxy type of near-infrared-selected samples. Finally, the effects of dust extinction are minimal at long wavelengths. In the target galaxies, this means that the total near-infrared luminosity is not dependent on galaxy orientation and so provides a reliable measure of galaxy mass. In our own Galaxy, it means the 6dFGS can map the local universe nearer to the plane of the Milky Way than would otherwise be possible through optical selection.  In this paper we describe the key components contributing to the realisation of the 6dF Galaxy Survey. Section~\\ref{sec:instrument} describes the Six-Degree Field instrument; section~\\ref{sec:design} details the compilation of the input catalogues and the optimal placement of fields and fibres on sources therein; section~\\ref{sec:implementation} outlines the methods used to obtain and reduce the spectra, and derive redshifts from them; section~\\ref{sec:datarelease} summarises the First Data Release of \\uniqueDB\\ unique galaxy redshifts and describes the 6dF Galaxy Survey on-line database; section~\\ref{sec:conclusions} provides concluding remarks. ",
        "conclusions": "\\label{sec:conclusions}  The 6dF Galaxy Redshift Survey (6dFGS) is designed to measure redshifts for approximately \\aimsixdf\\ galaxies and the peculiar velocities of 15\\,000. The survey uses the 6dF multi-fibre spectrograph on the United Kingdom Schmidt Telescope, which is capable of observing up to 150 objects simultaneously over a $5.7^\\circ$-diameter field of view.  The 2MASS Extended Source Catalog (Jarrett \\etal\\ 2000) is the primary source from which targets have been selected. The primary sample has been selected with $K_{\\rm tot} \\leq 12.75$, where $K_{\\rm tot}$ denotes the total $K$-band magnitude as derived from the isophotal 2MASS $K$ photometry. Additional galaxies have been selected to complete the target list down to $(H, J, r_F, b_J) = (13.05, 13.75, 15.6, 16.75)$. Thirteen miscellaneous surveys complete the total target list.  The survey covers the entire southern sky (declination $\\delta < 0^\\circ$), save for the regions within $|\\,b\\,| \\leq 10^\\circ$ of the Galactic Plane. This area is has been tiled with around 1500 fields that effectively cover the southern sky twice over. An adaptive tiling algorithm has been used to provide a uniform sampling rate of 94\\%. In total the survey covers some 17\\,046\\,deg$^2$ and has a median depth of $\\bar{z}$=0.05. There are three stages to the observations, which initially target the declination strip $-42^\\circ<\\delta<-23^\\circ$, followed by the equatorial region $-23^\\circ<\\delta<0^\\circ$, and conclude around the pole, ($\\delta<-42^\\circ$).  Spectra are obtained through separate V and R gratings and later spliced to produce combined spectra spanning 4000 -- 8400\\,\\AA. The spectra have 5 -- 6\\,\\AA\\ FWHM resolution in V and 9 -- 12\\,\\AA\\ resolution in R.  Software is used to estimate redshifts from both cross-correlation with template absorption-line spectra, and linear fits to the positions of strong emission lines. Each of these automatic redshift estimates is checked visually and assigned a quality $Q$ on a scale of 1 to 5, where $Q \\ge 3$ covers the range of reliable redshift measurements. The median signal-to-noise ratio is 9.4~\\perpix\\ for redshifts with quality $Q=4$, and 5.3~\\perpix\\ for $Q=3$ redshifts.  The data in this paper constitute the First Data Release of \\totalDB\\ observed spectra and the \\uniqueDB\\ unique extragalactic redshifts from this set. The rates of contamination by Galactic and failed spectra are 4\\% and 11\\% respectively. Data from the 6dF Galaxy Survey are publicly available through an online database at {\\tt http://www-wfau.roe.ac.uk/6dFGS/}, searchable through either SQL query commands or a online WWW form.  The main survey web site can be found at {\\tt http://www.mso.anu.edu.au/6dFGS}."
    },
    "0403/astro-ph0403098_arXiv.txt": {
        "abstract": "We study the Internal Linear Combination (ILC) method presented by the \\emph{Wilkinson Microwave Anisotropy Probe} (\\emph{WMAP}) science team, with the goal of determining whether it may be used for cosmological purposes, as a template-free alternative to existing foreground correction methods. We conclude that the method does have the potential to do just that, but great care must be taken both in implementation, and in a detailed understanding of limitations caused by residual foregrounds which can still affect cosmological results. As a first step we demonstrate how to compute the ILC weights both accurately and efficiently by means of Lagrange multipliers, and apply this method to the observed data to produce a new version of the ILC map. This map has 12\\% lower variance than the ILC map of the \\emph{WMAP} team, primarily due to less noise.  Next we describe how to generate Monte Carlo simulations of the ILC map, and find that these agree well with the observed map on angular scales up to $\\ell \\approx 200$, using a conservative sky cut. Finally we make two comments to the on-going debates concerning the large-scale properties of the \\emph{WMAP} data. First, we note that the Galactic south-eastern quadrant is associated with notably different ILC weights than the other three quadrants, possibly indicating a foreground related anisotropy. Second, we study the properties of the quadrupole and octopole (amplitude, alignment and planarity), and reproduce the previously reported results that the quadrupole and octopole are strongly aligned and that the octopole is moderately planar. Even more interestingly, we find that the $\\ell=5$ mode is spherically symmetric at about $3\\sigma$, and that the $\\ell=6$ mode is planar at the $2\\sigma$ level. However, we also assess the impact of residual foregrounds on these statistics, and find that the ILC map is not clean enough to allow for cosmological conclusions. Alternative methods must be developed to study these issues further. ",
        "introduction": "The first-year release of the \\emph{Wilkinson Microwave Anisotropy Probe} (\\emph{WMAP}; Bennett et al.\\ 2003a) data set has presented the cosmological community with an extraordinarily rich source of high-quality information, allowing the constraint of specific cosmological models and the parameters which define them to percentage accuracy.  Nevertheless, there remains an important goal beyond such a statistical assessment of the Cosmic Microwave Background (CMB) sky, namely to build an accurate image of the last-scattering surface which captures the detailed nature and morphology of our universe, and not simply some best-fit ensemble averaged view of it. Impediments to this program include instrumental noise and systematic artifacts, and foreground emission from local astrophysical objects. On a fundamental level non-cosmological foregrounds can easily compromise any conclusion regarding primordial physics unless properly accounted for, while on a practical level they complicate both algorithms and analyses. Methods for either removing, suppressing or at the very least constraining foregrounds are therefore of great importance, and, indeed, direct attacks on the raw data are very rarely justified. Practically any analysis must consider sky maps which have been processed in some way, either by explicit foreground corrections, or by introducing a sky cut.  The importance of foreground removal has been recognized in the community for a long time (e.g.,\\ Banday \\& Wolfendale 1991; Readhead \\& Lawrence 1992; Brandt et al.\\ 1994; Tegmark \\& Efstathiou 1996; Tegmark 1998; Tegmark et al.\\ 2000), as has the preferred method for discriminating against such contamination, namely multi-frequency observations. While the CMB itself contributes equally to all frequencies (as measured in thermodynamic temperature units) due to the black body nature of its spectrum, foregrounds are typically strongly frequency-dependent. One may therefore distinguish between foregrounds and genuine CMB anisotropy by studying how the signal varies with frequency.  However, detailed subtraction of foregrounds has generally required one of two assumptions to be made: either that all of the physical components of the foreground emission and their spectral behavior are known, or that accurate templates of all of the components are available and that the appropriate spectral behavior can be determined by fitting the templates to the available multi-frequency data. Neither method is easily adapted to accommodate real spatial variations in the spectral behavior of the foregrounds.  The \\emph{WMAP} project appears to have systematic issues under control, whilst considerable effort has been expended on foreground issues, and uncertainties may still remain.  Recent detections of non-Gaussianity (Chiang et al. 2003; Coles et al. 2004; Eriksen et al. 2004a; Naselsky et al. 2003; Park 2004; Vielva et al. 2004), the continuing debate over the low amplitude of large angular scale anisotropy (see e.g., Efstathiou 2004), and a possible preferred direction or alignment contained therein (de Oliveira-Costa et al.\\ 2004; Eriksen et al.\\ 2004b) may yet be affected by improvements of our ability to remove non-cosmological foregrounds.  The \\emph{WMAP} satellite observes the sky at five frequencies (23, 33, 41, 61 and 94 GHz), and using at least in part this information the \\emph{WMAP} team have applied three different methods for removing, constraining or describing the foregrounds (Bennett et al.\\ 2003b).  The first method is to produce a so-called internal linear combination map (from now on denoted ILC), which assumes nothing about the particular frequency dependencies or morphologies of the foregrounds. Instead, a CMB map is reconstructed by co-adding the data at the five frequencies (now convolved to a common angular resolution of 1 degree) with a set of weights that minimizes the final variance of the map. The details of the non-linear method adopted to derive these weights have not been described by the \\emph{WMAP} team. In order to accommodate spectral variability of the foregrounds, the sky has been partitioned into 12 separate regions and the minimum variance criterion applied to each in turn to determine the weights. Discontinuities between regions have been minimized by smoothing of the weights at the boundaries. The resultant CMB map is visually satisfactory but has complex noise properties, and indeed the \\emph{WMAP} team \\emph{explicitly} warns against its use for cosmological analysis. Nevertheless, the map has been subjected to such studies in the literature, and indeed the \\emph{WMAP} team do use the resultant map as an input to their second foreground removal technique.  This involves the application of a Maximum Entropy Method (MEM) in order to construct a model of the foregrounds, component by component. The strength of this method in principle is its ability to reconstruct the synchrotron, free-free and dust emission and their detailed frequency dependence on a pixel-by-pixel basis.  However, the initial stage of the analysis must still utilize templates for these dominant foreground components, and also establishes priors on their spectral behavior by using the \\emph{WMAP} data at the five frequencies after correction for a CMB component as determined by the WILC method above. As we will see later, ILC-like methods in general still allow some leakage of foreground signal into the CMB reconstruction, and whether this results in any feedback is difficult to determine. Again as a consequence of complex noise properties, the resultant map has not been considered useful for cosmological purposes.  Instead, the \\emph{WMAP} team has used the results to interpret the nature of the foreground emission. In particular, they identify a dust-correlated component in the lower frequency (23, 33 and 41 GHz) channels with a spectral index of $\\beta \\sim -2.5$ for a spectrum of the form $\\nu^{\\beta}$.  This is physically interpreted as a flat spectrum synchrotron component in regions of star formation near the Galactic plane, rather than to emission from spinning dust, which had become the preferred solution to this anomalous, low frequency dust correlated emission.  The issue remains open, but recent reanalyses by Lagache (2003) and Finkbeiner (2004) find evidence for the latter interpretation.  The origin of this controversy lies simply with the fact that the component separation as implemented by \\emph{WMAP} is allowed only to produce a combined synchrotron/spinning dust solution at each frequency, with no attempt made to separately disentangle these two components.  The final method for foreground correction is perhaps the simplest of the three, and it is also the preferred method for generating cleaned maps suitable for cosmological analysis.  Rather than inherently exploiting the frequency information contained in the data, one subtracts external templates of the various physical components (i.e., maps produced by non-CMB observations made preferably at frequencies where a specific component dominates the emission) with coupling coefficients determined by cross-correlation with the observed maps. This avoids the noise amplification which occurs when one co-adds the \\emph{WMAP} data alone, and has the added benefit that the resultant maps have well-known noise properties, provided that the templates themselves do not contribute significantly to this. Difficulties associated with the method include uncertainties in the detailed morphologies of the templates as scaled to the wavelengths of interest, and the propagation of errors in the coupling coefficients into the error budgets of specific analyses.  It should be noted that Tegmark, de Oliveira-Costa, \\& Hamilton (2003) have applied a generalization of the ILC method to the \\emph{WMAP} data. The basic idea is to allow the weighting of each map to be scale-dependent by performing the analysis in harmonic space, the assumption being that this allows any spatial variations in the spectral dependence of the foregrounds to be adequately tracked.  It is not clear to what extent real variations project onto the harmonic eigenmodes of the analysis.  As with the ILC method, complex noise properties result, and so it is unlikely that this method is suitable for high precision cosmological analyses. In what follows we will denote the map as TCM -- the Tegmark et al.\\ cleaned map.  In this paper a new look is taken at the ILC method presented by \\citet{bennett:2003b}, with the main goal of determining whether a map derived in this manner can be suitable for cosmological purposes. Specifically, we derive a new ILC map based on Lagrange multipliers (in what follows to be referred to as the LILC -- Lagrange Internal Linear Combination -- map), which has 12\\% lower variance than the \\emph{WMAP} science team's ILC (hereafter referred to as WILC) map. We then generate Monte Carlo simulations of this map by adding white noise and foreground templates to CMB realizations, and process these through our ILC pipeline. This allows us to quantify the efficiency of the ILC method, and realistic foreground residual estimates may be established.  In the final section we repeat the large-scale analysis of \\citet{de Oliveira-Costa:2004} both for our new LILC map and for the simulations, to assess the impact of residual foregrounds on these statistics. However, we study not only the quadrupole and the octopole, but also consider the properties of the $\\ell=4, 5$ and 6 modes. In fact, we find that the properties of the latter two are at least as intriguing as those of the quadrupole and octopole: the $\\ell=5$ mode is highly spherically symmetric, and the $\\ell=6$ mode is planar. ",
        "conclusions": "\\label{sec:conclusions}  The main goal of this paper was to study whether the ILC method is able to yield cosmologically useful maps, and if so, whether realistic simulations can be generated in reasonable time in order to calibrate the uncertainties associated with the use of such a map. The results presented earlier suggest a cautiously positive conclusion -- the ILC method does have the capability of producing relatively clean CMB maps without the use of external templates.  Nevertheless, great care should be taken in the practical implementation of the method (e.g., the proper definition of the individual regions is a crucial step), and beyond this one needs to be highly aware of its limitations.  On a more detailed level, we derived the equations for the ILC weights based on Lagrange multipliers, which were also discussed by Tegmark (1998). While a non-linear search algorithm is based on iterations, this method solves one single linear system of equations, and is therefore much faster. This is important when generating Monte Carlo simulations. Subsequently, we discussed how to produce realistic simulations of the ILC map, and used these simulations to study the properties of the method itself, with particular emphasis on the sensitivity to noise and sky cuts.  The method was applied to the real \\emph{WMAP} data, and the resultant LILC map was determined to have properties similar to the TCM map, but somewhat different from the \\citet{bennett:2003b} WILC map. We also computed ILC weights for four quadrants of the sky, and found that the south-eastern Galactic quadrant has significantly different properties than the other three, possibly shedding new light on the asymmetry issue discussed by \\citet{Eriksen:2004a}.  Finally, as a comment to the on-going debate on the nature of the large-angular scale anisotropy, we investigated the implications of the LILC map for estimates of the quadrupole and octopole modes, and found that the new quadrupole moment increases from $195\\;\\mu \\textrm{K}$ to $351\\;\\mu \\textrm{K}$, which is a perfectly acceptable amplitude compared to the best-fit spectrum. However, the alignment between the quadrupole and the octopole is stronger in our map than in the WILC and the TCM. We also pointed out that the $\\ell=5$ and 6 modes are most peculiar in their symmetry properties, as only 0.2\\% of the simulations have a more spherically symmetric $\\ell=5$ mode than the \\emph{WMAP} data, and 1.4\\% a more planar $\\ell=6$ mode. Further, since we have access to the full sky, these modes are all independent under the Gaussian, random-phase hypothesis, and the probabilities therefore accumulate quite straightforwardly. The major caveat, however, is that many of these measurements are derived from maps with complex foreground and noise properties, and definitive cosmological conclusions therefore remain elusive. Better foreground correction methods are required, or, alternatively, methods for studying the same properties on a cut sky should be developed. This work is already under way, and will be published in a future paper.  Returning to the ILC method, one may question whether the minimum variance criterion in itself is a meaningful measure of performance. As we have seen, this criterion implies a trade-off between suppressing noise and foregrounds, and moderate levels of foregrounds are often accepted in order to suppress noise. For most practical cosmological analyses this is not likely to be acceptable; noise is more easily quantified than residual foregrounds.  Note therefore that although we do provide a copy of the LILC map at H.K.E.'s home page\\footnote{http://www.astro.uio.no/$\\sim$hke/cmbdata/WMAP\\_ILC\\_lagrange.fits}, we strongly advise against using it for purposes beyond visual presentation, for which, of course, the official WILC map is perfectly acceptable."
    },
    "0403/astro-ph0403051_arXiv.txt": {
        "abstract": "{We present iron abundance analysis for 23 wide binaries with main sequence components in the temperture range 4900-6300~K, taken from the sample of the pairs currently included in the radial velocity planet search on going at the Telescopio Nazionale Galileo (TNG) using the high resolution spectrograph SARG. The use of a line-by-line differential analysis technique between the components of each pair allows us to reach errors of about 0.02 dex in the iron content difference. Most of the pairs have abundance differences lower than 0.02 dex and there are no pairs with differences larger than 0.07 dex. The four cases of differences larger than 0.02 dex may be spurious because of the larger error bars affecting pairs with large temperature difference, cold stars and rotating stars. The pair HD~219542, previously reported by us to have a different composition, here is shown to be normal. For non-rotating stars warmer than 5500 K, characterized by a thinner convective envelope and for which our analyis appears to be of higher accuracy, we are  able to exclude in most cases the consuption of more than 1 Earth Mass of iron (about 5 Earth masses of meteoritic material) during the main sequence lifetime of the stars, placing more stringent limits (about 0.4 Earth masses of iron) in five cases of warm stars. This latter limit is similar to the estimates of rocky material accreted by the Sun during its main sequence lifetime. Combining the results of the present analysis with those for the Hyades and Pleiades, we conclude that the hypothesis that pollution by planetary material is the only mechanism responsible for the highest metallicity of the stars with planets may be rejected at more than 99\\% level of confidence if the incidence of planets in these samples is as high as 8\\% and similar to the field stars included in current radial velocity surveys. However, the significance of this result drops considerably if the incidence of planets around stars in binary systems and clusters is less than a half of that around normal field stars. ",
        "introduction": "\\label{s:intro}  The evidence for a high metal content in stars harbouring planets is becoming stronger as planet discoveries cumulate and suitable control samples are studied using strictly the same procedures. It appears that the metal abundance of stars with planets is on average about 0.25 dex larger than that of typical control sample stars (Laws et al.~\\cite{laws03}) and that the frequency of planets is a strong function of metallicity (Santos et al.~\\cite{santos04}, Fischer et al.~\\cite{fischer04}).  Two alternative hypotheses have been proposed to explain these observations: either the high metallicity is responsible for the presence of planets, making their formation easier; or the planets are the cause of the high metallicity, because of pollution of metal-rich planetary material onto the (outer region of the) central star (Gonzalez \\cite{gonzalez}).  Infall of planetesimals on the star during the early phases of planet formation is generally expected on the basis of current models of planet formation. The orbital migration proposed to explain the occurrence of the close-in giant planets found by radial velocity surveys also points to the infall on the star of portions of the proto-planetary disk.  Most of the accretion is expected to take place during the early phases of the evolution of the planetary system. However, when a star is still in the phase of gravitational contraction, its convective zone is much thicker than for main sequence stars (see e.g. Murray et al.~\\cite{murray}). In this case, the metal-rich material should be uniformly distributed by convective mixing over a large portion of the star, resulting in a negligible photospheric chemical alteration even for rather large amounts of accreted material. Late accretion, when the star is approaching or has already reached the main sequence, is likely required to produce observable differences. This may happen due to infall of residual planetesimal/asteroids. The age distribution of the main sequence stars hosting debris disks (Habing et al.~\\cite{habing01}), and of the impact craters on Solar System solid bodies, indicate that a significant infall of gas-poor material may be present until 300-400 Myr after star formation; at this epoch the convective zone of a solar-mass star has been thinned enough that significant alterations of the surface abundances may be produced. The ingestion of planets scattered toward the star by dynamical interactions (Marzari \\& Weidenschilling \\cite{marzari02}) might also produce metallicity enhancements at late phases.  The signatures of accretion should be more evident for stars with smaller mixing zones,  and disappear for evolved stars, when the outer convective region penetrates deeply into the star. The observations of lithium depletion for stars in the range $1.2-1.5~M_{\\odot}$ (the so-called Li-dip, Boesgaard \\& Tripicco \\cite{boesgaard86}) indicate that an additional  mixing mechanism is at work, other than standard convective mixing\\footnote{Alternative explanations of Li-dip depletion such as mass loss or gravitational sedimentation do not seem to be supported by observations, see Do Nascimento et al.~(\\cite{donascimento}).}. It must be further noted that the lithium dip is not present in open clusters as young as Pleiades (Pilachowski et al.~\\cite{pilachowski}), indicating an age dependence. Such an extra-mixing mechanism, besides destroying lithium, should also efficiently flatten the radial abundance profiles of the star, if enrichment of external layers occurred. Therefore the stars that should show the most prominent signatures of accretion are early G-dwarfs and late F-dwarfs, with masses smaller than those in the Li-dip (Murray et al.~\\cite{murray}). The occurrence of 'Li-dip' extra-mixing might also help in explaining the lack of F stars with very high metallicity ([Fe/H]$ > +0.50$), predicted by the pollution scenario when considering only the extension of the convective zone.  The metal abundance of planet hosts does not appear to be correlated with the thickness of the convective zone (Pinsonneault et al.~\\cite{pinsonneault}, Santos et al.~\\cite{santos04}). This seems to indicate that pollution is not the most relevant contribution to the high metallicity of stars with planets. On the other hand, two possible effects might in principle explain such a lack of a trend as a function of the thickness of the convective enevelope. The first is the result of Vauclair (\\cite{vauclair}), who suggests that, when the metallicity difference between the inner and outer parts of a star exceedes a threshold, further mixing is induced in the transition zone, therefore diluting the atmospheric metallicty enhancement. The second effect can be considered if the metal enrichment is dominated by the ingestion of already formed planets. According to the hydrodynamical simulations by Sandquist et al.~(\\cite{sandquist98}), a $1.22~M_{\\odot}$ star should be able to only partially consume an engulfed giant planet within its thin convective zone. Depending on the actual distribution of heavy elements within the planet (possibly mostly confined in the planet core), the resulting alteration of chemical composition might be even smaller than for a solar-like star that is able to consume nearly completely a Jupiter-like planet within its convective zone.  On the other hand, several observational results support the occurrence of pollution phenomena.  Smith et al.~(\\cite{smith}) found indications of systematic element-to-element differences as a function of the condensation temperature on some planet host stars. This could be explained by the infall of material devoid of light elements, such as rocky planets or asteroids. However, such chemical composition trends can be confused with galactic chemical evolution trends, and can be more easily studied in binaries and clusters. Gratton et al.~(\\cite{paper1}), hereafter Paper I, performed a differential analysis of 6 visual main sequence binaries. They found one pair (HD~219542) with a 0.09 dex iron content difference, four pairs with very similar composition and one ambiguous case (HD~200466). Laws \\& Gonzalez (\\cite{laws01}) also found a small difference between the components of 16 Cyg. A similar approach can be pursued to look for chemical anomalies among open cluster stars. Paulson et al.~(\\cite{paulson}) studied 55 FGK dwarfs in the Hyades. They found two stars with abundances 0.2 dex larger than the cluster mean, but their membership is questionable. The Pleiades were studied by Wilden et al.~(\\cite{wilden}); they found one possible candidate out of 15 stars with abundance 0.1 dex above the cluster mean and very little scatter for the other stars (0.02 dex). No correlation appears to be present between iron and lithium abundances.  The presence of $^6$Li might be a strong indication of accretion phenomena, since this element is easily destroyed and cannot be observed in main sequence stars unless recently fallen onto the star from an external source. The presence of $^6$Li in HD~82943 is controversial (Reddy et al.~\\cite{reddy}, Israelian et al.~\\cite{isr03}). If confirmed, this would indicate the infall of a giant planet. Anomalies of the most abundant  $^7$Li isotope are instead easier to find (see e.g. Pasquini et al.~\\cite{pasquini_m67}, King et al.~\\cite{king}, Martin et al.~\\cite{martin02}) but they are more difficult to interpret as they might not be linked to planetary pollution.  Murray et al.~(\\cite{murray}) found that the Sun should have ingested some $2~M_{\\oplus}$ of meteoritic material (about $0.4~M_{\\oplus}$ of iron), considering the drop of iron density in the asteroid region and the time distribution of the impact craters. This corresponds to a metallicity enhancement of 0.017 dex. Helioseismology currently does not allow a confirmation of the presence of such a small metal abundance difference between the inner and the outer regions (Winnick et al.~\\cite{winnick02}), with upper limits on the amount of accreted iron about four times higher than the Murray et al.~(\\cite{murray}) estimate\\footnote{The comparison between the results quoted in the original papers show some inconsistencies because of different assumptions on the composition of meteoritic material. In this paper, we adopt the meteoritic abundances by Lodders (\\cite{lodders}). The mass fraction of iron is thus 18.3\\% of the meteoritic material.}.  From this discussion, it appears that the comparison of the chemical composition of wide binaries is a very powerful approach to study the occurrence of planetary pollution. The number of pairs studied in Paper I is too small to allow significant inferences about the frequency and the amount of chemical anomalies in visual binaries. In this paper, we significantly enlarge  the number of pairs studied (23), with various improvements in the analysis procedure. We will then combine the results of the present analysis with the literature results for Hyades and Pleaides,  to draw tentative conclusions about the impact of planetary material pollution. ",
        "conclusions": "The frequency of planets is estimated to be about 8\\% (Butler et al.~\\cite{butler_manchester}) and the stars hosting planets appear to have metallicity larger than 0.25 dex with respect to normal field population stars (Laws et al.~\\cite{laws03}). We have not found any pair among 23 with such large composition differences\\footnote{Considering the observed planet frequency, and  the differences of the mass of the convective zone between the components, it appears very unlikely that both components of a binary system have such metal rich material that their original composition is alterated by a similar amount.}. One star out of 15 in Pleiades might have a metallicity enhancement of about 0.1 dex (Wilden et al.~\\cite{wilden}), while two out of 55 in Hyades (with questionable membership) have metallicity difference of about 0.20 dex with respect to the cluster mean (Paulson et al.~\\cite{paulson}).  The fraction of stars with metal enrichment comparable to the observed metal overabundance of stars with planets appears to be lower than the fraction of stars with planets. Taken at face value, this indicates that pollution phenomena {\\it alone} cannot explain the planet-metallicity connection. High metallicity then seems to play a role in favouring the presence of planets. Enrichment phenomena might be present, as possibly indicated by  the still controversial presence of $^6$Li in HD~82943, but they should be rare and/or produce only a small composition alteration. Increasing the statistics of wide binaries and stars in clusters in the most suitable temperature range should allow one to constrain the accretion scenario more tightly.  If we consider our sample and those of Pleiades and Hyades, we have 2 possible cases with a metallicity difference comparable to the typical offset between planet hosts and normal field stars (the two Hyades stars) out of 126 stars. With the hypothesis that 8\\% of these stars have planets (see below for an important caveat) and that the $\\sim0.25$ dex metallicity offset is entirely due to planetary pollution, we obtain a probability of less than 1\\% to observe just two cases out of 126 stars. However, such a probability reaches 13\\% if the actual planet frequency for these stars is 4\\%.  Some selection effect may be present considering that the temperature and mass distribution of some of the samples studied for the high-precision differential abundance analysis may be different with respect to the field stars selected for radial velocity planet searches.  A more fundamental caveat to this discussion is that the components of wide  binaries considered in this study (typical projected separations of about 100-400 AU, Table \\ref{t:mag}) as well as stars in open clusters like Pleiades and Hyades might have planet frequencies and pollution histories different from the mostly single field star planets hosts. There are some hints that the properties of planets in binary systems are different with respect to those orbiting single stars (Zucker \\& Mazeh \\cite{zucker02}). This seems to indicate a significant role of orbital interactions in the formation and evolution of planetary systems, whose effects on the accretion of planetary material should be considered. Dedicated searches for planets in various kind of binaries (Desidera et al.~\\cite{hd219542}; Eggenberger et al.~\\cite{eggenberger}) and in star clusters characterized by different dynamical conditions (see e.g. Gilliland et al.~\\cite{gilliland}; Bruntt et al.~\\cite{bruntt}; Street et al.~\\cite{street}; Piotto et al.~\\cite{p8}) are thus very useful to understand the effects of the dynamical environment on the presence of planets and their orbital properties. This will allow us also to determine whether the pollution properties derived by the study of wide pairs and cluster stars can be extended to single stars (for which this is much more difficult to show)."
    },
    "0403/astro-ph0403321_arXiv.txt": {
        "abstract": "{Evidence is presented from numerical magneto-hydrodynamical simulations for the existence of magnetic activity in late-type giant stars. A red supergiant with stellar parameters similar to that of Betelgeuse ($\\alpha$ Orionis) is modeled as a ``star-in-a-box\" with the high-order ``Pencil Code\". Both linear kinematic and non-linear saturated dynamo action are found: the non-linear magnetic field saturates at a super-equipartition value, while in the linear regime two different modes of dynamo action are found. It is speculated that magnetic activity of late-type giants may influence dust and wind formation and possibly lead to the heating of the outer atmospheres of these stars. ",
        "introduction": "There are indications from both dynamo theory and observations that some late-type giant stars such as red supergiants and asymptotic-giant-branch stars (AGB stars) may harbor magnetic fields. On the theoretical side, it has been suggested that non-spherically symmetric planetary nebulae (PNe) may be a result of the collimating effect of a strong magnetic field: Blackman et al.\\ (\\cite{Blackman+ea01}) studied interface dynamo models similar to the mean field theory's solar $\\alpha \\omega$-dynamo and found that the generated magnetic fields were strong enough to shape bipolar outflows, producing bipolar PNe, while also braking the stellar core thereby explaining the slow rotation of many white dwarf stars. Also using mean field dynamo theory Soker \\& Zoabi (\\cite{Soker+Zoabi2002}) propose instead an $\\alpha^2 \\omega$ dynamo due to the slow rotation of AGB stars rendering the $\\omega$-effect ineffective. They find that the magnetic field may reach strengths of $\\sim 100$ Gauss, significantly less than that found by Blackman et al.\\ (\\cite{Blackman+ea01}). On the one hand, they believe that the large-scale field is strong enough for the formation of magnetic cool spots (see also Soker \\& Kastner \\cite{Soker+Kastner2003} on AGB star flaring). These spots in turn may regulate dust formation, and hence the mass-loss rate, but the authors argue that they cannot explain the formation of non-spherical PNe (see also Soker \\cite{Soker2002}): on the other hand, the locally strong magnetic tension could enforce a coherent flow that may favor a maser process.  On the observational side of things, maser polarization is known to exist in circumstellar envelopes of AGB stars (e.g.\\ Gray et al.\\ \\cite{Gray+ea99}, Vlemmings et al.\\ \\cite{Vlemmings+ea03}, and recently Sivagnanam \\cite{Sivagnanam2004}) and X-ray emission has been observed from some cool giant stars (e.g.\\ H\\\"{u}nsch et al.\\ \\cite{Hunch+ea98} and Ayres et al.\\ \\cite{Ayres+ea03}). These observations are generally taken as evidence for the existence of magnetic activity in late-type giant stars (cf.\\ Soker \\& Kastner \\cite{Soker+Kastner2003}).  The cool star Betelgeuse (a.k.a.\\ $\\alpha$ Orionis) is an example of an abundantly observed late-type supergiant that displays irregular brightness variations interpreted as large-scale surface structures (e.g.\\ Lim et al.\\ \\cite{Lim+ea98} and Gray \\cite{Gray2000}). It is one of the stars with the largest apparent sizes on the sky---corresponding to a radius in the interval 600--800 ${\\rm R}_{\\odot}$. Freytag et al.\\ (\\cite{Freytag+ea02}) performed detailed numerical 3-d radiation-hydrodynamic (RHD) simulations of the convective envelope of the star under realistic physical assumptions, while trying to determine if the star's known brightness fluctuations may be understood as convective motions within the star's atmosphere: the resulting models were largely successful in explaining the observations as a consequence of giant-cell convection on the stellar surface, very dissimilar to solar convection. Dorch \\& Freytag (\\cite{Dorch+Freytag2002}) performed a kinematic dynamo analysis of the convective motions (i.e.\\ not including the back-reaction of the Lorentz force on the flow) and found that a weak seed magnetic field could indeed be exponentially amplified by the giant-cell convection on a time-scale of about 25 years.  This paper reports on full non-linear magneto-hydrodynamical (MHD) numerical simulations of dynamo action in a late-type supergiant star with fundamental stellar parameters set equal to that of Betelgeuse. The paper is organized as follows: Section \\ref{model} contains a description of the numerical model, code and setup, Section \\ref{results} describes and discusses the results, the convective flows and dynamo action, and Section \\ref{conclusion} contains a short summary and conclusion. ",
        "conclusions": "\\label{conclusion}  In summary three different modes of dynamo action are recognized: \\begin{enumerate} \\item A relatively fast growing linear mode with an exponential growth time of $\\sim 4$ years. \\item A relatively slowly growing pseudo-linear mode with an exponential growth of $\\sim 25$ years. \\item A saturated non-linear mode operating a factor of two above equipartition. \\end{enumerate} More modes may of course exist but these must then have very low growth rates and/or very small initial amplitudes since they have not appeared in the simulations. It is worth noting that in case 2) of the pseudo-linear mode, the same value of the growth time (around 25 years) was found in the previous purely kinematic dynamo models Dorch \\& Freytag (\\cite{Dorch+Freytag2002}) although they employed a different computational method. This may in fact not be so strange, since the growth rate in a kinematic dynamo is set in part by the convergence of the flow across the field lines $- \\nabla_{\\perp}\\cdot\\uv$ and if the flows are similar so should the growth rates be.  Based on the results presented here, it is not possible to state conclusively if Betelgeuse actually has a magnetic field, since such a field is unobserved. However, one may conclude that it seems possible that late-type giant stars such as Betelgeuse can indeed have presently undetected magnetic fields. These magnetic fields are likely to be close to or stronger than equipartition; they may be difficult to detect directly, due to the relatively small filling factors of the strong fields, but even the moderately strong fields may have influence on their immediate surroundings through altered dust, wind and mass-loss properties. The formation of dust in the presence of a magnetic field will be the subject of a subsequent paper along the lines presented here: the ``Pencil Code'' has recently been augmented with modules for radiation and dust modeling.  The dynamos of the late-type giant studied here may be characterized as belonging to the class called ``local small-scale dynamos'' another example of which is the proposed dynamo action in the solar photosphere that is sometimes claimed to be responsible for the formation of small-scale flux tubes (cf.\\ Cattaneo \\cite{Cattaneo1999}). However, in the case of Betelgeuse this designation is less meaningful since the generated magnetic field is both global and large-scale, but because of the slow rotation, no large-scale solar-like toroidal field is formed.  It is interesting to note that very recently Lobel et al.\\ (\\cite{Lobel+ea04}) published spatially resolved spectra of the upper chromosphere and dust envelope of Betelgeuse. Based on various emission lines they provide evidence for the presence of warm chromospheric plasma away from the star at around 40 R. The spectra reveal that Betelgeuse's upper chromosphere extends far beyond the circumstellar envelope. They compute that temperatures of the warm chromospheric gas exceed 2600 K. The presence of a hot chromosphere lead this author to speculate on the possible connection to coronal heating in the Sun, which is likely to be magnetic in origin and caused by flux braiding motions in the solar photosphere (cf.\\ Gudiksen \\& Nordlund \\cite{Gudiksen+Nordlund2002}): it remains to be proven whether a similar process could be operating in late-type giant stars."
    },
    "0403/astro-ph0403081.txt": {
        "abstract": "Observations with the Extreme Ultraviolet Explorer (EUVE) have shown the Coma Cluster to be a source of EUV emission in excess of that produced by X-ray gas in the cluster. We have re-examined the EUVE data on this cluster in an attempt to obtain clues as to the origin of this emission. We find two important new results. First, the ratio between the azimuthally averaged EUV excess emission and the ROSAT hard X-ray flux is constant as a function of distance from the cluster center outward. Second, a correlation analysis between the EUV excess emission and the X-ray emission shows that on a detailed level the EUV excess is spatially closely related to the X-ray emission. These findings contradict previous suggestions as to the underlying source of the diffuse EUV emission in Coma and provide important information in regards to the true source of this emission.  We propose a new explanation for the source of this emission: inverse Compton scattering of microwave background photons by secondary electrons and positrons. We explore this possibility in some detail and show that it is consistent with all of the available observational evidence. The parent cosmic ray protons may have been produced by any of a number of sources, including supernovae, active galaxies, galactic winds, and cluster formation shocks, but we believe that the most likely source is cluster formation shocks. If the EUV emission in the Coma Cluster is, in fact, the result of secondary electrons, this may be the only direct evidence for secondary electrons in the intracluster medium of a cluster of galaxies, since recent work suggests that secondary electrons may not be the cause of radio halos. \\\\\\ \\\\ ",
        "introduction": "Observations with the Extreme Ultraviolet Explorer (EUVE) provided evidence that a number of clusters of galaxies emit excess EUV emission in the cores of the clusters. The first clusters reported to have EUV excesses were the Virgo cluster (Lieu \\etal 1996a; Bowyer \\etal 1996) and the Coma Cluster (Lieu \\etal 1996b).  Thereafter EUV emission was reported for Abell 1795 (Mittaz, Lieu \\& Lockman, 1998) and Abell 2199 (Lieu \\etal 1999a). These early works employed a variety of data analysis schemes that were later found to be incorrect (Bowyer, Bergh\\\"ofer, \\& Korpela, 1999), primarily because incorrect methods were used to account for the sensitivity profile, or exposure map, of the telescope.  The only clusters that have been determined to have an EUV excess using uncontested data analysis procedures are the Virgo cluster (Bergh\\\"ofer \\etal 2000) and the Coma Cluster (Bowyer \\etal 1999).  Subsequent to the analysis of Bowyer \\etal (1999) additional EUV data on the Coma Cluster were obtained with EUVE. In this paper we re-examine the excess EUV emission in the Coma Cluster using all the EUVE data available on this cluster. We obtain important new information on the character of the EUV emission in this cluster. Given these new results, we provide strong evidence that the EUV emission is produced by secondary electrons and positrons in the intracluster medium (ICM).  This finding may well be the only secure evidence of the presence of secondary electrons and positrons in an intracluster medium. ",
        "conclusions": "We have analyzed archival data obtained with EUVE on the core of the Coma cluster. We find the ratio of the azimuthally averaged EUV and X-ray intensities is essentially constant with increasing cluster radius. In addition, a correlation of the diffuse EUV emission with the diffuse X-ray emission shows that the detailed spatial distributions of these emissions are quite similar, but not identical.  XMM-{\\em Newton} and {\\em Chandra} observations show that there is no intracluster 10$^6$ K gas in the core of the cluster and hence that the EUV emission must be non-thermal. The only viable non-thermal source for the EUV excess is photons inverse Compton scattered by 200 MeV electrons from the cosmic microwave background. To account for the observed EUV intensity distribution, the scattering electrons must be distributed with a spatial density roughly in proportion to the square of the thermal plasma density. A scenario that naturally produces these results is that the EUV emitting particles are electron/positron secondaries produced by inelastic collisions between primary cosmic ray protons and thermal protons with similar spatial distributions. This type of distribution would be a natural outcome if the cosmic rays had been produced over the past several Gyr and had become well-mixed with the intracluster medium.  This model accounts naturally for the average spatial distribution of the EUV emission in relation to the X-ray emission. It accounts for the similar, but imperfect, pixel-to-pixel correspondence of the EUV and X-ray emission. Finally, it accomplishes this without violating observational limits in other bands of the spectrum.  We have demonstrated that the required underlying cosmic rays could reasonably have been produced in large scale structure shocks accompanying the cluster formation.  Secondary electrons as the source of the emission in radio halos were first suggested by Dennison 1980 and have since been discussed by a large number of authors. However, observational evidence for the presence of these particles has been lacking. There is, in fact, a growing body of evidence that, {\\em in general}, these particles may not be the underlying source of the emission in radio halos (Brunetti 2003; Kuo et al. 2003). Hence the EUV emission in the Coma cluster may be the only direct evidence for secondary electrons in an intracluster medium."
    },
    "0403/astro-ph0403429_arXiv.txt": {
        "abstract": "New high resolution CO~(3--2) interferometric map of the IR-bright interacting galaxy system VV~114 observed with the Submillimeter Array (SMA) reveal a substantial amount of warm and dense gas in the IR-bright but optically obscured galaxy, VV~114E, and the overlap region connecting the two nuclei.  A $1.8 \\times 1.4$ kpc concentration of CO~(3--2) emitting gas with a total mass of $4\\times 10^9 M_\\odot$ coincides with the peaks of NIR, MIR, and radio continuum emission found previously by others, identifying the dense fuel for the AGN and/or the starburst activity there.  Extensive CO~(2--1) emission is also detected, revealing detailed distribution and kinematics that are consistent with the earlier CO~(1--0) results.  The widely distributed molecular gas traced in CO~(2--1) and the distributed discrete peaks of CO~(3--2) emission suggest that a spatially extended intense starbursts may contribute significantly to its large IR luminosity.  These new observations further support the notion that VV~114 is approaching its final stage of merger, when violent central inflow of gas triggers intense starburst activity possibly boosting the IR luminosity above the ultraluminous threshold. ",
        "introduction": "Observations of luminous and ultraluminous infrared galaxies (LIRGs/ULIRGs) are the key to understanding the formation and evolution of galaxies and their associated star forming environment \\citep{sanders96,genzel00}.    Tidal interaction between two or more gas-rich progenitor galaxies is largely responsible for radially transporting the gas to the central kpc, condensing the gas and triggering subsequent starburst activity there \\citep{mihos96}.  It is now widely accepted that the elevated level of infrared luminosity originates from the reprocessed emission from the dust particles surrounding the starburst or AGN.  Based on a number of similarities, the infrared bright galaxies are thought to be the local analogs of the high redshift sub-mm sources discovered using SCUBA on JCMT during the last decade \\citep{blain02}, and understanding the nearby LIRGs/ULIRGs population is an important step toward better understanding the sub-mm galaxy phenomenon.  Investigation of LIRGs/ULIRGs using molecular gas at the low J-transitions (CO~(1--0) at 2.6 mm and CO~(2--1) at 1.3 mm) were carried out extensively in the past \\citep{bryant99,downes98}.  While these CO transitions were believed to be a good tracer of the optically thick dense gas identifying the extent of the starburst region and the distribution of the fuel for such activity, more recent studies suggest that the diffuse inter-clump medium may dominate the CO luminosity in these low J-transition CO emission \\citep{downes98}. Since these two J-transitions require only a small difference in excitation conditions, the large scale distribution of the molecular gas inferred is usually quite similar \\citep{downes98}.  The $J=3\\rightarrow 2$ transition of CO has a higher excitation temperature (33 K) and critical density ($\\sim 10^4$ cm$^{-3}$), making it a better tracer of the warmer and denser molecular gas of the starburst regions.  In this letter, we present  high resolution CO~(3--2) and CO~(2--1) interferometric maps of the IR-bright galaxy system VV~114 observed with 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} \\citep{ho04} to trace the distribution of the warmer dense gas and its relation to the colder, more diffuse molecular gas.  VV~114 is a gas rich \\citep[M$_{\\rm H_2} = 5.1 \\times 10^{10} \\rm M_{\\odot}$; ][YSK94 hereafter]{yun94} nearby ($\\rm D = 77$ Mpc) interacting system with high infrared luminosity \\citep[$\\rm L_{\\rm IR} = 4.0 \\times 10^{11} \\rm L_{\\odot}$; ][]{soifer87}.  The projected nuclear separation between the two optical galaxies (VV~114E and VV~114W) is $\\sim 6$ kpc.  \\citet{frayer99} found a large amount of dust ($\\rm M_{dust} = 1.2 \\times 10^8 \\rm M_{\\odot}$) distributed across the galaxy with dust temperature of $20 - 25$ K.  About half of the warmer dust traced in the MIR is associated with VV~114E, where both compact (nuclear region) and extended emission are found \\citep{lefloch02}.  The MIR spectrum also shows a sign of an AGN in VV~114E \\citep{lefloch02}.  \\citet{herrero02} detected abundant \\ion{H}{2} regions in VV~114E and in the overlap region using the narrow-band Pa$\\alpha$ images.  \\citet{scoville00} imaged the near infrared emission using NICMOS on-board HST and found that the highly optically obscured VV~114E \\citep{knop94} is the brighter of the two in the near infrared.  Far-UV imaging using STIS found several hundred young star clusters in VV~114W, while no UV emission was found in VV~114E \\citep{goldader02} which suggests that most of the activity in VV~114E is obscured by dust and not visible in the UV emission.  \\begin{deluxetable}{lccc} \\tablecaption{Observed and Derived Properties of VV~114\\label{tbl-1}} \\tablewidth{0pt} \\tablehead{ \\colhead{Parameter} & \\colhead{CO(3-2)} & \\colhead{CO(2-1)} & \\colhead{CO(1-0)\\tablenotemark{a}} } \\startdata $\\theta$ (FWHM)\\tablenotemark{b}\\\\ ~~arcsecond  & $2.9 \\times 2.2$ & $5.0 \\times 3.0$  &  $6.5 \\times 3.7$  \\\\ ~~kpc & $1.1 \\times 0.8$ & $1.9 \\times 1.2$ &  $2.4 \\times 2.2$ \\\\ Cont. RMS (mJy)\\\\ ~~Observed & 50 & 5 & -\\\\ ~~Theoretical & 20 & 5 & -\\\\ Peak (J2000)\\\\ ~~R.A.  & $01^{\\rm h} 07^{\\rm m} 47.5^{\\rm s}$ & $01^{\\rm h} 07^{\\rm m} 47.2^{\\rm s}$ &  $01^{\\rm h} 07^{\\rm m} 47.3^{\\rm s}$ \\\\ ~~Decl.  & $-17^{\\circ} 30' 25.0''$ & $-17^{\\circ} 30' 25.0''$ & $-17^{\\circ} 30' 25.7''$ \\\\ $S \\Delta v$ (Jy km/s)\\tablenotemark{c} & 2275 & 1620 & 674\\\\ $v_{sys}$ (km/s)&  6020 & 6027 & 6040\\\\ $\\Delta v_{FWZI}$ (km/s) & 330 & 520 & 520\\\\ \\enddata \\tablenotetext{a}{from \\citet{yun94}} \\tablenotetext{b}{Natural weighting} \\tablenotetext{c}{The integrated flux is subject  to a $20\\%$ (CO~(2--1)) uncertainty and factor of $\\sim 2$ overestimate (CO~(3--2)), mainly because of the uncertainties in flux and gain calibration.} \\end{deluxetable} ",
        "conclusions": "While several hundred young star clusters have been identified in VV~114W \\citep{goldader02}, most of the current vigorous activity is seen in the highly obscured galaxy, VV~114E, as identified by its infrared, radio and sub-mm continuum emission.  The origin of the intense IR emission, starburst or AGN or both, in VV~114 is still subject to debate.  While \\citet{lefloch02} suggests that 40\\% of the MIR emission is associated with the compact nuclear AGN in VV~114E, different diagnostics \\citep[i.e. NIR spectroscopy, ][]{doyon95} suggest that the IR emission is mostly dominated by young stars.  The three massive, compact gas complexes identified by our CO (3--2) observations mark the locations of the densest concentrations of gas and the sites of current activities fueled by this gas. The dominant peak C1 in VV~114E appears to mark the location of the most intense current activity.  The distributed nature of CO~(3--2) emission and the presence of several massive gas concentrations provide the clearest evidence yet that distributed intense starburst activity may provide a significant fraction of the large IR luminosity in VV~114.  Our new CO (3--2) and CO (2--1) observations strengthen the earlier suggestion that VV~114 is a late stage merger ready to undergo more massive bursts of star formation in the near future.  The widely extended CO~(2--1) and CO~(1--0) emission across the two galaxies implies that the cold and dense gas have begun to decouple from the gravitational potential of the host galaxies, and such decoupling only occurs in the well advanced stages ($t \\sim 5 \\times 10^8$ years since the initial collision) before the coalescence results in a massive elliptical galaxy \\citep[e.g., ][]{mihos96}.  The three discrete peaks traced in CO (3--2) emission identify the massive and compact molecular gas concentrations condensing out of the inflowing gas. By comparing the observed distribution of gas and optical morphology to those predicted in simulations \\citep{mihos96,iono04}, we estimate that the final coalescence may occur within the next $\\sim 10^8$ years. The CO~(1--0) and CO~(2--1) gas distribution of VV~114 is very similar to what is seen in the IR-bright galaxy NGC~6090 \\citep{bryant99, wang03}.  There the widely extended gas is seen over the edge-on and face-on system separated by 3.4 kpc, where not only the CO~(1--0) and CO~(2--1) but also the CO~(3--2) gas peaks in the overlap region \\citep{wang03}.  Thus, both VV~114 and NGC~6090 appears to be undergoing a similar transition where the bulk of the gas is in the process of decoupling from the host galaxies and funneling toward the dynamical center of the new combined potential, possibly resulting in the ultraluminous ($L_{IR}\\ge 10^{12} L_\\odot$) phase, following the scenario advocated by \\citet{yun94}.   The authors would like to thank J. Hibbard for kindly supplying the R-band image of VV~114.  D. I. and M.Y. are grateful for the warm and extensive support of the SMA group at Hilo during D.I.'s pre-doc residency at the SMA site, and the countless stimulating discussions with the members of the SMA.  This and other science observations were made possible only with the significant contributions from the staff members, engineers, operators and administrators at Hilo, Cambridge and ASIAA in Taiwan.   This research is partly supported by the Faculty Research Grant at the University of Massachusetts and the National Science Foundation grant AST 97-25951."
    },
    "0403/astro-ph0403103_arXiv.txt": {
        "abstract": "We present VLA HI observations at $\\sim$20$''$$\\simeq$1.5 kpc resolution of the highly inclined, HI-deficient Virgo cluster spiral galaxy NGC~4522, which is one of the clearest and nearest cases of ongoing ICM-ISM stripping. HI is abundant and spatially coincident with the stellar disk in the center, but beyond R = 3 kpc the HI distribution in the disk is sharply truncated and the only HI is extraplanar, and all on the northwest side. Forty percent of the total HI, corresponding to 1.5$\\times$10$^8$ M$_{\\sun}$, is extraplanar and has likely been removed from the galaxy disk by an ICM-ISM interaction. The kinematics and the morphology of the HI appear more consistent with ongoing stripping, and less consistent with gas fall-back which may occur long after peak pressure. Some of the extraplanar gas has linewidths (FWZI) of 150 km s$^{-1}$, including a blueshifted tail of weaker emission, and much of the extraplanar gas exhibits a modest net blueshift with respect to the galaxy's disk rotational velocities, consistent with gas accelerated toward the  mean cluster velocity. The SW side of the galaxy has less HI in the disk but more HI in the halo, suggesting more effective gas removal on the side of the galaxy which is rotating into the ICM wind. In recent simulations of ICM-ISM interactions, large surface densities of extraplanar gas like that in NGC~4522 are seen at relatively early stages of active stripping, and not during later gas fall-back stages. The galaxy is 3.3$\\deg$$\\simeq$800 kpc from M87, somewhat outside the region of strongest cluster X-ray emission. The ram pressure at this location, assuming a static smooth ICM and standard values for ICM density and galaxy velocity, appears inadequate to cause the observed stripping. We consider the possibility that the ram pressure is significantly stronger than standard values, due to large bulk motions and local density enhancements of the ICM gas, which may occur in a dynamic, shock-filled ICM experiencing subcluster merging. The HI and H$\\alpha$ distributions are similar, with both truncated in the disk at the same radius and HII regions located throughout much of the extraplanar HI. This implies that the star-forming molecular ISM has been effectively stripped from the outer disk of the galaxy along with the HI. The inferred peak stripping rate of $\\sim$10 M$_{\\sun}$ yr$^{-1}$ is much larger than the galaxy's total star formation rate of $\\sim$0.1 M$_{\\sun}$ yr$^{-1}$, implying that the rate of triggered star formation due to ICM pressure is presently minor compared to the rate of gas lost due to stripping. ",
        "introduction": "1. The Virgo cluster Sc galaxy NGC~4522 is a clear case of ISM gas selectively removed from an undisturbed stellar disk by an ICM-ISM interaction. The existence of a one-sided distribution of extraplanar gas and an undisturbed stellar disk clearly indicates a non-gravitational external process which has selectively removed ISM gas. HI is truncated within the disk at R=3 kpc = 0.4R$_{25}$. 40\\% of the HI appears to be extraplanar, and all the extraplanar gas is displaced toward one side of the major axis, extending 60$''$=5 kpc above the disk plane. The peak surface density of the extraplanar HI is 10 M$_{\\sun}$ pc$^{-2}$, a high value for extraplanar gas.  2. There is direct evidence for ongoing ICM pressure on NGC~4522 from radio continuum maps (Vollmer etal 2004), but it remains unclear whether the present ram pressure is sufficiently strong to to have caused the observed removal of the disk gas, or whether the galaxy may presently be observed some time after peak ICM pressure. Both the kinematics and the morphology of the HI are more suggestive of ongoing stripping than with gas fall-back which may occur long after peak pressure.  3. The HI content within the disk is asymmetric, with 50\\% less emission on the SW, high-velocity half of the disk. An opposite asymmetry is observed in the halo, which has more HI in the SW than the NE. This could be due to the stronger ICM pressure and enhanced stripping which is expected in the half of the disk which is rotating into the oncoming ICM.  4. Much of the extraplanar HI exhibits a modest net blueshift with respect to the systemic velocity of the galaxy, consistent with stripped galaxy gas being accelerated toward the mean cluster velocity. Extraplanar gas on the  SW, high-velocity half of the disk shows large linewidths (FWZI) of 150 km s$^{-1}$, including a tail of weaker emission extending to lower velocities. We suggest that the lower velocity tails represent gas which is accelerated toward the  mean cluster velocity.  5. The galaxy is 3.3$\\deg$$\\simeq$800 kpc from M87, which is somewhat outside the region of strongest X-ray emission, where the ICM gas density is highest. Assuming a smooth and static ICM, the calculated ram pressure at this location appears inadequate by an order of magnitude to cause the observed stripping. We consider the possibility that in a dynamic, shock-filled ICM with bulk motions and local density enhancements, the ram pressure presently on NGC~4522 may be significantly stronger than it would be in smooth, static ICM.      6. NGC~4522 has only $\\sim$25\\% of the HI expected for a galaxy of its mass, so $\\sim$75\\% , or 1.2$\\times$10$^9$ M$_{\\sun}$ of HI has presumably been stripped, heated, and ionized. The location and distribution of this gas is not known, but could in principle be traced by its X-ray emission, and if detected would provide evidence about the interaction history of NGC~4522.  7. The inferred peak gas stripping rate of $\\sim$10 M$_{\\sun}$ yr$^{-1}$ is much larger than the galaxy's total star formation rate of $\\sim$0.1 M$_{\\sun}$ yr$^{-1}$, implying that the present rate of triggered star formation due to ICM pressure is minor compared to the peak rate of gas loss due to ICM pressure.  8. The HI and H$\\alpha$ distributions in NGC~4522 are similar. Both are truncated at R=3 kpc in the disk, and they are coextensive in the extraplanar components. This implies that the star-forming ISM, presumably including the molecular gas, has been effectively removed from the galaxy beyond a radius of R=3 kpc. If this result is general, it would indicate that molecular gas does not strongly mitigate the effect of stripping on cluster galaxy evolution. The outer disks of cluster spirals can be stripped clean by ram pressure stripping. In clusters with a denser ICM and stronger ram pressure, most of the disk could be completely stripped, transforming spiral galaxies into lenticulars. ",
        "conclusions": "1. The Virgo cluster Sc galaxy NGC~4522 is a clear case of ISM gas selectively removed from an undisturbed stellar disk by an ICM-ISM interaction. The existence of a one-sided distribution of extraplanar gas and an undisturbed stellar disk clearly indicates a non-gravitational external process which has selectively removed ISM gas. HI is truncated within the disk at R=3 kpc = 0.4R$_{25}$. 40\\% of the HI appears to be extraplanar, and all the extraplanar gas is displaced toward one side of the major axis, extending 60$''$=5 kpc above the disk plane. The peak surface density of the extraplanar HI is 10 M$_{\\sun}$ pc$^{-2}$, a high value for extraplanar gas.  2. There is direct evidence for ongoing ICM pressure on NGC~4522 from radio continuum maps (Vollmer etal 2004), but it remains unclear whether the present ram pressure is sufficiently strong to to have caused the observed removal of the disk gas, or whether the galaxy may presently be observed some time after peak ICM pressure. Both the kinematics and the morphology of the HI are more suggestive of ongoing stripping than with gas fall-back which may occur long after peak pressure.  3. The HI content within the disk is asymmetric, with 50\\% less emission on the SW, high-velocity half of the disk. An opposite asymmetry is observed in the halo, which has more HI in the SW than the NE. This could be due to the stronger ICM pressure and enhanced stripping which is expected in the half of the disk which is rotating into the oncoming ICM.  4. Much of the extraplanar HI exhibits a modest net blueshift with respect to the systemic velocity of the galaxy, consistent with stripped galaxy gas being accelerated toward the mean cluster velocity. Extraplanar gas on the  SW, high-velocity half of the disk shows large linewidths (FWZI) of 150 km s$^{-1}$, including a tail of weaker emission extending to lower velocities. We suggest that the lower velocity tails represent gas which is accelerated toward the  mean cluster velocity.  5. The galaxy is 3.3$\\deg$$\\simeq$800 kpc from M87, which is somewhat outside the region of strongest X-ray emission, where the ICM gas density is highest. Assuming a smooth and static ICM, the calculated ram pressure at this location appears inadequate by an order of magnitude to cause the observed stripping. We consider the possibility that in a dynamic, shock-filled ICM with bulk motions and local density enhancements, the ram pressure presently on NGC~4522 may be significantly stronger than it would be in smooth, static ICM.      6. NGC~4522 has only $\\sim$25\\% of the HI expected for a galaxy of its mass, so $\\sim$75\\% , or 1.2$\\times$10$^9$ M$_{\\sun}$ of HI has presumably been stripped, heated, and ionized. The location and distribution of this gas is not known, but could in principle be traced by its X-ray emission, and if detected would provide evidence about the interaction history of NGC~4522.  7. The inferred peak gas stripping rate of $\\sim$10 M$_{\\sun}$ yr$^{-1}$ is much larger than the galaxy's total star formation rate of $\\sim$0.1 M$_{\\sun}$ yr$^{-1}$, implying that the present rate of triggered star formation due to ICM pressure is minor compared to the peak rate of gas loss due to ICM pressure.  8. The HI and H$\\alpha$ distributions in NGC~4522 are similar. Both are truncated at R=3 kpc in the disk, and they are coextensive in the extraplanar components. This implies that the star-forming ISM, presumably including the molecular gas, has been effectively removed from the galaxy beyond a radius of R=3 kpc. If this result is general, it would indicate that molecular gas does not strongly mitigate the effect of stripping on cluster galaxy evolution. The outer disks of cluster spirals can be stripped clean by ram pressure stripping. In clusters with a denser ICM and stronger ram pressure, most of the disk could be completely stripped, transforming spiral galaxies into lenticulars."
    },
    "0403/astro-ph0403335_arXiv.txt": {
        "abstract": "\\bigskip  We propose a new mechanism for magnetic field generation from inflation, by which strong magnetic fields can be generated on cosmological scales. These fields may be observable by cosmic microwave background radition measurements, and may have a dynamical impact on structure formation. The mechanism is based on the observation that a light nearly minimally coupled charged scalar may be responsible for the creation of a negative photon mass-squared (provided the scalar field coupling to the curvature scalar is small but negative), which in turn results in abundant photon production -- and thus in growing magnetic fields -- during inflation. ",
        "introduction": "\\label{Cosmological magnetic fields}  There is a growing evidence that magnetic fields permeate intergalactic medium~\\cite{Widrow:2002,Kronberg:1993,CarilliTaylor:2001,VogtEnsslin:2003}. The galactic $\\alpha-\\Omega$ dynamo mechanism~\\cite{Parker:1971} represents the standard explanation, according to which small seed fields in protogalaxies are magnified to a microgauss strength correlated on kiloparsec scales observed today in many galaxies. Yet the dynamo does poorly as regards to explaning the observed cluster fields, which are typically correlated over several kiloparsecs and can reach $10$ microgauss strength. On larger scales the field strength drops quite dramatically, and it is characterized by a negative spectral index between $-1.6$ and $- 2$~\\cite{VogtEnsslin:2003}, consistent with the Kolmogorov spectrum of turbulence.  The seed field that fuels the dynamo, which amplifies the galactic (and possibly cluster) fields, has either primordial origin, or it was created by the Biermann battery mechanism, operative at the time of structure and galaxy formation, when the Universe was about one billion years old~\\cite{PudritzSilk:1989,KulsrudCenOstrikerRyu:1996,DaviesWidrow:2000}. The Biermann battery mechanism is operative in the presence of oblique shocks, which generate nonideal fluid that violates proportionality between the gradients of pressure and energy density. Such a fluid can produce the vorticity required for generation of seed fields, whose strength is typically of the order $10^{-20}$~gauss. While in the protogalactic environments these seed fields may be sufficiently strong to fuel the galactic dynamo mechanism, whether they can be used to explain the observed cluster fields is a highly controversial question. Moreover, these fields may be insufficiently strong to explain the galactic fields observed in a couple of galaxies at higher redshifts~\\cite{Kronberg:1993,Widrow:2002}.   The quest for the fields correlated on much larger scales has so far been unsuccessful. A notable exception is the result of~\\cite{AnchordoquiGoldberg:2001}, where a local field of about $10^{-8}$ gauss correlated on  megaparsec scale is quoted. If established, the existence of extragalactic fields would point at a primordial origin. Perhaps the most promising method for detecting primordial fields correlated over supragalactic scales is through their impact on the cosmic microwave background radiation (CMBR)~\\cite{BarrowFerreiraSilk:1997,DurrerFerreiraKahniashvili:1999}, and possibly through their dynamical influence on large scale structure and galaxy formation. The current upper bounds from the CMBR are of the order $10^{-9}$~gauss, while the bounds from structure formation are somewhat weaker. Fields stronger than about $10^{-5}$ gauss at the time of galaxy formation (which correspond to about $10^{-9}$~gauss comoving field strength) are dynamically relevant.  In this Letter we propose a novel mechanism for generation of large scale magnetic fields during cosmological inflation. Our mechanism is based on a model with a rather minimal set of assumptions: We study the dynamics of (one-loop) quantum  scalar electrodynamics~\\cite{ProkopecPuchwein:2003,ProkopecTornkvistWoodard:2002AoP, ProkopecTornkvistWoodard:2002prl,ProkopecWoodard:2003ajp, ProkopecWoodard:2003photons} in cosmological space-times, which can be naturally embedded into the Standard Model (where the role of the scalars is played by the charged components of the Higgs field), or into supersymmetric extensions of the Standard Model (the charged scalars are the sypersymmetric partners of the Standard Model fermions). ",
        "conclusions": "\\label{Discussion}  The spectrum~(\\ref{spectrum:radiation}) can be thought of as a function of the scalar coupling to gravity $\\xi$ and its mass $m_\\Phi$, which are fundamental parameters characterizing a scalar field. To illustrate this point more precisely, we show the spectral index $n = 3 - 2\\nu$ characterizing the spectrum envelope of ${\\cal P}_B\\propto k^n$ in Eq.~(\\ref{spectrum:radiation}) on subhorizon scales, $k/a > H(t)$ (or $B_\\ell \\equiv \\langle\\vec B_\\ell^2\\rangle^{1/2} \\propto \\ell^{-n/2}$ ($n\\neq 0$), $B_\\ell \\propto -\\ln(\\ell)$ ($n = 0$) in Eq.~(\\ref{spectrum:definition})). Between $\\xi = \\xi_{\\rm crit} \\equiv - m_\\Phi^2/(12H^2)$  and $\\xi = \\xi_3 \\equiv (\\alpha/\\pi) - m_\\Phi^2/(12H^2)$, $n=3$, equal to that of thermal spectrum. Above $\\xi > \\xi_3$, $n$ drops, reaching an asymptotic value $n\\rightarrow n_{\\infty} = 2$ when $\\xi \\rightarrow \\infty$. The spectral index on superhorizon scales is obtained simply by replacing $n + 2 \\rightarrow n = 5-2\\nu$. When compared with the vacuum spectrum,  ${\\cal P}_{B}^{\\rm vac} \\propto k^4$ ($k<H$), a spectrum with $n \\in (2,3)$ exhibits an enhancement of magnetic fields on subhorizon scales, which is due to a conversion of the electric energy (which is enhanced during inflation) into the magnetic energy during radiation era. The conversion is efficient provided radiation era is characterised by a low conductivity. This mechanism was used in Refs.~\\cite{ProkopecPuchwein:2003,ProkopecTornkvistWoodard:2002AoP, ProkopecTornkvistWoodard:2002prl,ProkopecWoodard:2003ajp, ProkopecWoodard:2003photons,DavisDimopoulosProkopecTornkvist:2000, DimopoulosProkopecTornkvistDavis:2001} to argue that a spectrum $B_\\ell \\propto \\ell^{-1}$ can be obtained from inflation, which could be sufficient to seed the galactic dynamo mechanism.  When $\\xi\\in(-\\infty,\\xi_{\\rm crit})$ however, gauge fields exhibit instability and are enhanced during inflation, such that the spectral index of subhorizon modes drops from $n=2$ when $\\xi \\rightarrow -\\infty$, to flat spectrum ($n=0$), when \\begin{equation} \\xi_{\\rm flat} = -\\frac{\\alpha}{8\\pi}-\\frac{m_\\Phi^2}{12H^2} \\,, \\label{xi flat} \\end{equation} to $n<0$, when $\\xi \\in (\\xi_{\\rm flat},\\xi_{\\rm crit})$.  % A negative spectral index $n \\leq - 2$ would imply a growth in magnetic field energy during inflation, resulting in a  divergent magnetic (and electric) field energy, as it can be inferred from Eqs.~(\\ref{spectrum:radiation}), (\\ref{spectrum:large sigma}) and~(\\ref{spectrum:definition}), with ${\\cal W} = 1$. A proper study of this case would require an inclusion of backreaction of the electromagnetic field on the background space-time, which is beyond the scope of this Letter. A spectral index $n > -2$ on subhorizon scales (or $n > 0$ on superhorizon scales), results in an acceptable magnetic field spectrum on cosmological scales, whose dynamical impact on CMBR and large scale structure formation should be considered. \\begin{figure}[h] \\begin{center} \\hskip -0.05in \\includegraphics[scale=1]{spectral-indexB.eps} \\end{center} \\vskip -0.25in \\caption{Spectral index $n$ for subhorizon modes in the low conductivity case ({\\it cf.} Eq.~(\\ref{spectrum:radiation})) as a function of the scalar coupling to gravity $\\xi$ (with $\\alpha = 1/137$ and the scalar mass $m_\\Phi^2 = 10^{-3}H^2$). } \\label{fig3:spectral index} \\end{figure}   Spectrum normalization today can be determined from magnetic field strength at a comoving scale corresponding to the Hubble scale at the end of inflation ($k=H$). At the comoving inflation scale $\\ell_H\\sim 3~{\\rm m} \\sim 10^{-16}$ parsec (corresponding to $H \\sim 10^{37}~{\\rm Hz} \\sim 10^{13}~{\\rm GeV}$) the field strength is about $B_H \\sim 10^{-12}-10^{-11}$~gauss~\\cite{Prokopec:2001}. Therefore, if the spectrum is (almost) flat, the magnetic field is potentially observable by the next generation of CMBR experiments (PLANCK satellite), as well as of importance for the dynamics of large scale structures of the Universe.  \\bigskip"
    },
    "0403/astro-ph0403045_arXiv.txt": {
        "abstract": "We show that nonlinear electrodynamics (NLED) modifies in a fundamental basis the concept of gravitational red-shift (GRS) as it was introduced by Einstein's general relativity. The effect becomes apparent when light propagation from super strongly magnetized compact objects, as pulsars, is under focus. The analysis, here based on the (exact) nonlinear Lagrangean of Born \\& Infeld (1934), proves that alike from general relativity (GR); where the GRS independs on any background magnetic ($B$) field, when NLED is taken into the photon dynamics an effective GRS appears, which happens to decidedly depend on the $B$-field pervading the pulsar. The resulting GRS tends  to infinity as the $B$-field grows larger, as opposed to the Einstein prediction. As in astrophysics the GRS is admittedly used to infer the mass-radius relation, and thus the equation of state of a compact star, e.g. a neutron star (Cottam, Paerels \\& Mendez 2002), this unexpected GRS critical change may misguide observers to that fundamental property. Hence, a correct procedure to estimate those valuous physical properties demands a neat separation of the NLED effects from the pure gravitational ones in the light emitted by ultra magnetised pulsars. ",
        "introduction": "The idea that the nonlinear electromagnetic interaction, i. e., light propagation in vacuum, can be geometrized was developed by Novello et al. (2000) and Novello \\& Salim (2001). Since then a number of physical consequences for the dynamics of a variety of systems have been explored. In a recent paper, Mosquera Cuesta \\& Salim (2003,2004) presented the first astrophysical context where such nonlinear electrodynamics effects were accounted for: the case of a highly magnetized neutron star or pulsar. In that paper NLED was invoked {\\it a l\\`a} Euler-Heisenberg, which is an infinite series expansion of which only the first nonlinear term was used for the analisys. An immediate consequence of that study was an overall modification of the space-time geometry around the pulsar as ``perceived'' by light propagating out of it. This translates into a fundamental change of the star surface redshift, the GRS, which might have been inferred from the absorption (or emission) lines observed in a super magnetized pulsar by Ibrahim et {al.} (2002;2003). The result proved to be even more dramatic for the so-called magnetars: pulsars said to be endowed with magnetic ($B$) fields higher then the Schafroth quantum electrodynamics critical $B$-field. In this {\\it Letter} we demonstrate that the same effect still appears if one calls for the NLED in the form of the one rigorously derived by Born \\& Infeld (1934), which is based on the special relativistic limit to the velocity of approaching of an elementary charged particle to a pointlike electron. As compared to our previous results, here we stress that from the mathematical point of view the Born \\& Infeld (1934) NLED is described by an exact Lagrangean whose dynamics has been successfully studied in a wide set of physical systems (Delphenich 2003). The analysis presented next proves that this physics affects not only the magnetar electrodynamics (our focus here) but also that one of newly born; highly magnetised proto-neutron stars and the dynamics of their progenitor supernovae. ",
        "conclusions": "In a very interesting couple of {\\it Letters} by Ibrahim et {al.} (2002;2003) was reported the discovery of cyclotron line resonance features in the source SGR 1806-20, said to be a candidate to a magnetar by Kouveliotou {et al.} (1998).  The 5.0 keV feature discovered with {\\it Rossi XTE} is strong, with an equivalent width of $\\sim 500$~eV and a narrow width of less than 0.4~eV (Ibrahim et {al.} 2002;2003). When these features are interpreted in the context of accretion models one arrives to a $M/R > 0.3$ M$_\\odot$ km$^{-1}$, which is inconsistent with NSs, or that requires a low $(5-7) \\times 10^{11}$~G magnetic field, which is said not to correspond to any SGRs (Ibrahim et {al.} 2003). In the magnetar scenario, meanwhile, the features are plausibly explanined as being ion-cyclotron resonances in an ultra strong $B$-field $B_{\\rm sc} \\sim 10^{15}$~G (see Eq.(\\ref{line})).  The spectral line is said to be consistent with a proton-cyclotron fundamental state whose energy and width are close to model predictions (Ibrahim et { al.} 2003). According to Ibrahim et {al.} (2003), the confirmation of this findings would allow to estimate the gravitational redshift (the GRS), mass ($M$) and radius ($R$) of the quoted ``magnetar'': SGR 1806-20.   Although the quoted spectral line in Ibrahim et {al.} (2002;2003) is found to be a cyclotron resonance produced by protons in that high field, we alert on the possibility that it could also be due to NLED effects in the same super strong magnetic field of SGR 1806-20, as suggested by Eq.(\\ref{tt-component}).  If this were the case, no conclusive assertion on the $M/R$ ratio of the compact star glowing in SGR 1806-20 could be consistently made, since a NLED effect might well be emulating the standard gravitational one associated with the pulsar surface redshift.  In summary, because we started with a more general and exact Lagrangean, as that of Born \\& Infeld (1934), we can assert that the result here derived is inherently generic to any kind of nonlinear theory describing the electromagnetic interaction, and thus is universal in nature. Thence, absorption or emission lines from magnetars, if these stellar objects do exist in nature (see P\\'erez Mart\\'{\\i}nez et al. 2003 for arguments contending their formation), cannot be saftely used as an unbiased source of information regarding the fundamental parameters of a NS pulsar such as its $M$, $R$ or EOS."
    },
    "0403/astro-ph0403273_arXiv.txt": {
        "abstract": "The high resolution obtained through the use of VLBI gives an unique opportunity to directly observe the interaction of an expanding radio supernova with its surrounding medium. We present here results from our VLBI observations of the young supernovae SN~1979C, SN~1986J, and SN~2001gd. \\footnote{MAPT, AA, and ER thank the organizers for funding to attend the conference. This research has been supported by the Spanish DGICYT grants AYA2001-2147-C02-01 and AYA2001-2147-CO2-02. KWW wishes to thank the Office of Naval Research for the 6.1 funding supporting this research. } ",
        "introduction": "\\label{sec:intro}  In the standard model of radio emission from supernovae, a blast wave is driven into an ionized, dense, slowly expanding wind. As a result, a high-energy-density shell is formed. The relativistic electrons present in this shell spiral along the magnetic field and respond for the observed radio synchrotron emission. The supernova quickly increases its radio brightness with time, due to the increasingly smaller electron column density in the line of sight. When the optical depth at cm-wavelengths is about unity, the supernova reaches its maximum of emission, after which the emission monotonically decreases due to expansion losses. Very-Long-Baseline Interferometry (VLBI) observations of radio supernovae are a powerful tool to probe the circumstellar interaction that takes place after a supernova explodes. Indeed, high-resolution radio observations permit us to trace the presupernova mass loss history by directly imaging the structure of the supernova as it expands. The wealth of information includes also a direct estimate of the deceleration of the supernova expansion, estimates of the ejecta and circumstellar density profiles (which has implications on the progenitor system), distortion of the shock front, and the potential observation of Rayleigh-Taylor instabilities. Unfortunately, the usually large distances to radio supernovae makes their VLBI imaging a challenging task, and only radio supernovae that are young, bright, nearby, and rapidly expanding are appropriate VLBI targets. Since high-resolution radio observations of the young supernovae 1993J and 1987A, as well as VLBI observations of the supernova remnants in M82 are covered by other authors in these proceedings, we shall limit our discussion to the cases of SN~1979C, SN~1986J, and SN~2001gd. ",
        "conclusions": ""
    },
    "0403/astro-ph0403209_arXiv.txt": {
        "abstract": "A combination of evidence is presented suggesting that the majority of the stars in today's galaxies were born during a luminous infrared phase (LIRP) triggered by the local environment of galaxies. The CIRB is a fossil record of these LIRPs and therefore reflects the influence of triggered star formation through galaxy-galaxy interactions, including non merging tidal encounters.  This scenario, in which galaxies experienced several LIRPs in their history, is consistent with the measured redshift evolution of the cosmic density of star formation rates and of stellar masses of galaxies. ",
        "introduction": "Stars represent only 15\\,\\% of the cosmic baryonic density, itself only equal to about 4\\% of the critical density ($\\Omega_{\\rm b}\\simeq$ 0.04), and are unequally distributed into spheroids (10\\% of $\\Omega_{\\rm b}$, including spiral bulges) and disks (5\\% of $\\Omega_{\\rm b}$, Fukugita, Hogan \\& Peebles). It is usually assumed that disk stars formed quiescently while bulge stars formed more efficiently and rapidly, as suggested by their redder colors and overabundance in $\\alpha$-element over iron ratio, typical of a SNII origin. However recent studies of the history of the star formation of the disk of the Milky Way indicate that during the last 2 Gyr it has experienced about five major events of star formation, starbursts, whose signatures can be found in the peaked ages of these open clusters (de la Fuente Marcos \\& de la Fuente Marcos 2004, and references therein). As a result we may wonder whether quiescent star formation did play a major role in the formation of present-day stars at all. There are some evidence that star formation mostly takes place in globular clusters and is rarely isolated. These clusters are thought to evolve into unbound stellar associations, which evolve and dissolve in a time-scale of about 50 Myr. This timescale is longer than the lifetime of massive stars which dominate the luminosity of starbursting galaxies or regions of galaxies. Hence, it is logical to expect that if star formation occured mainly in starburst episodes within galaxies, then the bulk of galaxies luminosity will be absorbed by dust in the giant molecular clouds, while if star formation is quiescent then only a minor fraction of a galaxy's luminosity will be affected by dust extinction. We will argue in the following that there is presently a solid collection of evidence suggesting that most stars that we see in the local universe formed during starburst episodes.  A major peace of evidence for that comes from the detection of a strong diffuse cosmic infrared background (CIRB, Puget et al. 1996, Hauser \\& Dwek and references therein) which is majoritarily produced by luminous infrared phases (LIRP) within galaxies located around $z\\sim$ 0.7, for the peak of the CIRB at $\\lambda \\sim$ 140\\,$\\mu$m, while the $\\lambda \\geq$ 240 \\,$\\mu$m is due to galaxies at $z\\sim$ 2 and above (Elbaz et al. 2002, Chary \\& Elbaz 2001). We have introduced the term LIRP instead of the classical one, luminous and ultra-luminous infrared galaxies, i.e. LIRGs and ULIRGs, because we wish to emphasize the idea that the scenario that is emerging from the study of distant galaxies is that LIRGs do not represent a population of galaxies that would require to be studied independantly in order to determine which present-day galaxies are the remnants of these LIRGs, but what is suggested instead is that they illustrate the omnipresence of rapid and efficient star formation as a leading process in shaping the present-day universe, i.e. that any galaxy that we see today must have experienced a phase when it radiated the bulk of its light in the infrared (see also Elbaz \\& Cesarsky 2003). This phase should not be restricted to LIRGs and ULIRGs, i.e. galaxies with infrared (IR) luminosities larger than 10$^{11}$ L$_{\\odot}$ or star formation rates (SFR) larger than $\\sim$ 20 M$_{\\odot}$ yr$^{-1}$, since the closest starburst M82, for example, presents a spectral energy distribution typical of most LIRGs, with the bulk of its luminosity radiated in the IR although its luminosity is only 4$\\times$10$^{10}$ L$_{\\odot}$.  We present evidence suggesting that the bulk of local stars formed during a LIRP. A spectroscopic diagnostic is used to quantify the typical duration of this phase and the amount of stars that formed during it. From the combination of both we will advocate that not only did all galaxies experience a LIRP but that they must have experienced several of them. Finally we will discuss the physical origin of the LIRP and present evidence that a major event in the lifetime of galaxies was probably underestimated: the effect of the ``local environment of galaxies'' (LEG) and its impact in terms of driving the conversion of gas into stars through passing-by galaxies. ",
        "conclusions": ""
    },
    "0403/astro-ph0403453_arXiv.txt": {
        "abstract": "We report the discovery of the binary millisecond pulsar J0514$-$4002A, which is the first known pulsar in the globular cluster NGC~1851 and the first pulsar discovered using the Giant Metrewave Radio Telescope (GMRT). The pulsar has a rotational period of 4.99\\,ms, an orbital period of 18.8\\,days, and the most eccentric pulsar orbit yet measured ($e = 0.89$). The companion has a minimum mass of 0.9\\,M$_{\\odot}$ and its nature is presently unclear. After accreting matter from a low-mass companion star which spun it up to a (few) millisecond spin period, the pulsar eventually exchanged the low-mass star for its more massive present companion.  This is exactly the same process that could form a system containing a millisecond pulsar and a black hole; the discovery of NGC~1851A demonstrates that such systems might exist in the Universe, provided that stellar mass black holes exist in globular clusters. ",
        "introduction": "\\label{sec:intro}  Since 1987, several globular cluster surveys at Jodrell Bank, Arecibo, Parkes and more recently the Green Bank telescope have found 79 pulsars in a total of 23 globular clusters (GCs)\\footnote{See http://www2.naic.edu/$\\sim$pfreire/GCpsr.html and references therein}.  These findings have confirmed that most of the binary millisecond pulsars (MSPs) in GCs have low-mass white dwarf (WD) companions and nearly circular orbits, as observed in the Galactic disk. This is an important confirmation of the evolutionary scenarios proposed by Alpar~et~al.~(1982)\\nocite{acrs82} for the formation of MSPs.  In GCs, {\\em exchange encounters}, which only have a significant probability of occurring in dense stellar environments, occasionally exchange one of the components of a binary system with a typically more massive star.  The exchanges may occur during encounters with either other binaries or with isolated stars.  In GCs such encounters can place isolated neutron stars into binaries with a main sequence (MS) star which eventually evolves, ``recycles'' the neutron star, and finally forms a MSP$-$WD binary system.  Such a process explains the anomalously large numbers of MSPs in GCs (by mass) when compared to the Galaxy. If the neutron star placed in orbit around a low-mass MS star is already a pulsar, we observe ``irregular'' eclipsing binary pulsars, such as PSR~J0024$-$7204W in 47~Tucanae \\cite{egc+02}, PSR~B1718$-$19 in NGC~6342 \\cite{lbhb93}, and PSR~J1740$-$5340 in NGC~6397 \\cite{fpds01}, which are peculiar to GCs.  In this paper we report and discuss the discovery of a unique binary millisecond pulsar, PSR~J0514$-$4002A (henceforth NGC~1851A) in the GC NGC~1851. We discovered this system in a new 327-MHz survey of GCs carried out using the Giant Metrewave Radio Telescope (GMRT), at Khodad near Pune, India. ",
        "conclusions": "We have discovered a remarkable 5-ms binary pulsar, the first to be found in the GC NGC~1851 and the first pulsar to be discovered with the GMRT. Its orbit is the most eccentric known for any system containing a pulsar, while its rotational period is much shorter than that of any other pulsar in an eccentric binary system.  This indicates that, after becoming an MSP by accreting matter from a low-mass companion star, this neutron star has almost certainly exchanged it for its present, significantly more massive companion. If black holes exist in GCs, a MSP-black hole binary could be formed in exactly the same way.  If the companion is a compact object, then two relativistic effects will be measurable, namely the rate of advance of periastron and the Einstein delay. The measurement of both effects would lead to the determination of the masses of the two components of this system. A third relativistic effect, the orbital decay due to emission of gravitational waves, will be masked by the unpredictable acceleration of the binary in the gravitational field of the cluster, so no tests of general relativity are likely to be possible. If the companion is a MS star, we might measure the orbital precession caused by the quadrupole moment of the companion star and possibly a rotation of the orbital plane around the axis of rotation of the companion \\cite{wjm+98}.  In addition, variations in DM with orbital phase may be detectable \\cite{jml+96}, particularly if the companion is an active star. For NGC~1851A none of these effects have been conclusively measured to date."
    },
    "0403/astro-ph0403386_arXiv.txt": {
        "abstract": "We have numerically explored the stable planetary geometry for the multiple systems involved in a 2:1 mean motion resonance, and herein we mainly concentrate on the study of the HD 82943 system by employing two sets of the orbital parameters (Mayor et al. 2004). In the simulations, we find that all stable orbits are related to the 2:1 commensurability that can help to remain the semi-major axes for two companions almost unaltered over the secular evolution for $10^{7}$ yr, and the apsidal phase-locking between two orbits can further enhance the stability for this system, because the eccentricities are simultaneously preserved to restrain the planets from frequent close encounters. For HD 82943, there exist three possible stable configurations:(1) Type I, only $\\theta_{1} \\approx 0^{\\circ}$, (2) Type II, $\\theta_{1}\\approx\\theta_{2}\\approx\\theta_{3}\\approx 0^{\\circ}$ (aligned case), and (3) Type III, $\\theta_{1}\\approx 180^{\\circ}$, $\\theta_{2}\\approx0^{\\circ}$, $\\theta_{3}\\approx180^{\\circ}$ (antialigned case), here the lowest eccentricity-type mean motion resonant arguments are $\\theta_{1} = \\lambda _{1} - 2\\lambda _{2} + \\varpi_{1}$ and $\\theta_{2} = \\lambda _{1} - 2\\lambda _{2} + \\varpi_{2}$, the relative apsidal longitudes  $\\theta_{3} = \\varpi_{1}-\\varpi_{2}=\\Delta\\varpi$ (where $\\lambda _{1,2} $ are, respectively, the mean longitudes of the inner and outer planets;$\\varpi_{1,2}$ are the longitudes of periapse). And we find that the other 2:1 resonant systems (e.g., GJ 876 or HD 160691) may possess one of three stable orbits in their realistic motions. In addition, we also propose a semi-analytical model to study $e_{i}-\\Delta\\varpi$ Hamiltonian contours, which are fairly consistent with direct numerical integrations. With the updated fit, we then examine the dependence of the stability of this system on the relative inclination, the planetary mass ratios, the eccentricities and other orbital parameters: in the non-coplanar cases, we find that stability requires the relative inclination being $\\sim 25^{\\circ}$ or less; as to the planetary mass ratio, the stable orbits for HD 82943 requires $\\sin i \\geq 0.50$ for a fixed  value or $m_{1}/m_{2}\\leq 2$ where $m_{2}$ remains for the varying mass ratio; concerning the eccentricities, the system can be always steady when $0 < e_{2}\\le 0.24$ and $0 < e_{1}< 0.60$. Moreover, we numerically show that the assumed terrestrial bodies cannot survive near the habitable zones for HD 82943 due to the strong perturbations induced by two resonant companions, but these low-mass planets can be dynamically habitable in the GJ 876 system at $\\sim 1$ AU in the numerical surveys. Finally, we present a brief discussion on the origin of the 2:1 resonance for HD 82943. ",
        "introduction": "The discovery of the extrasolar planets is opening a new world beyond our solar system. Ever since 1995, Mayor \\& Queloz (1995) detected the first extrasolar giant Jupiter--51 Peg, and to date there are more than 100 planetary systems discovered \\footnote{see also http://cfa-www.harvard.edu/planets/bibli.html and http://exoplanets.org/}(Marcy, Cochran, \\& Mayor 2000; Butler et al. 2003) using the radial velocity technique in the surveys of nearby young stars. At the time of writing, a dozen of multiple planet systems--HD 82943, GJ 876, HD 168443, HD 74156, 47 Uma, HD 37124, HD 38529, HD 12661, 55 Cnc, $\\upsilon$ And (Fischer et al. 2003), HD 169830 (Mayor et al. 2004) and HD 160691 (Jones et al. 2002) were discovered in recent years and this number is undoubtedly cumulative as the measurements are being carried on. Hence, it is necessary to categorize the discovered multiple planetary systems according to their statistical characteristics (such as the distribution of the planetary masses, semi-major axes, eccentricities and metallicity) (Marcy et al. 2003), then to study the correlation between mass ratio and period ratio (Mazeh \\& Zucker 2003) and further to improve the understanding of the relationship between the planet occurrence rate and stellar metallicity (Santos et al. 2003; Fischer, Valenti, \\& Marcy 2004). The other key point is to investigate possible stable configurations for the multiple systems, in which the observations reveal that most of them are typically characterized by mean motion resonance (MMR) and (or) apsidal phase-locking between their orbiting companions (Fischer et al. 2003; Ji et al. 2003a; Lee \\& Peale 2003), so that one can better understand the full dynamics of these systems. In the present study, we mainly focus our attention on the HD 82943 system.  As of April 4, 2001, the Geneva Extrasolar Planet Search Team \\footnote {see http://obswww.unige.ch/$\\sim$udry/planet/hd82943syst.html} announced the discovery of the HD 82943 planetary system at an ESO press release, consisting of two Jupiter-like planets orbiting the parent star. Subsequently, Israelian et al. (2001, 2003) reported the discovery of $^6$Li in the atmosphere of this metal-rich solar-type star and indicated that the presence of $^6$Li can probably be shown as evidence for a planet or planets having been engulfed by the star HD82943. At first, let us review several stellar features for this star:HD 82943 is a G0 star with $B-V=0.623$, Hipparcos parallax of 36.42 mas, distance of 27.46 pc, and [Fe/H] = $0.29 \\pm 0.02$ (Santos et al. 2001, 2003). The star is aged 2.9 Gyr with the stellar mass of 1.15 $M_{\\odot}$. Similar to the GJ 876 system (Marcy et al. 2001), which is believed that stability of the planetary system might be sustained by the 2:1 mean-motion resonance (MMR) and apsidal alignment in the semi-major axes (Kinoshita \\& Nakai 2001; Lee \\& Peale 2002; Ji, Li, \\& Liu 2002), the two planets of HD 82943 are now also close to a 2:1 commensurability, with orbital periods $435.1 \\pm 1.4$ and $219.4\\pm 0.2$ \\,d, and semi-major axes 1.18 and 0.75\\,AU (Mayor et al. 2004). However, the previous investigations (Gozdziewski \\& Maciejewski 2001) of the stability of the HD 82943 system showed that the system given by the earlier best-fit orbital solution was extremely unstable  due to the strong interaction of two massive planets with high eccentricities, hence a natural question is that whether the two companions for this system are really in a 2:1 orbital resonance and whether the system is secularly stable with the observed best-fit orbital parameters. Additionally, the stability for HD 82943 places constraints on both planetary masses--if the mass of the outer companion is smaller than that of the inner planet (Hadjidemetriou 2002), a system that harbors two planets moving on elliptic 2:1 resonant periodic orbits will be destabilized as a finale. Nevertheless, the recently updated orbital solutions (Mayor et al. 2004), with more abundant observations as of October 13, 2003, show that the inner planet is a bit more massive than the outer one, thus it is indeed noteworthy to thoroughly reanalyze the observed HD 82943 system, still these mentioned issues of this fascinating system are to arouse our great interests and inspire us to seek after the answers. Consequently, our first goal is to carefully explore the system to fully understand in what kind of likely configurations that two companions may orbit about their host star and if exist, to reveal the dynamical mechanisms to remain the system.  In a viewpoint of celestial mechanics, the former studies are concerned for the dynamical analysis for the global stability of the exosystems (Gozdziewski \\& Maciejewski 2001, 2003; Ji et al. 2002; Kiseleva-Eggleton et al. 2002 ; Barnes \\& Quinn 2004), the investigations of resonant picture in the phase space (Hadjidemetriou 2002; Ji et al. 2003a; Haghighipour et al. 2003; Callegari Jr., Michtchenko, \\& Ferraz-Mello 2004) or chaotic behavior (Jiang \\& Yeh 2004) in these systems, and the existence of Habitable Zones (HZ) for Earth-like planets that supply stable liquid-water environment to cultivate extraterrestrial intelligent beings (Kasting, Whitmore, \\& Reynolds 1993; Jones \\& Sleep 2002; Laughlin, Chambers, \\& Fischer 2002; Dvorak et al. 2003; Menou \\& Tabachnik 2003). Other recent works on the 2:1 resonant geometry of the extrasolar planetary systems include those by Hadjidemetriou \\& Psychoyos (2003), Beauge, Ferraz-Mello, \\& Michtchenko (2003) and Lee (2004). They studied the presence and location of stable equilibrium solutions or dynamical evolution of such exosystems in more general way. In this work, we revisit the dynamics of the HD 82943 system and outline several possible stable configurations for a system harboring two companions in a 2:1 resonance, and our new findings are that this system can be locked in aligned or antialigned orbital motions, which these results are never presented nor analyzed in other studies. Specifically, at first, we introduce the numerical setup (see \\S2) for the dynamical simulations. In \\S3, we make a study of the various planetary configurations for HD 82943 by means of the semi-analytical treatment and direct integrations over long-term orbital evolution. We present the numerical results of HD 82943 by employing two series of the best-fit orbital solutions, and we discover three kinds of stable planetary geometry for this system that all are linked to a 2:1 resonance. Furthermore, we advance a semi-analytical model that can avoid the difficulties in  the perturbation expansions for the larger eccentricities and still help to explain the numerical evolutions. Still, we investigate how the stability for the resonant topology depends on the planetary mass ratio and the orbital parameters on the basis of the new fit (Mayor et al. 2004). In \\S4, we explore whether there exist an Earth-like planet surviving about the Habitable zones for the authentic systems of HD 82943 and GJ 876.  The other extraordinary phenomenon in the exosystems is that many of them can host giant Jupiter-like planets with strikingly larger eccentricities, and in Figure 1 is shown the distribution of the eccentricities of 111 planets\\footnote{The data were taken from http://exoplanets.org, as of Aug. 5, 2003} discovered at present day. The figure exhibits that more than 50\\% of the planets have the eccentricities larger than 0.30, and HD 80606 b (Naef et al. 2001) can occupy the eccentricity amounting up to 0.93, which indicates that such unusual discoveries are quite different from the circumstances in our solar system, where most of the major planets revolve around Sun on the near-circular orbits. In addition, from the view of the observations, the larger eccentricities seem to favor the planet detectability, where the semi-amplitude of wobble velocity $K \\propto {M_{p} \\sin i}/ {\\sqrt{a(1-e^2)}}$ (with $M_{p}\\ll M_{c}$), herein $M_{c}$, $M_{p}$, $a$, $e$ and $i$ are, respectively, the stellar mass, the planetary mass, the orbital semi-major axis, the eccentricity and the inclination of the orbit relative to the sky plane, requires more massive planetary mass ($\\sim$ the order of Jupiter's mass), a lower $a$ (e.g., close-in giant planets) and a higher $e$ (see Fig. 1). Moreover, many previous authors have put forward diverse theories and possible mechanisms to interpret the evolution and origin of the orbital eccentricity variation in the exosystems: the disk-planet or planet-planet interaction (Goldreich \\& Tremaine 1980; Weidenschilling \\& Marzari 1996; Lin \\& Ida 1997; Ford, Rasio, \\& Sills 1999) render the orbital migration of gas giant planets embedded in the protoplanetary systems (Lin, Bodenheimer, \\& Richardson 1996; Ward 1997; Bryden et al. 2000; Nelson et al. 2000) and the eccentricities are usually believed to excite to the observed values through complicated processes (Snellgrove, Papaloizou, \\& Nelson 2001; Kley 2003) or hybrid mechanisms (Lee \\& Peale 2002; Chiang, Fischer, \\& Thommes 2002) related to both the planetary migration and the resonant capture. Therefore, as a second objective, we intend to preliminarily make a discussion of the possible origin of the eccentricities or resonant configurations (see \\S5) for two planets of HD 82943. As a final part, in \\S6, we summarize our principal results and give a brief discussion. ",
        "conclusions": "  Using the earlier best-solutions of Fit 1, we found three types of stable orbits for HD 82943, while we observed the steady configurations of Types I and II by employing the new solution of Fit 2 supplied by Mayor et al. (2004). In the meantime, we also proposed a semi-analytical model to study the $e_{i}-\\Delta\\varpi$ Hamiltonian contour determined by the initial parameters, and  the closest level curves encompass the direct numerical results, which presents a good agreement between them. The theoretical figures can still provide valuable information on the dynamics for two planets.  Given the stable geometry based on Fit 2, we then extensively examined the dependence of the stability of HD 82943 in the orbital parameter space and the planetary mass ratios. In the case of the non-coplanar circumstances, we found that stability requires that the relative inclination be $\\sim 25^{\\circ}$ or less. For a fixed planetary mass ratio $\\sim 1$, the stable orbits for HD 82943 requires $\\sin i \\geq 0.50$. For the non-constant mass ratio case (where $m_{2}$ is kept), the requirement of the stability is $m_{1}$/$m_{2}\\leq 2$, which indicates an upper boundary for the mass of the inner planet. Concerning the eccentricities, the system can be always steady when $0 < e_{2}\\le 0.24$ and $0 < e_{1}< 0.60$. In a word, these outcomes do demonstrate that the stability for HD 82943 is strongly sensitive to the sound planetary masses, rather lower relative inclination between two orbits, the eccentricities and other orbital parameters. In addition, we showed that the assumed terrestrial bodies cannot exist in the habitable zones of HD 82943 thanks to strong perturbations induced by two massive planets, but the Earth-like planets can be dynamically habitable in the GJ 876 system at $\\sim 1$ AU in the numerical surveys.  However, as the additional measurements will improve the present best-fit solution for HD 82943 with a corrected model by taking into account the planet-planet interactions. Thus, we hope to have more precise orbital solution not only to check for our presented predictions in this paper but to better make sense of the complete dynamics and the origin of this system in the near future."
    },
    "0403/astro-ph0403665_arXiv.txt": {
        "abstract": "We present the first detection of orbital motion in the cataclysmic variable \\exhydra\\ based on X-ray data from the {\\it Chandra X-ray Observatory}. The large collecting area of the telescope and the high resolution of the HETG spectrometers allow for an unprecedented velocity accuracy of $\\sim$15 \\kms\\ in the X-ray wavelength regime. We find an emission line velocity amplitude of 58.2$\\pm$3.7~\\kms \\ and infer a white dwarf mass of $0.49\\pm0.13~{\\rm M}_\\odot$, in good agreement with previous studies using optical, ultraviolet, and far ultraviolet data. ",
        "introduction": "Magnetic cataclysmic variables (CVs) are semi-detached, interacting binaries consisting of a late-type star that overflows its Roche-lobe and a white dwarf with a magnetic field sufficiently strong to influence the accretion of matter originating from the late-type companion.  The magnetic field directs the accretion flow away from the orbital plane and onto both of the white dwarf's magnetic poles. Accretion behaves differently in the two classes of magnetic cataclysmic variables. In polars the stars are tidally locked; this ``static'' orientation results in an accreting flow from the Lagrangian point L1 of the binary along the magnetic field lines of the white dwarf onto its magnetic pole(s). In intermediate polars (IPs), where the white dwarf is spinning faster than the orbital period, matter also accretes onto the white dwarf along the magnetic field lines; however, the matter first accumulates in an accretion disk surrounding the white dwarf. In both polars and IPs the accreting matter approaches the poles of the white dwarf supersonically and so passes through a stand-off shock before cooling and settling onto the white dwarf \\citep{aiz1973,kyl1982}.  The cataclysmic variable \\exhydra\\ belongs to the IP class. It consists of a white dwarf with a spin period of 67 minutes and has an orbital period of 98 minutes. Both the spin period and binary period can be identified in its light curves, from optical \\citep*[e.g.,][]{vog1980,ste1983}, UV \\citep*[e.g.,][]{mau1999,bel2003}, and X-ray emission \\citep*[e.g.,][]{wat1978,all1998}. Furthermore, the inclination of the \\exhya\\ system is such that its light curve shows a partial eclipse of the lower accretion pole of the white dwarf by the companion \\citep{all1998}.  Optical and UV spectroscopy of the \\exhya\\ system has detected the orbital motion associated with the white dwarf \\citep*[e.g.,][]{bre1980,hel1987,mau1999,bel2003}. This motion, combined with the orbital velocity of the companion and the presence of the eclipse, allows for a mass determination of the white dwarf independent of models of the accretion column.  In this letter we use {\\it Chandra} High-Energy Transmission Grating (HETG) data for \\exhya\\ and show that its orbital motion can be detected with high resolution X-ray spectra. ",
        "conclusions": "We have confirmed the orbital velocity of the white dwarf in the cataclysmic variable \\exhya\\ using the {\\it Chandra} HETG. This measurement is the first detection of orbital motion in a CV binary system based on X-ray spectroscopy. By adding various spectral lines visible in the \\exhya\\ spectrum enough signal-to-noise ratio can be accumulated to obtain velocities accurate to 10--15 \\kms. No new conclusions are drawn on the white dwarf mass (0.49$\\pm$0.13 ${\\rm M}_\\odot$)."
    },
    "0403/astro-ph0403314_arXiv.txt": {
        "abstract": "s{ We positionally match the 6 cm GB6, 20 cm FIRST and NVSS, and 92 cm WENSS radio catalogs and find 16,500 matches in $\\sim$3,000 deg$^2$ of sky. Using this unified radio database, we construct radio ``color-magnitude-morphology\" diagrams and find that they display a clear structure, rather than a random scatter. We propose a simple, yet powerful, method for morphological classification of radio sources based on FIRST and NVSS measurements. For a subset of matched sources, we find optical identifications using the SDSS Data Release 1 catalogs, and separate them into quasars and galaxies. Compact radio sources with flat radio spectra are dominated by quasars, while compact sources with steep spectra, and resolved radio sources, contain substantial numbers of both quasars and galaxies.} ",
        "introduction": " ",
        "conclusions": "The distribution of sources in radio ``color-magnitude-morphology\" space displays clear structure which encodes detailed information about the astrophysical processes that are responsible for the observed radio emission, and thus provides strong constraints for the models of these processes.   \\vskip -0.1in \\phantom{x}"
    },
    "0403/astro-ph0403122_arXiv.txt": {
        "abstract": "We present a numerical study of an idealized magnetohydrodynamic (MHD) configuration consisting of a planar wake flow embedded into a three-dimensional (3D) sheared magnetic field. Our simulations investigate the possibility for in-situ development of large-scale compressive disturbances at cospatial current sheet -- velocity shear regions in the heliosphere. Using a linear MHD solver, we first systematically chart the destabilized wavenumbers, corresponding growth rates, and physical parameter ranges for dominant 3D sinuous-type instabilities in an equilibrium wake--current sheet system. Wakes bounded by sufficiently supersonic (Mach number $M_s > 2.6$) flow streams are found to support dominant fully 3D sinuous instabilities when the plasma beta is of order unity. Fully nonlinear, compressible 2.5D and 3D MHD simulations show the self-consistent formation of shock fronts of fast magnetosonic type. They carry density perturbations far away from the wake's center. Shock formation conditions are identified in sonic and Alfv\\'enic Mach number parameter space. Depending on the wake velocity contrast and magnetic field magnitude, as well as on the initial perturbation, the emerging shock patterns can be plane-parallel as well as fully three-dimensionally structured. Similar large-scale transients could therefore originate at distances far above coronal helmet streamers or at the location of the ecliptic current sheet. ",
        "introduction": "\\label{SectionIntroduction} The magnetic and plasma flow configuration in a solar coronal streamer belt beyond its helmet cusp is characterized by a region of cospatial magnetic shear and a cross-stream velocity variation reminiscent of a `wake' flow. Suitably simplified, a local study of streamer belt dynamics may then consider a `boxed' model of a fluid wake -- current sheet configuration as sketched in Fig.~\\ref{ModelFig}. A large number of previous studies have identified the various linear MHD instabilities supported by such an equilibrium configuration, with purely incompressible studies~\\cite{EinaudiBoncinelliDahlburgKarpen1999}, as well as extensions to the compressible regime~\\cite{DahlburgKeppensEinaudi2001,WangLeeWeiAkasofu1988,EinaudiChibbaroDahlburgVelli2001,DahlburgEinaudi1999}. Their applications cover a range of physical problems, as they are used to model dynamics at the heliospheric current sheet~\\cite{WangLeeWeiAkasofu1988}, mechanisms for variable slow solar wind formation and acceleration~\\cite{EinaudiBoncinelliDahlburgKarpen1999}, and plasmoid formation and evolution in the solar streamer belt~\\cite{EinaudiChibbaroDahlburgVelli2001,DahlburgKarpen1995,Einaudi1999}.  Similarly idealized configurations equally containing velocity shear layers co-spatial with an electric current sheet or neutral sheet have been used to investigate the linear and nonlinear stability properties of planar and non-planar current -- vortex sheets~\\cite{Dahlburg1998,ShenLiu1999,DahlburgBoncinelliEinaudi1997,AntognettiEinaudiDahlburg2002,DahlburgEinaudi2001,DahlburgEinaudi2000,OfmanChenMorrisonSteinolfson1991} and planar magnetized jets~\\cite{Dahlburg1998,BiskampSchwarzZeiler1998,DahlburgKarpen1994,DahlburgBoncinelliEinaudi1998,OfmanChenMorrisonSteinolfson1991}. These configurations were investigated in the context of solar flares and surges, among other applications. The possibility of having both shear-flow related instabilities, and ideal and resistive magnetic instabilities, allow the plasma to transit to turbulent regimes via various routes in its further nonlinear evolution. Energy can be tapped from the background mean flow, with localized reconnection events releasing magnetic energy.  In this paper, we revisit the stability properties for a compressible wake -- current sheet model of streamer belt dynamics, and extend earlier two-dimensional nonlinear simulations~\\cite{DahlburgKeppensEinaudi2001,EinaudiChibbaroDahlburgVelli2001,AntognettiEinaudiDahlburg2002} to three-dimensional MHD computations. The necessity to model the compressive aspects when applying the `box-model' from Fig.~\\ref{ModelFig} to encompass both the slow and fast solar wind regions is obvious since the wind becomes supersonic sufficiently close to the Sun (at roughly $5\\, R_\\odot$, as follows from Parker's model \\cite{Parker1958}). Compressibility can give rise to new effects as compared to incompressible studies, such as nonlinear steepening of ideal sinuous instabilities under supersonic flow conditions to form fast magnetosonic shock fronts~\\cite{DahlburgKeppensEinaudi2001,AntognettiEinaudiDahlburg2002}, and the domination of fully three-dimensionally structured unstable modes over two-dimensional ones \\cite{DahlburgEinaudi2000,DahlburgKeppensEinaudi2001}. In the incompressible case, the validity of Squires theorem~\\cite{DahlburgKarpen1995} precluded the latter possibility and allowed one to concentrate on 2D dynamics. It is the main result of this paper to show that (1) strongly supersonic magnetized wakes, characterized by a plasma beta of order unity, also have dominant 3D sinuous instabilities which (2) can lead to fully 3D structured fast magnetosonic shock fronts in their nonlinear evolution.  Earlier, formation and amplification of fast magnetosonic shocks was shown for supersonic super-Alfv\\'enic wakes by means of numerical simulations~\\cite{DahlburgKeppensEinaudi2001}. These shocks were also observed in studies of 2D current-vortex sheets in the transonic regime \\cite{AntognettiEinaudiDahlburg2002}. We supplement these earlier findings by a complete parametric study of initial state parameters which lead to shock formation. Our findings that dominant 3D sinuous instabilities are possible correct earlier claims to the contrary~\\cite{DahlburgKeppensEinaudi2001} where wavenumber space was not sufficiently explored. At the same time this calls for a study of the nonlinear development of a sinuous instability in three dimensions, which we perform as well.  It was pointed out by Wang et al.~\\cite{WangLeeWeiAkasofu1988}, on the basis of a linear stability analysis of a similar wake -- current sheet configuration, that instabilities may develop in-situ at the heliospheric current sheet at radial distances of order 0.5 -- 1.5 AU. They emphasized the importance of streaming sausage modes leading to traveling magnetic islands (clouds) and plasmoids. Our results focus on the streaming kink mode, and our nonlinear simulations demonstrate that it can lead to shock-dominated large-scale transients. This re-emphasizes the possibility that not all satellite solar wind observations of such events need to be directly mapped back to a specific solar coronal event (eruptive prominence or CME).  This paper is organized as follows. Section~\\ref{SectionModelEQS} presents the system of MHD equations describing the wake-current sheet evolution and the initial equilibrium configuration studied. Section~\\ref{NumericalTools} describes the numerical tools used for the linear stability analysis and for the nonlinear evolution studies. In Section~\\ref{SectionLinear}, we present the linear stability analysis of the wake-current sheet configuration and identify the previously unexplored domain of the dominating three-dimensional sinuous instability. In Section~\\ref{SectionNonlinear}, the nonlinear MHD simulations with the Versatile Advection Code~\\cite{Toth1996} for different values of wake velocity contrast and magnetic field magnitude are discussed. Formation of fast magnetosonic shocks is demonstrated and their properties and generation conditions are described. Conclusions and a discussion of the relevance to in-situ streamer belt dynamics are placed in Section~\\ref{SectionConclusions}. ",
        "conclusions": "\\label{SectionConclusions} In this paper we performed a numerical study of the linear properties and nonlinear evolution of wake -- current sheet configurations by means of compressible resistive MHD simulations. We focused on three-dimensional effects for the ideal sinuous-type instability. In contrast to previous studies, we have shown that there exist ranges of Mach number $M_s$ and Alfv\\'en Mach number $M_a$, where the most unstable sinuous mode is oblique to the shear layer. Such pure three-dimensional instabilities develop nonlinearly into fully 3D deformations of the wake system containing fast magnetosonic shock fronts. For dominant 3D instabilities, the Mach number should exceed a threshold value $M_s>2.6$. With increase of $M_s$, the range of $M_a$ corresponding to 3D instability decreases, but even for highly supersonic cases, a dominant 3D instability still exists in some range of $M_a$. At the nonlinear stage of this sinuous instability in the supersonic regime, self-consistent formation of fast magnetosonic shocks is observed. They carry density jumps (by factors of 2) in cross-stream fast flow regions far away from the wake center. As a result of mode-mode interaction, the shock strengths and interseparations can vary. Depending on the wake velocity contrast and magnetic field magnitude, as well as on the streamwise and spanwise scale of the initial perturbation, these shocks may have plane-parallel as well as complicated three-dimensional structure. Shock fronts appear only for Mach numbers above some critical value, which depends on magnetic field strength, but is always larger than the hydro limit $M_s=2.1$.  Finally, let us estimate the heliospheric regions, where our box-model predicts the possibility for in-situ compressible transients. Using the Parker model~\\cite{Parker1958} for the interplanetary magnetic field one finds, in a manner originally done by Wang et al.~\\cite{WangLeeWeiAkasofu1988}, Alfv\\'en and sonic Mach number radial dependencies of the form: \\begin{displaymath} M_s(r)=\\frac{\\delta v(r)}{v_a^0} \\sqrt{\\frac{2}{\\gamma \\beta_0}}\\sqrt[4]{\\frac{r^2}{r^2+1}}, \\end{displaymath} \\begin{displaymath} M_a(r)=\\frac{\\delta v(r)}{v_a^0}\\sqrt{\\frac{2r^2}{r^2+1}}, \\end{displaymath} where $\\delta v(r)$ is the velocity contrast between fast and slow wind streams, and $\\beta_0$ and $v_a^0$ indicate the plasma beta and Alfv\\'en speed at 1 AU. For solar wind parameters at the Earth's orbit given by $n=5$ cm$^{-3}$, $B=5$nT, $T_e=2 \\cdot 10^5$K, $T_p=4 \\cdot 10^4$K (taken from~\\cite{Winterhalter1994}), one finds $\\beta_0 \\approx 1.6$, while $v_a^0 \\approx 50$ km/sec. Assuming that the fast-slow solar wind velocity contrast, as determined by Ulysses measurements~\\cite{McComas1998} to be of the order $\\delta v \\approx 350$km/sec, does not change significantly with radius beyond the various magnetosonic transitions, one can compute $M_s(r)$ and $M_a(r)$ dependencies.  Our linear theory predicted that the sinuous mode is destabilized for $M_a \\geq 2.5$ which then corresponds to distances beyond 0.3 AU from the Sun. Relevant heliospheric distances turn out to be 0.3 AU -- 1.5 AU. Since we used the fluid wake scale length and the fast-slow velocity contrast to define dimensionless variables, we can verify whether the instability grows sufficiently fast to influence conditions at 1 AU. The flow wake width is observationally determined to be of order 300.000km at 1 AU, while at a heliospheric distance of order 0.5 AU it is estimated to be around 150.000km~\\cite{Winterhalter1994}. Our time unit will then correspond to approximately 7 minutes. At 0.5 AU, the most unstable sinuous mode has a dimensionless growth rate of 0.0272 and corresponding streamwise wavenumber of $k_x=0.325$. Hence, the sinuous instability could fully develop within about 25 hours, in close correspondence with the transit time to 1 AU. The lengthscale of these compressive perturbations are then of order $2 \\cdot 10^{-2}$AU. This points to the distinct possibility of in-situ shock formation at wake-current sheet configurations within the heliosphere.  \\vspace*{1cm} {\\bf Acknowledgements.} This work was performed as part of the Computational Science programme $\\,$ \"Rapid Changes in Complex Flows\" coordinated by JPG, and funded by the Netherlands Organization for Scientific Research (NWO). Computing facilities were provided by NCF (Stichting Nationale Computerfaciliteiten). RK and JPG performed this work supported by the European Communities under the contract of Association between EURATOM/FOM, carried out within the framework of the European Fusion Programme. Views and opinions expressed herein do not necessarily reflect those of the European Commission."
    },
    "0403/astro-ph0403408_arXiv.txt": {
        "abstract": "{  We have conducted multi-epoch EVN observations of the 22.2~GHz water masers towards the high-mass young stellar object in AFGL 5142. With four observing epochs, spanning a time of $\\sim$1~year, 12 distinct maser features have been detected, 7 out of these detected in more than one epoch. The positions and the velocities of the VLBI features agree well with those of the emission centers previously identified by means of VLA observations. For a few features, persistent over three or four epochs, accurate values of the proper motions are derived. The observed proper motions have an amplitude of 15--20~km~s$^{-1}$, significantly larger than the range of variation of the line-of-sight velocities ($\\pm$4~km~s$^{-1}$ around the systemic velocity). On the basis of their spatial distribution, the observed maser features can be divided into two groups. A model fit to the positions and velocities of the maser features of Group I, detected in the same region (within $\\sim$500~mas) where the massive YSO should be located, demonstrates that these might arise on the surface of a nearly edge-on Keplerian disk, rotating around a massive young stellar object. The maser features of Group II, found at large distances from the YSO ($\\geq$ 1$''$), have positions and line-of-sight velocities in agreement with the blue-shifted lobe of a large scale molecular outflow (traced by the HCO$^{+}$ and SiO emission), and might result from the interaction between the gas flowing away from the young stellar object and the ambient gas of the progenitor molecular core. ",
        "introduction": "The star formation process is  better understood  for low-mass stars ($\\sim$1~M$_\\odot$) than for high-mass stars ($\\geq$ 10~M$_\\odot$). The massive stars are less numerous, on average more distant from the Sun, and  enter the ZAMS phase still enshrouded in their progenitor dust and gas envelope, making optical and near-infrared observations impracticable. This explains why, to date, only a handful of high-mass protostellar candidates have been identified \\citep{Ces97, She98, Hun98, Mol98, Fon04}. To study the formation process of massive stars, high resolution observations at radio, millimetre and far-infrared wavelenghts are needed. The highest resolutions ($\\leq $ 1 mas) are obtainable through the  Very Long Baseline Interferometry (VLBI) technique at radio wavelenghts, which can be used to observe the maser transitions of several molecular species, such as OH, H$_{2}$O, CH$_{3}$OH, observed in the proximity of the high-mass proto-star. Multi-epoch VLBI observations can provide accurate (relative) positions, line-of-sight and transversal velocities of the maser spots (the individual, mas-scale compact, centers of emission), so that the 3-dimensional velocity distribution of the gas traced by the maser transition can be derived. Moreover, the VLBI measurements of the emission properties of the single maser spots can be usefully compared with the maser excitation models to constrain the relevant physical and geometrical parameters.  The first multi-epoch VLBI experiments were performed about 20~years ago towards a few of the strongest 22.2~GHz H$_2$O interstellar masers in the Galaxy, i.e. Orion-KL \\citep{Gen81a}, W51 \\citep{Gen81b}, Sagittarius B2  \\citep{Rei88}, W49 \\citep{Gwi92}, and allowed both the determination of the characteristic pattern of the kinematics of these regions and, by comparing the line-of-sight velocities with the proper motions of the spots, the derivation of accurate source distances. Since then, VLBI observations of 22.2~GHz water masers have been mostly focused on those sources which, selected on the basis of interferometrical observations of (thermal) molecular  tracers, are considered to be among the best high-mass protostellar candidates. The Very Long Baseline Array (VLBA) observations carried out by \\citet{Mos00} towards the water masers in IRAS 20126+4104, one of the best studied examples of a high-mass Young Stellar Object (YSO) associated with a Keplerian disk and a jet/outflow system \\citep{Ces97}, suggest that the H$_2$O masers could arise on a bipolar conical surface, excited by the interaction of an ionized jet with the surrounding molecular gas. VLBA observations of the 22.2~GHz masers performed towards intermediate-mass (NGC 2071, \\citealt{Set02}) and high-mass (IC~1396N, \\citealt{Sly99}; W3 IRS 5, \\citealt{Ima00}) YSOs indicate that this maser emission can originate both in the protostellar disks and at the base of the molecular outflows.  This paper presents multi-epoch VLBI observations of the 22.2~GHz H$_2$O masers towards the high-mass star forming region AFGL~5142 (IRAS05274+3345). At a distance of 1.8 kpc \\citep{Sne88}, its far-infrared luminosity is estimated to be \\ $ 3.8 \\times 10^3$ L$_\\odot$ \\citep{Car90}. Using the Very Large Array (VLA) at 8.4~GHz \\citet{Tof95} have revealed a faint (1 mJy), almost unresolved, continuum source (best interpreted as free-free emission from an ionized wind), which subsequently has been found to be coincident in position with the center of a CO bipolar outflow and with the origin of a jet observed in the H$_2$ near-infrared emission \\citep{Hun95}. Owens Valley Radio Observatory (OVRO) data by \\citet{Hun99} show: \\ 1) a well-collimated SiO-jet and an HCO$^+$-outflow, both aligned with the axis of the CO-outflow, and emanating from the 8.4 GHz continuum source; \\ 2) a compact 88 GHz continuum source coincident in position (within the observational errors) with the 8.4 GHz continuum. The radio flux and the bolometric luminosity (estimated using the IRAS fluxes) can be explained if the exciting source is a massive object, with spectral type B2 or earlier. The VLA NH$_3 $-observations performed by \\citet{Zha02} show a 1$''$ diameter (1800 AU) compact structure, whose morphology and  kinematics are compatible with a rotating disk surrounding a high-mass young star.  The 22.2~GHz water masers in AFGL~5142 have been observed with the VLA at three epochs (1991, \\citealt{Tor92}; 1992, \\citealt{Hun95}; 1998, \\citealt{Hun99}), and found to be distributed within a few arcseconds from the 8.4 GHz continuum source. However, the VLA angular resolution ($\\sim$0\\pas1) is inadequate to determine the detailed spatial distribution of the maser spots, to measure their proper motions, and to investigate the kinematics traced by the H$_{2}$O masers in this source. At the distance of AFGL 5142, the linear resolution attainable using the VLBI technique at 22.2~GHz is $\\sim$2~AU, sufficient in principle to resolve and determine the velocity profile of a protostellar disk (of size of hundreds of AU) and/or to study the mechanism of collimation and acceleration of a jet.  Section~2 of this paper describes our multi-epoch VLBI observations and gives technical details of the data analysis, while Section~3 presents the observational results. Section~4 describes plausible kinematical models for interpreting the measured positions and velocities of the maser features. The conclusions are drawn in Section~5.   \\begin{figure} \\includegraphics[angle= -90, width=\\hsize, trim=0.5cm 0cm 0cm 5cm, clip]{single_vs_vlbi_spec.eps} \\caption{The Effelsberg total-power spectrum observed in Oct 1996 (connected plot) is compared with the integrated flux densities (histogram plot) of the VLBI channel maps. The velocity resolution of both spectra is equal to the channel width of 0.12~km~s$^{-1}$.}   \\label{spec} \\end{figure} ",
        "conclusions": "Using the EVN we have observed the 22.2 GHz H$_2$O masers towards the massive star forming region AFGL 5142 for four epochs (from October 1996 to November 1997). Previous high-angular resolution observations of several thermal tracers in the cm and mm-wavelength band indicate that the exciting source of the water maser emission is a high-mass YSO, of spectral type B2 or earlier.  We identified a final set of 12 distinct water maser features, 7 out of these detected in more than one epoch. For a few features, persistent over three or four epochs, accurate values of the proper motions are derived. The observed proper motions have amplitudes of 15--20~km~s$^{-1}$, significantly larger than the range of variation of the line-of-sight velocities ($\\pm$4~km~s$^{-1}$ around the systemic velocity). The positions and the velocities of the VLBI features agree well with those of the emission centers previously identified by means of VLA observations.  On the basis of their spatial distribution, we have divided the maser features into two groups. Group I, comprising eight VLBI features, is found in the same region (within $\\sim$500~mas) where the massive YSO, believed to be the exciting source of two compact, continuum sources (at 8.4 GHz and 88 GHz), should be located. A model fit to the positions and velocities of these  features indicates that they might arise on the surface of a  Keplerian disk seen nearly edge-on. The fitted value of the YSO mass, $M_{YSO} = 38~M_{\\odot} \\pm 20~M_{\\odot}$, although determined with high uncertainty, is in agreement with the results of  previous observations. Maser features of Group II, found at larger distances ($\\geq$ 1$''$) southward from the YSO, have positions and line-of-sight velocities in agreement with the blue-shifted lobe of a large scale molecular outflow (detected in the HCO$^{+}$ and SiO emission), and might be excited by the interaction of the gas outflowing from the YSO with the ambient gas of the progenitor molecular core.  This work demonstrates that multi-epoch EVN observations are able to measure the proper motions of the strongest and longer-living  22.2~GHz maser features. Our EVN observations towards the source AFGL 5142 suffered two major drawbacks: \\ 1) the time separation between consecutive epochs ($\\geq$ 3~months) is too long compared to the average life time of the maser features; \\ 2) modest sensitivity (with an average detection threshold of $\\sim$0.35~Jy~beam$^{-1}$), owing to the fact that only 5--7 antennae (out of the 11 available to observe at 22.2~GHz) took part in each run. Next we plan to use the VLBA to better constrain the kinematical scenario suggested by the EVN observations, using a shorter time separation ($\\sim$1~month) between two consecutive epochs and higher sensitivity."
    },
    "0403/astro-ph0403252_arXiv.txt": {
        "abstract": "Observations and $N$-body simulations both support a simple relation for the disruption time of a cluster as a function of its mass of the form: $t_{\\rm dis} = t_{4}\\times(M/10^{4}\\,M_{\\odot})^{\\gamma}$.  The scaling factor $t_{4}$ seems to depend strongly on the environment. Predictions and observations show that $\\gamma \\simeq 0.64 \\pm 0.06$. Assuming that $t_{\\rm dis} \\propto M^{0.64}$ is caused by evaporation and shocking implies a relation between the radius and the mass of a cluster of the form: $r_{\\rm h} \\propto M^{0.07}$, which has been observed in a few galaxies. The suggested relation for the disruption time implies that the lower mass end of the cluster initial mass function will be disrupted faster than the higher mass end, which is needed to evolve a young power law shaped mass function into the log-normal mass function of old (globular) clusters. ",
        "introduction": "This study combines the results of theoretical and observational studies to the life time of stellar clusters. It is of vital importance to understand how disruption processes influence a cluster population in order to get better understanding of the observed luminosity and mass distribution. A key question in studying cluster mass functions is whether the power law mass distribution of young clusters, as observed in interacting/merging and starburst galaxies, will evolve into a log-normal mass distribution as observed for old globular cluster populations. If so, a preferential depletion of low mass clusters is necessary. The two most important dynamical processes influencing the cluster mass function are 1) evaporation, which is due to internal relaxation of the cluster and 2) tidal shocking. The first is a result of two body interactions within the cluster which drive the low mass stars through interaction with higher mass stars to the edge of the cluster until they reach a velocity greater than the escape velocity and leave the cluster. Tidal shocks can be caused by passing giant molecular clouds (GMCs) or when the cluster travels through a disk, bulge or spiral arm. The stars in the cluster gain kinetic energy which causes the cluster to expand and become less gravitationally bound. ",
        "conclusions": ""
    },
    "0403/astro-ph0403587_arXiv.txt": {
        "abstract": "Pulsar dynamic spectra sometimes show organised interference patterns; these patterns have been shown to have power spectra which often take the form of parabolic arcs, or sequences of inverted parabolic arclets whose apexes themselves follow a parabolic locus. Here we consider the interpretation of these arc and arclet features. We give a statistical formulation for the appearance of the power spectra, based on the stationary phase approximation to the Fresnel-Kirchoff integral. We present a simple analytic result for the power-spectrum expected in the case of highly elongated images, and a single-integral analytic formulation appropriate to the case of axisymmetric images. Our results are illustrated in both the ensemble-average and snapshot regimes. Highly anisotropic scattering appears to be an important ingredient in the formation of the observed  arclets. ",
        "introduction": "\\label{sec:intro} At low radio frequencies, diffraction from inhomogeneities in the ionised component of the InterStellar Medium (ISM) leads to multi-path propagation from source to observer. Together with the refraction introduced by large-scale inhomogeneities, this leads to a variety of observable phenomena --- image broadening, image wander and flux variations, for example, see Rickett (1990). If the source is very small, as in the case of radio pulsars, then the electric field has a high level of coherence amongst the various propagation paths, and persistent, highly visible interference fringes are seen in the low frequency radio spectrum. These fringes evolve with time, due to the motions of source, observer and ISM. Mostly the observed fringes can be understood in terms of wave propagation through a random medium, with the electron density inhomogeneities being described by a power-law spectrum, close to the Kolmogorov spectrum, and physically identified with a turbulent cascade (Armstrong, Rickett and Spangler 1995; Cordes, Weisberg and Boriakoff 1985; Lee and Jokipii 1976). However, it has been known for many years that some pulsars, at some epochs, also exhibit organised patterns in their dynamic spectra -- drifting bands, criss-cross patterns and periodic fringes, for example (Gupta, Rickett and Lyne 1994; Wolsczan and Cordes 1987; Hewish, Wolszczan and Graham 1985; Roberts and Ables 1982; Ewing et al 1970). The interpretation of these patterns has never been entirely clear, although several authors have suggested an association with enhanced refraction in the ISM (Rickett, Lyne and Gupta 1997; Hewish 1980; Shishov 1974).  Recently a significant breakthrough in data analysis has occurred, with the recognition that these organised patterns are more readily apprehended by working with the so-called ``secondary spectrum'', i.e. the power spectrum of the dynamic spectrum (Stinebring et al 2001): in this domain the very complex patterns seen in the dynamic spectrum typically appear as an excess of power along the locus of either a single parabola, or a collection of inverted parabolic arcs (Hill et al 2003; Stinebring et al 2004). Previous investigations have demonstrated why parabolic structures are generic features of the secondary spectrum (Stinebring et al 2001; Harmon and Coles 1983), but the properties of the observed secondary spectra are not yet understood in detail. In this paper we develop a statistical theory of the power distribution in pulsar secondary spectra, based on the stationary phase approximation, and present some analytic results for simple forms of the scattered image. We compare our results with existing data and argue that highly anisotropic scattering is required to form some of the observed secondary spectra. Our results are consistent with, and complementary to, those of Cordes et al (2003).  This paper is organised as follows. We begin by developing a model secondary spectrum, based on the stationary phase approximation, in \\S2, leading to analytic models for the ensemble-average secondary spectra of linear and axisymmetric images (\\S3), and Monte Carlo simulations of snapshot secondary spectra (\\S4). In \\S5 we consider the appearance of additional, off-centre image components; we compare our models with existing data in \\S6. ",
        "conclusions": "Approximating the Fresnel-Kirchoff Integral by a sum over the stationary phase points of an arbitrary phase screen has allowed us to develop a probabalistic description of the power distribution in pulsar secondary spectra. We have presented a single-integral formulation of the power distribution for general axisymmetric images, and a closed form expression for highly anisotropic (i.e. linear) images. These results are descriptions of the (ensemble-)average properties, but our development makes it obvious how to generate examples of snapshot secondary spectra using the Monte Carlo method, and we have illustrated this technique with two examples. Our results clarify the origins of the parabolic arcs in observed secondary spectra, and provide an explanation for the inverted parabolic arclets which are also sometimes seen.  The observed arclets are due to structured, highly elongated images. Long-lived arclets require correspondingly durable features in the phase-screen, and the data show that compact, isolated clumps of refracting/scattering material must exist in the interstellar medium. Consideration of the lifetime of highly anisotropic image components highlights an observational bias in favour of those which are elongated along the direction of the effective velocity vector."
    },
    "0403/astro-ph0403064_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "Current cosmological data imply the existence of non-baryonic dark matter. We have discussed some of the most popular candidates and shown that none of the candidates known to exist, i.e.\\ the active neutrinos, can be non-baryonic cold dark matter. Hence to explain the nature of cold dark matter we need to invoke hypothetical particles that have not been detected yet. Some of these hypothetical particles have been suggested for reasons different than the dark matter problem (such as sterile neutrinos, neutralinos, and axions), some others have been proposed mainly as a solution to the cold dark matter problem (e.g., self-interacting dark matter, WIMPZILLAs, etc.). Although most studies focus on the first category of candidates, especially neutralinos and axions, we should keep an open mind.  To illustrate how we can find out if dark matter is made of elementary particles, we have used neutralino dark matter as our guinea pig to survey several methods  to search for non-baryonic dark matter. These methods range from a direct detection of dark matter particles in the laboratory to indirect observation of their annihilation products produced in the core of the Sun or of the Earth and in galactic halos, including our own. Direct searches may have found a signal from WIMPs (the annual modulation), but this claim is highly controversial at the moment. Future direct searches have great promise, and might even be able to explore the local velocity distribution of WIMPs. These searches are complemented by indirect searches for high-energy neutrinos from the core of the Sun or of the Earth. Indirect searches using gamma-rays and cosmic rays from annihilations in the galactic halo are subject to uncertainties related to the detailed structure of the dark matter halo. Even more so are predictions for dark matter signals from the Galactic Center. Despite this, some anomalies in cosmic ray fluxes, namely a positron excess, may be explained by neutralino annihilation, and future gamma-ray observations may discover a gamma-ray line from neutralino annihilation in our galactic halo.  All of the examples we have presented are without doubt simple, elegant, and compelling explanations for the nature of non-baryonic dark matter. As we ponder on which one of them is realized in Nature, we must remember the words of astrophysicist Thomas Gold (as quoted by Rocky Kolb): ``For every complex natural phenomenon there is a simple, elegant, compelling, wrong explanation.\""
    },
    "0403/astro-ph0403447.txt": {
        "abstract": "The POINT-AGAPE survey is an optical search for gravitational microlensing events towards the Andromeda Galaxy (M31).  As well as microlensing, the survey is sensitive to many different classes of variable stars and transients. Here we describe the automated detection and selection pipeline used to identify M31 classical novae (CNe) and we present the resulting catalogue of 20 CN candidates observed over three seasons. CNe are observed both in the bulge region as well as over a wide area of the M31 disk. Nine of the CNe are caught during the final rise phase and all are well sampled in at least two colours. The excellent light-curve coverage has allowed us to detect and classify CNe over a wide range of speed class, from very fast to very slow. Among the light-curves is a moderately fast CN exhibiting entry into a deep transition minimum, followed by its final decline. We have also observed in detail a very slow CN which faded by only 0.01~mag~day$^{-1}$ over a 150 day period. We detect other interesting variable objects, including one of the longest period and most luminous Mira variables. The CN catalogue constitutes a uniquely well-sampled and objectively-selected data set with which to study the statistical properties of classical novae in M31, such as the global nova rate, the reliability of novae as standard-candle distance indicators and the dependence of the nova population on stellar environment. The findings of this statistical study will be reported in a follow-up paper. ",
        "introduction": "\\begin{table*} \\begin{minipage}{114mm} \\caption{Principal M31 classical nova surveys} \\label{novasurveys} \\begin{tabular}{lllllll} \\hline Author(s)              & Epoch      & Filter(s)        & Detector & Novae & Annual          & Reference(s) \\\\                       &            & &          &       & rate            &              \\\\ \\hline Hubble                 & 1909--1927 & B                & Plates   &  85   & $\\sim 30$       & 1       \\\\Arp                    & 1953--1954 & B & Plates   &  30   & $24 \\pm 4$      & 2       \\\\Rosino {\\it et al.}    & 1955--1986 & B                & Plates   & 142   & -               & 3, 4, 5 \\\\Ciardullo {\\it et al.} & 1982--1986 & B, H$\\alpha$     & CCD      &  40   & - & 6, 7    \\\\Sharov \\& Alksins      & 1969--1989 & B & Plates   &  21   & -               & 8       \\\\Tomaney \\& Shafter     & 1987--1989 & H$\\alpha$        & CCD      &   9   & -               & 9 \\\\Shafter \\& Irby        & 1990--1997 & H$\\alpha$        & CCD      &  72   & $37_{-8}^{+12}$ & 10      \\\\Rector {\\it et al.}    & 1995--1999 & H$\\alpha$ & CCD      &  44   & -               & 11      \\\\Darnley {\\it et al.}   & 1999--2002 & $r'$, $i'$, $g'$ & CCD      &  20   & $*$             & 12 \\\\ \\hline \\end{tabular} %\\medskip References: (1)~\\citet{1929ApJ....69..103H}; (2)~\\citet{1956AJ.....61...15A}; (3)~\\citet{1964AnAp...27..498R}; (4)~\\citet{1973A&AS....9..347R}; (5)~\\citet{1989AJ.....97...83R}; (6)~\\citet{1987ApJ...318..520C}; (7)~\\citet{1990ApJ...356..472C}; (8)~\\citet{1991Ap&SS.180..273S}; (9)~\\citet{1992ApJS...81..683T}; (10)~\\citet{2001ApJ...563..749S}; (11)~\\citet{1999AAS...195.3608R}; (12)~This work. \\newline $^*$ To be reported in a follow-up paper. \\end{minipage} \\end{table*}   Classical novae (CNe) undergo unpredictable outbursts with a total energy that is surpassed only by gamma-ray bursts, supernovae and some luminous blue variables.  However, CNe are far more commonplace than these other phenomena \\citep{1989clno.book....1W}.  CNe are a sub-class of cataclysmic variables (CVs).  The canonical model for CVs \\citep{1956ApJ...123...44C} is that they are close binary systems, generally of short period, containing a low-mass G or K type near-main-sequence late-type dwarf that fills its Roche lobe (the secondary) and a more massive carbon-oxygen or magnesium-neon white dwarf companion (the primary).  As the secondary fills its lobe, any tendency for it to increase in size through evolutionary processes causes a flow of material through the inner Lagrangian point into the primary's lobe.  The high angular momentum of the transferred material causes it to form a disc around the white dwarf.  Viscous forces within this accretion disc act to transfer material inwards, so that a small amount of the accreted hydrogen-rich material falls on to the primary's surface.  As this layer grows, the temperature of the material increases.  Eventually the temperature may become high enough to initiate hydrogen burning.  Given the correct conditions, this can lead to a thermonuclear runaway in which the accreted material (and possibly some of the ``dredged-up'' white dwarf) is expelled from the system in a nova eruption \\citep{1989clno.book...17K,1989clno.book...39S}.  CNe exhibit outburst amplitudes of $\\sim 10-20$ magnitudes and, at maximum light, display an average absolute blue magnitude of $M_{B}=-8$ with a limit of around $M_{B}=-9.5$ for the fastest \\citep{1981ApJ...243..268S,1989clno.book....1W}.  They are potentially useful as standard candles for extragalactic distance indication, as they exhibit a correlation between their luminosity at maximum light and the rate of their decline \\citep{1929ApJ....69..103H,1945PASP...57...69M, 1956AJ.....61...15A}. Unfortunately, poor light-curve coverage, small sample sizes and a current lack of understanding of how the properties of CNe vary between different stellar populations, have limited their usefulness in this context.  Due to their relatively high frequency, novae may also be used as a tool for mapping the spatial distribution of the population of close binary systems in nearby galaxies. Despite the high luminosities of CNe, our position within the Galaxy prevents us from directly observing more than a small fraction of Galactic CNe that erupt each year \\citep{1997ApJ...487..226S}. Therefore the Galactic nova rate is poorly known, with estimates ranging from 11~yr$^{-1}$ \\citep{1990AJ.....99.1079C} to 260~yr$^{-1}$\\citep{1972SvA....16...41S}. Fortunately, CNe can be readily identified in external galaxies. Surveys of CNe in M31 have been carried out by \\citet{1929ApJ....69..103H}, \\citet{1956AJ.....61...15A}, \\citet{1964AnAp...27..498R,1973A&AS....9..347R}, \\citet{1987ApJ...318..520C}, \\citet{1989AJ.....97.1622C}, \\citet{1991Ap&SS.180..273S}, \\citet{1992ApJS...81..683T}, \\citet{1999AAS...195.3608R} and \\citet{2001ApJ...563..749S} amongst others.  These surveys have resulted in the discovery of around 450 novae and have indicated the global nova rate to be $\\sim 30-40$~yr$^{-1}$ \\citep{2001ApJ...563..749S}. Table~\\ref{novasurveys} summarises the findings of many of these past surveys, with most of the data reproduced from \\citet{2001ApJ...563..749S}.  The relatively high nova rate of M31 and its close proximity to our own Galaxy are major advantages of targeting M31 for nova surveys.  However, since M31 is nearer edge-on than face-on, the task of distinguishing between possible separate disk and bulge nova populations is difficult, and there remains debate surrounding the distribution and rate of novae within M31. POINT-AGAPE (Pixel-lensing Observations with the Isaac Newton Telescope - Andromeda Galaxy Amplified Pixels Experiment) is searching for gravitational microlensing events against the mostly unresolved stars in M31 \\citep{2003A&A...405...15P}.  It uses the wide-field camera (WFC) on the Isaac Newton Telescope (INT) to survey a 0.6~deg$^2$ region of M31 in at least two broadband filters. The survey has very good temporal sampling over the M31 observing season (August -- January) for three seasons and is therefore an excellent database within which to look for novae. Whilst H$\\alpha$ observations have been shown to be a particularly helpful diagnostic for CN identification \\citep{1957QB841.G24......,1983ApJ...272...92C}, the excellent sampling of our survey more than makes up for an absence of H$\\alpha$ data as it allows us to classify the light-curve profiles of different novae.  A commonly used method of describing the overall timescale of an eruption and classifying CNe, is the concept of the nova ``speed class'', developed by \\citet{1939PA.....47..410M} and \\citet{1948AnAp...11....3B}.  Their definition of the various nova classes depends on the time taken for a nova to diminish by three magnitudes below maximum light, $t_{3}$.  Throughout this paper we will use the speed-class definition, modified by \\citet{1957QB841.G24......} for $t_{2}$ times, given by \\citet{1989clno.book....1W}, reproduced in Table~\\ref{speed}. The 2.5-year baseline of our survey, along with the high sampling rate, allows us to classify novae over a wide range of speed classes.  \\begin{table} \\caption{The classification of classical nova light-curves into various speed classes according to the time taken to decrease in brightness by two magnitudes $(t_{2})$ and their $V$-band rate of decline $(dV/dt)$ from maximum light \\protect\\citep{1989clno.book....1W}.} \\label{speed} \\begin{tabular}{lll} \\hline Speed class     & $t_{2}$  & $dV/dt$\\\\ & (days)   & (mag day$^{-1}$)\\\\ \\hline Very fast       & $\\leq 10$    & $\\geq 0.20$ \\\\ Fast            & 11--25   & 0.18--0.08 \\\\ Moderately Fast & 26--80   & 0.07--0.025 \\\\ Slow            & 81--150  & 0.024--0.013\\\\ Very slow       & 151--250 & 0.013--0.008 \\\\ \\hline \\end{tabular} \\end{table}  In order for robust statistical statements to be made about CN populations in external galaxies it is important to take account of the bias induced by the unresolved galactic surface brightness component, which in the inner regions of galaxies may mask much of the CNe light-curve evolution, making identification and classification more difficult. One approach to this is to make the selection procedure for CN candidates completely automated \\citep{2002cne..conf..481D}, so that the detection efficiency for the various CN light-curve morphologies may be assessed as a function of position objectively through Monte Carlo simulation. This is a far from trivial task because the light-curve structure of CNe varies considerably with nova speed class. Automation of the pipeline also requires that, in the absence of H$\\alpha$ observations, the light-curves are well sampled.  The aim of this study is to present the first fully automated search for CN which can be used to assess in an objective manner the statistical properties of the CN population in M31. The study includes some of the finest examples of extragalactic nova light-curves observed to date. In a companion work we will study the spatial distribution and rate of novae in M31.  We will also assess the potential of our nova candidates as distance indicators by calibrating the maximum-magnitude versus rate-of-decline and and other proposed relationships \\citep{1955Obs....75..170B}.  The outline of this paper is as follows. In Section~\\ref{dataset} we describe the POINT-AGAPE survey dataset. In Section~\\ref{reduction} we detail the initial data reduction stages. Section~\\ref{detectpipe} describes the main nova detection pipeline used to define our CN catalogue. The catalogue itself is presented in Section~\\ref{catalogue} and we discuss our main findings in Section~\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  The primary aim of the POINT-AGAPE survey is to search for microlensing events due to compact dark matter in M31. However, the requirements for such a survey also enable us to compile a substantial catalogue of variable stars and transients. In this study we have conducted the first fully-automated search for classical novae (CNe), using objective selection criteria to define our sample. Our aim has been to devise criteria which, as far as possible, do not bias against the detection of CNe of certain speed classes in favour of others, despite the difficulty that the CN light-curve morphology is inextricably linked to its speed class. In the absence of H$\\alpha$ observations, which are an important diagnostic for CN identification, excellent sampling is a crucial pre-requisite for this task.  Our final catalogue of 20 CNe obtained from 3 seasons of data covering 18 months of total observing time spans a wide range of speed class from very fast to very slow. Their light-curve morphologies vary considerably from the smoothly declining very slow CN, PACN-00-02, through to the moderately fast CN, PACN-00-05, which exhibits multiple maxima as well as a deep transition minimum. Among the objects which did not make our catalogue is V9205~D31C, first seen by the DIRECT survey \\citep{1999AJ....117.2810S}. Combining the DIRECT data and POINT-AGAPE data reveals this object to be a Mira with a period of either 700 or 1400 days, making it one of the longest-period and (if as seems most likely it lies in M31) most luminous Miras known. Our PSF-fitting photometry also confirms the anomalous behaviour of the high signal-to-noise ratio microlensing event PA-99-N2 \\citep{2004ApJ...601..845A}, first discovered using our ``superpixel'' photometry pipeline. Overall, despite the absence of H$\\alpha$ data, we are confident that most, if not all, of our sample are indeed CNe. The only possible borderline candidates are PACN-00-07 and PACN-01-01.  In a follow-up study we intend to use our CN catalogue to undertake an objective study of the spatial and speed class distribution of CNe in M31. This may help to answer key questions such as whether there is more than one population of CN and whether CNe of different speed classes tend to be associated with different stellar populations. From Monte-Carlo completeness tests we should be able to assess objectively the underlying global nova rate, the spatial distribution of CNe, and their potential as standard candle distance estimators via the maximum-magnitude versus rate-of-decline and other relationships. The detailed profiles observed for several of the novae,especially prior to maximum, may also prove helpful for constraining theoretical models of the nova outburst."
    },
    "0403/astro-ph0403228_arXiv.txt": {
        "abstract": "Recent measurements of type Ia supernovae as well as other concordant observations suggest that the expansion of our universe is accelerating. A dark energy component has usually been invoked as the most feasible mechanism for the acceleration. However, the effects arising from possible extra dimensions can mimic well the role of a dark energy through a modified Friedmann equation. In this work, we investigate some observational constraints on a scenario in which this modification is given by $H^2 = {8\\pi G \\over 3} (\\rho +C\\rho^n)$. We mainly focus our attention on the constraints from recent measurements of the dimensionless coordinate distances to type Ia supernovae and Fanaroff-Riley type IIb radio galaxies compiled by Daly and Djorgovski (2003) and the X-ray gas mass fractions in clusters of galaxies published by Allen et al. (2002,2003). We obtain the confidence region on the power index $n$ of the modificative term and the density parameter $\\Omega_m$ of the universe from a combined analysis of these databases. It is found that $n=0.06^{+0.22}_{-0.18}$ and $\\Omega_m=0.30^{+0.02}_{-0.02}$, at the 95.4\\% confidence level, which is consistent within the errors with the standard $\\Lambda$CDM model. These parameter ranges give a universe whose expansion swithes from deceleration to acceleration at a redshift between 0.52 to 0.73. ",
        "introduction": "The Hubble expansion, the Cosmic Microwave Background Radiation(CMBR), the primordial Big Bang Nucleosynthesis and the structure formation are the four pillars of the standard Big Bang cosmology. Recent years, it seems that all these cornerstoness combined to point out that the expansion of the universe is speeding up rather than slowing down (for a recent review see Peebles and Ratra 2003). The main evidence comes from the recent well known distance measurements of some distant Type Ia supernovae (Perlmutter et al. 1998, 1999; Riess et al. 1998, 2001). Possible explanations for such an acceleration include: a cosmological constant $\\Lambda$ (Weinberg 1989; Carroll et al. 1992; Krauss and Turner 1995; Ostriker and Steinhardt 1995; Chiba and Yoshii 1999), a decaying vacuum energy density or a time varying $\\Lambda$-term (Ozer and Taha 1987; Vishwakarma 2001; Alcaniz and Maia 2003; Jain, Dev and Alcaniz 2003), an evolving scalar field (referred to by some as quintessence: Ratra and Peebles 1988; Caldwell et al. 1998; Wang and Lovelace 2001; Li, Hao and Liu 2002; Weller and Albrech 2002; Li et al. 2002a,b; Chen and Ratra 2003; Mukherjee et al. 2003), the phantom energy, in which the sum of the pressure and energy density is negative (Caldwell 2002; Hao and Li 2003a,b; Dabrowski et al. 2003), the so-called ``X-matter\" (Turner and White 1997; Zhu 1998,2000; Waga and Miceli 1999; Podariu and Ratra 2001; Zhu, Fujimoto and Tatsumi 2001; Sereno 2002; Alcaniz, Lima and Cunha 2003; Lima, Cunha and Alcaniz  2003), and the Chaplygin gas (Kamenshchik et al. 2001; Bento et al. 2002; Alam et al. 2003; Alcaniz, Jain and Dev 2003; Dev, Alcaniz and Jain 2003; Silva and Bertolami 2003; Makler et al. 2003). All the above mechanisms for accelerating are obtained by introducing a new hypothetical energy component with negative pressure -- the dark energy.  On the hand, many models have appeared that make use of the very ideas of branes and extra dimensions to obtain an accelerating universe (Randall and Sundrum 1999a,b; Deffayet et al. 2002; Avelino and Martins 2002; Alcaniz, Jain and Dev 2002; Jain, Dev and Alcaniz 2002). The basic idea behind these braneworld cosmologies is that our observable universe might be a surface or a brane embedded in a higher dimensional bulk spacetime in which gravity could spread (Randall 2002). The bulk gravity see its own curvature term on the brane which accelerate the universe without dark energy. Here we are concerned with the cosmological model from the modified Friedmann equation as follows \\begin{equation} \\label{eq:ansatz} H^2 = {8\\pi G\\over 3} (\\rho+C\\rho^n). \\end{equation} Freese and Lewis (2002) showed the above term proportional to $\\rho^n$ (dubbed as Cardassion term by the authors) may arise as a consequence of embedding our observable universe as a (3+1)-dimensional brane in extra dimensions, though an elegant and unique 5-dimensional energy-momentum tensor $T_{\\mu\\nu}$ that gives rise to equation (1) has not yet been found. If $n<1$, the new term dominates at late times which implies a modification of gravity at very low energy scales. Particularly, if $n<2/3$, it gives rise to a positive acceleration. Note that, in this scenario, although the universe is flat and accelerating, it contains matter (and radiation) only without any dark energy contribution.  The main goal of this paper is to set observational constraints on this Cardassian expansion model and check whether it is consistent with current cosmological data. We perform a combined analysis of data including the dimensionless coordinate distances to type Ia supernovae (SNeIa) and Fanaroff-Riley type IIb (FRIIb) radio galaxies compiled by Daly and Djorgovski (2003) and the X-ray gas mass fraction in clusters of galaxies published by Allen et al. (2002,2003). These results are timely and complementary to the previous constraints from the angular size of high-$z$ compact radio sources (Zhu and Fujimoto 2002), CMBR (Sen and Sen 2003a,b), the SNeIa database (Zhu and Fujimoto 2003a; Wang et al. 2003; Cao 2003; Szydlowski and Czaja 2003; Godlowski and Szydlowski 2003), large scale structures (Multamaki et al. 2003) and from optical gravitational lensing surveys (Dev, Alcaniz and Jain 2003).  The plan of the paper is as follows. In the next section, we provide a brief summary of the Cardassian expansion scenario relevant to our work. Constraints from dimensionless coordinate distance data of SNeIa and FRIIb radio galaxies are discussed in section~3. In section~4 we discuss the bounds imposed by X-ray gas mass fraction in galaxy clusters. Finally, we present a combined analysis, our concluding remarks and discussion in section~5. ",
        "conclusions": "Now we present our combined analysis of the constraints discussed in the previous sections and summarize our results. In Figure 6, we display the 68.3\\%, 95.4\\% and 99.7\\% confidence level contours in the ($\\Omega_m$,$n$) plane from a combination of databases of dimensionless coordinate distances to SNeIa and FRIIb radio galaxies and the X-ray gas mass fraction in clusters of galaxies. The best fit happans at $\\Omega_m=0.30$ and $n=0.06$. As it shown, although the two Cardassion parameters are not very sensitive to the dimensionless coordinate distance data of SNeIa and FRIIb radio galaxies and the Cardassion power index is not sensitive to the X-ray gas mass fraction data of clusters, a combination of the two data sets gives at the 95.4\\% confidence level that $\\Omega_m=0.30^{+0.02}_{-0.02}$ and $n=0.06^{+0.22}_{-0.18}$, a very stringent constraint on the Cardassion expansion scenario.        Given the two model parameters, $n$, and $\\Omega_m$, Zhu and Fujimoto (2003b) derived the redshift $z_{\\rm car}$ and $z_{q=0}$ satisfying the relation as follows \\begin{equation} \\label{eq:zq=0} (1+z)_{q=0} = (2-3n)^{1\\over {3(1-n)}} (1+z_{\\rm car}) = \\left[ (2-3n)({1\\over \\Omega_m} -1) \\right]^{1\\over {3(1-n)}} \\,\\,. \\end{equation} where $z_{\\rm car}$ is the redshift at which the two terms inside the bracket of Eq.(1) are equal, while $z_{q=0}$ is the redshift at which the universe switches from deceleration to acceleration, or in other words the redshift at which the deceleration parameter vanishes. It was shown (Zhu and Fujimoto 2003b) that, for every value of $\\Omega_m$, there exists a value for the power index of the Cardassian term, $n_{\\rm peak}(\\Omega_m)$, satisfying $(2-3n_{\\rm peak})^{-1} \\exp\\left[3(1-n_{\\rm peak})/(2-3n_{\\rm peak})\\right] = \\Omega_m^{-1} -1$ which makes the turnaround redshift $z_{q=0}$ reach the maximum value, $[ z_{q=0} ]_{\\rm max} = \\exp\\left[1/ (2-3n_{\\rm peak})\\right] -1$. Lower $\\Omega_m$ is, higher $[ z_{q=0} ]_{\\rm max}$ will be. For the lower bound obtained here, $\\Omega_m=0.28$, we have $n_{\\rm peak}=0.0576$ and $[ z_{q=0} ]_{\\rm max}=0.73$. In conclusion, our combined analysis result, $\\Omega_m=0.30^{+0.02}_{-0.02}$ and $n=0.06^{+0.22}_{-0.18}$ at the 95.4\\% confidence level, suggest a Cardassion expansion universe which swithes from deceleration to acceleration around $z_{q=0} \\in (0.52,0.73)$. However, the modified term of the Friedmann equation would dominate at a redshift around $z_{\\rm car} \\in (0.25,0.55)$, a little bit later than the expansion turnaround happans (note that, $z_{\\rm car}$ is generally not equal to $z_{q=0}$), which is simply due to the resulting power index $n$ is well below $1/3$ (Zhu and Fujimoto 2003b).  Alternative cosmologies from a modified Friedmann equation may provide a possible mechanism for the present acceleration of the universe congruously suggested by various cosmological observations. In this paper we have focused our attention on one of these scenarios, the so-called Cardassion expansion in which the universe is flat, matter (and radiation) dominated and accelerating but without any dark energy component. We have shown that stringenet constraints on the parameters $n$ and $\\Omega_m$, that completely characterize the scenario, can be obtained from the combination analysis of the dimensionless coordinate distance data of SNeIa and FRIIb radio galaxies and the X-ray mass fraction data of clusters. It is natually hopeful that, with a more general analysis such as a joint investigation on various cosmological observations, one could show clearly if this scenario constitutes a feasible alternative to other acceleration mechanisms."
    },
    "0403/astro-ph0403702_arXiv.txt": {
        "abstract": "Using the WMAP cosmic microwave background data it is shown that collisions between Alfven wave packets and the cascades generated by these collisions (the Iroshnikov model) can determine the photon temperature fluctuations for arcminute scales on the last scattering surface. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403534_arXiv.txt": {
        "abstract": " ",
        "introduction": "In the paper of Stelmach and Jakacka \\cite{jak01} it was suggested that the effect of acceleration of the universe can be driven by non-homogeneities in the Stephani model \\cite{kra85,kra97}. We perform the statistical verification of this hipothesis using current available astronomical data. We use three different samples of supernovae where Perlmutter sample \\cite{Perlmutter99} is treated as the fiducial data set. For comparison both the Tonry {\\it et al.} \\cite{Tonry03} (improved by Barris {\\it et al.} \\cite{Barris03}) and Knop {\\it et al.} \\cite{Knop03} are used.  To find which model fits to the data in the best way we use the $\\chi^2$ statistics. The parameters of the models are estimated using the maximum likelihood method. One-dimensional probability distribution function (pdf) over model parameters is presented to deepen statistical insight.  Recently found accelerated expansion of our Universe \\cite{sup98} is explained in the literature by the presence of the $\\Lambda$-term or the so-called quintessence matter satisfying the equation of state with negative pressure \\cite{stein98}. Observation of type Ia supernovae (SN Ia) revealed that this form of energy called dark energy could dominate the present evolution of the universe. Therefore it is important to consider different candidates for dark energy. In the paper of Stelmach and Jakacka it was shown that non-homogeneity in spherically symmetric Stephani model can be treated as some kind of dark energy due to which the universe accelerates (we call this model ``S-J model''). It is a sort of fictitious fluid for which density parameter $\\Omega_{\\rm non}$ can be defined. Then dynamics can be formally reduced to the FRW flat model with additional fluid satisfying the  equation of state $p=(-2/3)\\rho$ corresponding to the equation of state for topological defects in a form of domain walls \\cite{IAU87, AJII,stelbd93}.  The best fitting procedure applied to supernovae data as well as confidence level for the redshift-magnitude relation shows that $\\Omega_{\\rm non}$ should dynamically dominate the present evolution of the Universe if we want to explain observational data without cosmological term (see also \\cite{Celerier00}). ",
        "conclusions": "Our investigations show that the Stephani model is an excellent fit to both Perlmutter data points and currently available Knop's  data points. Summarizing, relatively large value of $\\Omega_{{\\rm non}0}=0.61^{+0.08}_{-0.10}$ can explain supernovae data of Perlmutter. In the range of observed redshifts of supernovae ($z<1$) the curve corresponding to $m-z$ relation in our model is almost indistinguishable from the $\\Lambda$CDM (Perlmutter) model with cosmological constant. With future data from SNAP the error bars in the extimation of the model parameters will be reduced significantly. Together with new limits for $\\Omega_{m,0}$ (obtained from extragalactic data) it will be possible discriminate between Stephani model and $\\Lambda$CDM (Perlmutter) model.   Locally our spherically symmetric Stephani universe is indistinguishable from the Friedman one (the same equation of state, the same local geometry), hence local physical cosmology in both models must be the same. In other words standard scenario of the evolution of the universe remains valid. Especially in the early universe effects of non-homogeneity become negligible. This is easily seen if we evaluate the energy density parameter $\\Omega_\\alpha$ for matter as a function of the scale factor $R$ \\begin{equation} \\Omega_\\alpha(R)={1\\over 1-3\\beta C_\\alpha R^{2+3\\alpha}}. \\end{equation} In the early universe relativistic matter dominates ($\\alpha=1/3$) and from the above relation it follows that $\\Omega_\\alpha(R)$ tends very quickly to $1$ for $R \\to 0$. It means that $\\Omega_{\\rm non}(R) \\to 0$. This justifies application of standard physical cosmology (including e.g. primordial nucleosynthesis) for our nonhomogeneous model.  Freese and Lewis in \\cite{Freese02} have recently proposed interesting alternative model explaining the currently accelerating Universe. In this model which is called the Cardassian model the standard Friedmann-Robertson-Walker (FRW) equation is modified by the presence of an additional term $\\rho^n$, namely \\begin{equation} H^2 = \\frac{\\rho}{3} + B\\rho^n, \\end{equation} where $\\rho$ is a total energy density (matter $+$ radiation), and $B$ is a positive constant. This proposal seems to be attractive because the expansion of the universe is accelerated automatically due to the presence of the additional term (if we put $B=0$ then the standard FRW equation is recovered) without postulating existence of unknown form of dark energy.  Note that putting $\\rho\\equiv 3\\rho_{m 0}/R^3$ (and hence $\\alpha\\equiv 0$) and $n=1/3$ and $B\\equiv -3\\beta/\\sqrt{3\\rho_{m 0}}$ the above equation could be rewritten in the form: \\begin{equation} H^2 = \\frac{\\rho}{3} + B\\rho^n={\\rho_{m 0} \\over R^3}+{(-3\\beta)\\over R}. \\end{equation} Comparing to (1) we realize that we obtain correspondence between Cardassian and Stephani models.  When for simplicity we assume that the energy density parameter for radiation matter vanishes ($\\Omega_{r 0}=0$) we recover the model analysed in our previous paper \\cite{Godlowski04}. We note that  considered by us spherically symmetric Stephani model is a special realization or in other words is dynamically equivalent to the Cardassian model for n=1/3. It is interesting that this value of $n$ parameter is reasonable from the observational joint test analysis of both CMB from and  SNIa \\cite{Sen03}.  Let us note that in our model there is the cosmic coincidence problem, i.e. problem why did the nonhomogeneity start to dominate the present evolution of the Universe only fairly recently? There is no satisfactory solution of this problem and some  other face of fine tuning coincidence is required."
    },
    "0403/gr-qc0403090_arXiv.txt": {
        "abstract": "We study the slow-roll inflationary dynamics in a self-gravitating induced gravity braneworld model with bulk cosmological constant. For $E\\gg M_{5}^{3}/M^{2}_{4}$ we find important corrections to the four-dimensional Friedmann equation which bring the standard chaotic inflationary scenario in closer agreement with recent observations. For $M^{3}_{5}/M^{2}_{Pl}\\ll E \\ll M^{3}_{5}/M^{2}_{4}$ we find five-dimensional corrections to the Friedmann equation, which give the known Randall-Sundrum results of the inflationary parameters. ",
        "introduction": "The recent observation data from the Wilkinson Microwave Anisotropy Probe (WMAP) \\cite{wmap}, show strong support for the standard inflationary predictions of a flat Universe with adiabatic density perturbations in agreement with the simplest class of inflationary models \\cite{inflation}. In particular, these observations had significantly narrowed the parameter space of slow-roll inflationary models. We are entering an area where the physics in the early universe can be probed by upcoming high-precision observational data.  In light of these developments, it is important to understand further the inflationary scenario from the theoretical point of view and also from a more phenomenological approach. A new idea that was put forward is, that our universe lies in a three-dimensional brane within a higher-dimensional bulk spacetime and this idea may have important consequences to our early time Universe cosmology. The most common scenario, is the Randall-Sundrum model of a single brane in an AdS bulk \\cite{randall}. In this model the Friedmann equation is modified in early times while standard cosmology is recovered at late times \\cite{binetruy,csaki}. The slow-roll inflation was applied to this model \\cite{maartens}, and it was found that the extra terms in the Friedmann equation act as friction terms which damp the rolling of the scalar field, allowing steeper potentials. However, the recent observational data from CMB anisotropies, seems to exclude steep potentials and strongly constrain monomial potentials in the Randall-Sundrum scenario \\cite{shinji}.  We can also study inflation in other braneworld cosmological models. These models are mainly generalizations of the Randall-Sundrum model. The most interesting models are those which are based on the generalization of the Randall-Sundrum gravitational action. The gravitational action can be generalized in two ways. The first is to add a four-dimensional scalar curvature term in the brane action. This induced gravity correction arises because the localized matter fields on the brane, which couple to bulk gravitons, can generate via quantum loops a localized four-dimensional world-volume kinetic term for gravitons~\\cite{hc,dvali1}. The second is a Gauss-Bonnet correction to the five-dimensional action. This gives the most general action with second-order field equations in five dimensions~\\cite{lovelock}.  The induced gravity model, with no brane tension (i.e., the brane is not self-gravitating) and no bulk cosmological constant, leads to very different behaviour than the Randall-Sundrum model. In the Randall-Sundrum case, gravity becomes five-dimensional at high energies, so that general relativity is modified in the early universe. By contrast, in the induced gravity case, corrections to general relativity become significant at low energies/ late times, and the early universe evolution agrees with standard general relativity~\\cite{dvali1,indgrav,deffayet}. Astrophysical implications of induced gravity have also been considered~\\cite{astro}.   The Gauss-Bonnet model is like the Randall-Sundrum model in the sense that modifications to general relativity arise in the early universe at high energies. The graviton zero mode is also localized at low energies, as in the Randall-Sundrum case~\\cite{gbloc}. The brane cosmology of the Gauss-Bonnet theory has been investigated in Refs.~\\cite{dl,germani,charmousis,gbmore}. In \\cite{lidnun} the slow-roll inflation was applied to the Randall-Sundrum model with a Gauss-Bonnet correction term and it was found that if the inflation is driven by an exponential inflaton field, the Gauss-Bonnet term allows for a spectral index of the scalar perturbation spectrum to take values closer to the recent observational data.  Recently the cosmology of the Randall-Sundrum braneworld was studied included both curvature correction terms: a four-dimensional scalar curvature from induced gravity on the brane, and a five-dimensional Gauss-Bonnet curvature term \\cite{papa}. The combined effect of these curvature corrections to the action removes the infinite-density big bang singularity, although the curvature can still diverge for some parameter values. A radiation brane undergoes accelerated expansion near the minimal scale factor, for a range of parameters. At late times, conventional cosmology is recovered.  In the induced gravity model when the brane tension and bulk cosmological constant are included, the late-time modifications persist, with less fine-tuning needed, but further modifications are introduced ~\\cite{hc,ktt,sahni,burin,mmt}. In this work we show that the inclusion of the brane tension and bulk cosmological constant, have also important consequences in the early time evolution. For $E\\gg M_{5}^{3}/M^{2}_{4}$ we find significant corrections to the standard four-dimensional Friedmann equation. For $M^{3}_{5}/M^{2}_{Pl}\\ll E \\ll M^{3}_{5}/M^{2}_{4}$ we find Randall-Sundrum type five-dimensional corrections to the Friedmann equation for an AdS or Minkowski bulk as expected. Applying the slow-roll inflationary formalism in the case of a quadratic potential, we find that the correction terms to the standard Friedmann equation, bring the standard chaotic inflationary scenario closer to the observations. We also find that because gravity becomes stronger in the high energy regime, chaotic inflation can occur for values of the inflaton field below the Planck scale.  The paper is organized as follows. In Sec. II, we analyse the cosmological evolution of a self-gravitating braneworld model with induced gravity where a cosmological constant is also included, and we find all possible corrections, in the allowed parametric space, to the Friedmann equation. In Sec. III, we apply the slow-roll inflationary formalism with a quadratic potential to the resulted modified Friedmann equations. In Sec. IV, we discuss chaotic inflationary models allowed for various choices of parameters, and Sec. V gives our conclusions. ",
        "conclusions": "We studied the early time cosmological evolution of a braneworld model with induced gravity when the brane tension and the bulk cosmological constant is included. For $E\\gg M_{5}^{3}/M^{2}_{4}$ the evolution is four-dimensional and the Friedmann equation is the standard Friedmann equation of GR supplemented with correction terms due to induced gravity. For $M^{3}_{5}/M^{2}_{Pl}\\ll E \\ll M^{3}_{5}/M^{2}_{4}$ the evolution is five-dimensional and the Friedmann equation is of the Randall-Sundrum type. In late times and at energies $E\\ll M_{5}^{3}/M^{2}_{4}$ the evolution is four-dimensional and the Friedmann equation is the standard Friedmann equation of GR.  We applied the slow-roll inflationary formalism to this model and investigated chaotic inflationary scenarios with AdS or Minkowski bulk. In the induced gravity phase of $E\\gg M_{5}^{3}/M^{2}_{4}$, we found better agreement with the recent observations. The spectral index of the scalar perturbation spectrum is $n_{s}=0.9775$ for the number of e-folds $N=55$, while it becomes $n_{s}=0.9824$, for $N=70$. In the five-dimensional phase we recover the Rundall-Sundrum model with the known results of the spectral index of the scalar perturbation spectrum.   An important issue is whether the spectrum of tensor perturbations is altered in the presence of induced gravity, and also what is the relation between the scalar and tensor perturbations. The calculation of tensor perturbations is more involved than the corresponding calculations in the Rundall-Sundrum model \\cite{mdl}. Up to now it is not known what is the spectrum of tensor perturbations in braneworls with induced gravity.  This work can be extended to a braneword cosmological model where both corrections are included in the five-dimensional action: a four-dimensional scalar curvature from induced gravity on the brane, and a five-dimensional Gauss-Bonnet curvature term. In \\cite{papa} it was shown, that the combined effect of induced gravity and the Gauss-Bonnet term, is to smooth out the initial singularity that is encountered in the individual theories. The resulted model may improve further the value of the spectral index of the scalar perturbation spectrum. It may also help to stabilize the bulk-brane system in the case of a negative tension brane. \\\\ \\\\ {\\bf{Note added:}} The paper \\cite{zhang} recently appeared which studies the slow-roll inflationary dynamics in the five-dimensional RS sector of the induced gravity model."
    },
    "0403/astro-ph0403158_arXiv.txt": {
        "abstract": "Quasar accretion disks are believed to form stars by self-gravity. Low Luminosity Active Galactic Nuclei (LLAGN) are much dimmer galactic centers, and are often believed to be quasars that ran out of gaseous fuel. LLAGN accretion disks should thus co-exist with thousands to millions of stars or proto-stars left from the previous stronger accretion activity. In principle, these stars may produce several important effects: (i) contribute to the optical/UV spectra of some LLAGN; (ii) reprocessing of the stellar radiation in the dusty disks could dominate the LLAGN infra-red spectra; (iii) deplete the (accretion) gas disk much faster than it can accrete onto the supper-massive black hole (SMBH); (iv) stars, individually or in groups, may slow down and modulate the accretion flow significantly due to their inertia. In this way they may produce the LLAGN cut-off disks; (v) alternatively, frequent enough stellar collisions and resulting stellar disruptions could keep the inner disk empty. Here we explore these ideas. We find that, despite ``low'' luminosities of LLAGN, unrealistically high stellar densities are required to make a sizable radiative contribution to the (HST) optical/UV spectra of these galactic nuclei. Stellar contribution to the infrared spectrum is more likely. Further, if LLAGN are in a quasi steady-state for as long as $10^7$ years or more, too high stellar densities would again be required to significantly affect the dynamics of accretion flow. However, if LLAGN are ``short''-lived phenomena, e.g. $t \\simlt 10^5$ years, quiescent states of quiescence-outburst cycles, then embedded stars may be much more important through the mass effects (iii) -- (v). With observations of LLAGN becoming progressively better, it will be more and more difficult to neglect the presence of close stars in and around nuclear accretion disks. ",
        "introduction": "LLAGN are galactic centers intermediate in luminosity and in some key properties between quasars (bolometric luminosity $L_{\\rm b}\\sim 10^{46}$ erg s$^{-1}$) and inactive galactic centers (IGC), such as our own (\\sgra), whose luminosity is about $L_{\\rm b}\\sim 10^{36}$ erg s$^{-1}$.  LLAGN are also known as LINERs (low ionization nuclear emission-line regions; Heckman 1980) or as ``dwarf'' Seyfert nuclei (e.g. Ho, Filippenko \\& Sargent 1997, 2003). Typical nuclear bolometric luminosities of these sources are $L_b \\simlt 10^{42}$ erg s$^{-1}$ (e.g. Ho 2003).  It is well known that most present day galaxies harbor supper massive black holes (SMBHs), implying that LLAGN and IGC too had to be bright accreting sources at some earlier times. LLAGN and IGC are thus quasar relics or quiescent phases of a long-term evolution of galactic nuclei. Some of the processes operating in the quasar phase may have long-lasting effects and continue to influence SMBH activity in the present dim epoch.  One such process is star formation inside quasar accretion disks due to self-gravity at distances greater than $\\sim 10^3$ \\sch radii (e.g. Paczy\\'nski 1978; Kolykhalov \\& Sunyaev 1980, Shlosman \\& Begelman 1989; Goodman 2003), where $R_S = 2 G M_{\\rm BH}/c^2$, the \\sch radius for the black hole with mass $M_{\\rm BH}$. The goal of this paper is to investigate some of the possible effects of the presence of these close and fast moving stars on the LLAGN phenomenon.  Let us summarize the observational and theoretical framework in which bright AGN, LLAGN and inactive galactic centers are currently understood.  Quasar accretion disks are believed to be the ``standard'' accretion disks (Shakura \\& Sunyaev 1973), possibly supplemented by an X-ray emitting corona. The cold accretion disk apparently extends down to the last stable orbit, as evidenced by a very prominent thermal-like UV feature called the Big Blue Bump (BBB) and by the broad \\fe lines observed in many cases (it is our personal opinion that the non-detection of broad \\fe lines in other sources is caused by an over-ionization of the accretion disk surface). The near-Eddington luminosities of quasars, short time scale variability, and the high radiative efficiency with which most of the SMBHs are inferred to have grown (Yu \\& Tremaine 2002) are also best explained in the context of the standard accretion disks since they indeed have very high radiative efficiency.   LLAGN, and our Galactic Center (GC) black hole, \\sgra, are distinctly different in this regard, possibly requiring a radiatively inefficient accretion. Chandra X-ray observations constrain the amount of hot gas at the capture radius in \\sgra\\ and M87 (Baganoff et al. 2003; Di Matteo et al. 2003), which allows one to estimate the Bondi accretion rate onto the SMBHs. For other sources where direct X-ray data are not constraining enough, Ho (2003) notes that the centers of the galaxies hosting LLAGN should be filled with hot winds from young stars (similar to our GC), yielding an accretion rate estimate. In addition, the accretion rate through a cold accretion disk may be estimated by using the optical/UV data in some LLAGN (e.g. Quataert et al. 1999; Ptak et al. 2004). All of these methods lead to the same conclusion -- the expected luminosity from a thin standard accretion disk fed at the deduced accretion rate is some $1-4$ orders of magnitude higher than actually observed.  In addition, in stark contrast to quasars and Seyfert galaxies, LLAGN do not show the BBB in their spectra (Ho 1999), have no short time scale X-ray variability and no broad \\fe lines (e.g. Dewangan et al. 2004; Ptak et al. 2004). This strongly argues for the absence of the cold accretion disks in the inner $\\sim 10^2-10^3 R_S$. Theoretical interpretation of this fact is that the cold disk ``evaporates'' at few hundred to few thousand $R_S$ (e.g. Quataert et al. 1999; Fig. 2 in Ho 2003; Ptak et al. 2004) to form an inner hot Non-Radiative Accretion Flow (e.g. Narayan \\& Yi 1994; Blandford \\& Begelman 1999). The most frequently discussed evaporation mechanism is that of Meyer \\& Meyer-Hofmeister (1994). If the ions of the hot flow are much hotter than the electrons, then the flow is radiatively inefficient and is much less luminous than the standard flow, explaining the discrepancy. In the least luminous (in terms of $L_{\\rm b}/\\ledd$, where $\\ledd = 1.3 \\times 10^{46} M_8$ erg s$^{-1}$, the Eddington limit for the black hole mass $M_8 = \\mbh/10^8 \\msun$) cases, such as \\sgra, it is possible that a cold accretion disk does not even exists (e.g. Falcke \\& Melia 1997; Narayan 2002).   The best studied example of an inactive galactic center is the center of our own galaxy, \\sgra. Genzel et al. (2003; see also Levin \\& Beloborodov 2003) discovered that most of hot young ($t < 10^7$ years) stars in the GC lie in one of the two {\\em stellar disks}, strongly suggesting these stars were formed by the self-gravity in two massive accretion disks existing in earlier times. In the case of the GC, these young stars actually dominate and determine the SMBH activity, since their radiative output is some $\\sim 5$ orders of magnitude greater than that of the accretion flow, and since the hot winds produced by these stars are believed to be the main source of fuel for \\sgra\\ (e.g.  Coker \\& Melia 1997).  It is thus likely that important interactions between close stars and the accretion flow may be taking place in LLAGN as well. Star formation may even continue in the LLAGN phase as well (e.g., see Fig. 1 in Goodman 2003; also Milosavljevic \\& Loeb 2004). Quite generally, close stars should assume more prominent roles in the LLAGN (and IGC) phase than they do in quasars because the mass and the luminosity output of the gaseous accretion disk are modest in these sources. The possible effects of the close stars can be divided into two categories: the luminosity and the mass effects. The former are due to the combined stellar luminosities; whereas the latter are due to stellar mass -- dynamical and accretion effects. We find that the luminosity effects, besides reprocessing of the stellar light into the infrared region, are unlikely to be important due to existing observational constraints. The mass effects are much more likely to be important especially if the disk inner regions are ``light'' as is the case when the disk is unstable to a limit cycle instability (e.g. Siemiginowska, Czerny, \\& Kostyunin 1996; Lin \\& Shields 1986) or is emptied out in the inner part for other reasons. ",
        "conclusions": "Here we suggested that close stars ($R\\ll 1$ pc from the SMBH) created in the accretion disk by self-gravity in the bright quasar phase, may be present in the LLAGN disks. We found (\\S \\ref{sec:rad}) that it is unlikely that these disk stars significantly contribute to the LLAGN optical/UV spectrum directly.  Hot young massive stars are ruled out because most LLAGN are weak in the UV; and low mass stars would have to be too numerous. Stellar contribution in the infra-red range is plausible, however, because SED of standard accretion disks are weak in that region.  The effects of stars' mass, i.e. inertia, are much more realistic because they require a ``small'' number of low-mass stars embedded in the disk (say $10^2-10^4$ stars). These stars can consume the disk faster than it can accrete on the SMBH (if no radial gap forms in the disk). In this case little gas may be flowing into the central SMBH. However, when fitting the optical/UV spectrum of a ``typical'' LLAGN, M81, we found that this model would only work if the disk were quickly and abruptly consumed by the embedded stars in the inner $\\sim$ few hundreds \\sch radii. Further work is needed to determine whether this is natural and may be sustained for long enough times to explain the LLAGN phenomenon.  If a radial gap does form, then stars embedded in the disk do not accrete gas from the disk. However their presence will be still felt by the disk if the disk mass interior to the star is smaller than the mass of the star (or group of stars). This appears to be especially important in the quiescent phase of limit-cycle unstable disks, in which case the stars will clog-up the inner accretion flow for time necessary for accumulation of enough disk matter to push the star in the disk. Since the disk interior to the star's location is quickly emptied out (Syer \\& Clarke 1995), such a star-terminated disk would appear as a cut-off disk.  Finally, we noted that hard stellar collisions in the inner disk region may disrupt the disk -- create temporary inner holes.  We shall however notice that it is unlikely that the star-disk interactions discussed in this paper are the only explanation to the LLAGN ``inner holes'' in the accretion disk. As we have seen in \\S \\ref{sec:load}, the character of star-disk interactions depends strongly on the mass of the SMBH. Observations span a very large range in the SMBH mass, i.e. almost $10^3$ (see e.g. Ho 1999, Ptak et al. 2004). Moreover, the accreting (stellar mass) binary systems may have inner holes in their accretion disks in the low luminosity states as well (e.g. see Esin, McClintock \\& Narayan 1997; McClintock et al. 2003). Therefore the star-disk interactions may be an {\\em additional reason} for a complicated time-variability in a disk that is already unstable (e.g. Lin \\& Shields 1986; Siemiginowska et al. 1996) or time-dependent because of a changing supply of fuel at large distances, but they are unlikely to be the primary reason for such variations in the disk structure.  Nevertheless with observations becoming progressively better, we should expect many surprises, especially exotic variability, from the effects of star-disk interactions."
    },
    "0403/astro-ph0403191_arXiv.txt": {
        "abstract": "We present new VLA C-configuration \\HI\\ imaging of the \\lya-emitting starburst galaxies \\tol\\ and \\iras.  The effective resolution probes neutral gas structures larger than 4.7 kpc in \\tol, and larger than 8.1 kpc in \\iras. Both systems are revealed to be tidally interacting:  \\tol\\ with \\eso\\ (6.6\\arcmin\\ = 72 kpc minimum separation), and \\iras\\ with \\twomass\\ (2.4\\arcmin\\ = 56 kpc minimum separation).  The \\HI\\ emission is extended in these systems, with tidal tails and debris between the target galaxies and their companions.  Since \\lya\\ emission has been detected from both of these primary systems, these observations suggest that the geometry of the ISM is one of the factors affecting the escape fraction of \\lya\\ emission from starburst environments.  Furthermore, these observations argue for the importance of interactions in triggering massive star formation events. ",
        "introduction": "\\label{S1}  Here, we present new VLA \\HI\\ imaging of the starburst galaxies \\tol\\ [$\\alpha$,$\\delta$ (J2000) $=$ 19:27:58.2, $-$41:34:32; D = 37.5 Mpc {({\\\" O}stlin, Bergvall, \\&\\ Roennback 1998}\\nocite{ostlin98})] and \\iras\\ [$\\alpha$,$\\delta$ (J2000) $=$ 08:38:23.2, 65:07:15; D $\\simeq$ 80 Mpc (Assuming H$_0$ = 72 km\\,sec$^{-1}$\\,Mpc$^{-1}$; {Freedman \\etal\\ 2001}\\nocite{freedman01})].  Both galaxies reveal extended tidal structure in neutral hydrogen and are clearly interacting with nearby companions. Each appears to provide evidence for tidally-induced starburst episodes.  These two systems were chosen for imaging because they exhibit prominent \\lya\\ emission \\citep{leitherer95,giavalisco96,kunth98,leitherer02}. Hydrogen \\lya\\ is one of the most important diagnostic emission lines in astrophysics. It is predicted to be luminous in star-forming galaxies \\citep{charlot93}, and could potentially be used as an indicator of star formation activity in high-redshift systems.  However, this application is not straightforward, since the propagation of \\lya\\ photons in starburst galaxies appears to be a complex process that depends on many factors intrinsic to both the starburst itself (e.g., star formation rate, metallicity) and to the surrounding galaxy (e.g., geometry, kinematics, and dust content).  In a large spectroscopic sample of comparatively high-redshift Lyman break galaxies (LBGs; these systems have typical star formation rates $\\gsim$ 10 \\msun\\,yr$^{-1}$ and are thus directly comparable to more local ``starburst'' galaxies), \\citet{shapley03} find various correlations between the emergent \\lya\\ profile (i.e., strengths of absorption or emission) and characteristics of the ISM.  In particular, dustier galaxies have less prominent \\lya\\ emission, or equivalently, redder UV continuum slopes.  \\lya\\ emission strength is shown to anti-correlate with the kinematic offset implied by the redshift difference between \\lya\\ emission and low-ionization interstellar absorption lines; that is, systems with larger outflow speeds usually show less prominent \\lya\\ emission.  Finally, systems with stronger \\lya\\ emission show lower star formation rates, on average, than systems with weaker \\lya\\ emission.  These correlations suggest that many factors can contribute to the escape fraction of \\lya\\ from starburst regions.  One of the most important factors that influences \\lya\\ propagation is resonant scattering.  It is well known that due to high resonant scattering by neutral hydrogen in the ISM, \\lya\\ photons can be attenuated by even small amounts of dust.  It is then expected that only young, relatively dust-free galaxies should be prodigious sites of \\lya\\ production.  Low-metallicity starburst galaxies in the local universe may be considered nearby analogs to such objects which are expected in greater numbers at higher redshifts.  Thus it was surprising that HST observations of the most metal-poor galaxy known, I\\,Zw\\,18, showed only damped \\lya\\ absorption and no emission \\citep{kunth94}.  In stark contrast, the more metal-rich starburst galaxy Haro\\,2 showed prominent \\lya\\ emission \\citep{lequeux95}.  The effects of resonant scattering can be lessened if there is a velocity difference between the site of \\lya\\ emission and the surrounding material from which resonant scattering can occur.  This will Doppler shift the \\lya\\ photon out of resonance and increase its chance of escaping the surrounding ISM.  While this process seems to be complicated (note, for example, that higher outflow speeds suggest less prominent \\lya\\ emission profiles in the sample of {Shapley \\etal\\ 2003}\\nocite{shapley03}), it appears that an important factor is the presence or absence of an outflow (see below). Indeed, \\citet{shapley03} measure an average \\lya\\ offset $>$ 300\\,\\kms\\ for the sample of \\lya-emitting LBGs.  These observations and the studies of {Giavalisco \\etal\\ (1996)}\\nocite{giavalisco96} and \\citet{kunth98} argue for the importance of ISM geometry and kinematics in affecting the \\lya\\ escape fraction.  The studies of \\citet{kunth98}, \\citet{tenoriotagle99}, and \\citet{mashesse03} have examined the characteristics of the ISM kinematics which appear to influence the propagation of \\lya\\ photons in star-forming galaxies.  If static, homogeneous neutral gas with column densities $\\gsim$ 10$^{18}$\\,cm$^{-2}$ shields the ionized gas, no emission will be detected. The resonant scattering of the \\lya\\ photons will lead to increased probability of destruction by any dust which is present.   On the other hand, there may be diffuse \\lya\\ emission which is detectable on sightlines not coincident with the sources of UV photons.  Similarly, if the areal coverage of the neutral gas is not uniform but clumpy, some \\lya\\ emission may be detectable on favorable sightlines.  Finally, if the velocity structure of the neutral gas is not static but rather outflowing from the ionizing regions (outflow velocities \\gsim\\,200\\,\\kms),  \\lya\\ photons to the red of 1216\\,\\AA\\ can escape and \\lya\\ emission may be significant.  This explains the strong \\lya\\ emission detected in some starburst galaxies with complete spatial coverage by neutral gas which is also comparatively rich in both metals and dust.  These observations were undertaken with these arguments in mind.  Our group has obtained new HST/ACS imaging of the \\lya\\ line in a small sample of galaxies (see {Kunth \\etal\\ 2003}\\nocite{kunth03}), two of which are \\tol\\ and \\iras. While these data are only a small segment of the sample that was imaged in \\lya, they will aide in the interpretation of the \\lya\\ data.  We note that our \\HI\\ imaging sample does not include systems showing only \\lya\\ absorption, so we must await a larger sample to draw definitive conclusions on the correlation between \\HI\\ kinematics and the appearance of \\lya\\ emission. However, the results presented here support the scenario that the geometry and kinematics of the ISM are important (although not the only) factors that affect the escape probability of \\lya\\ photons from massive star formation regions. ",
        "conclusions": "\\label{S4}  VLA \\HI\\ imaging of the starburst galaxies \\tol\\ and \\iras\\ has been presented. These two systems are remarkably similar in \\HI\\ content, mass, and current evolutionary state.  In each, we find extended neutral gas between the target and nearby neighbors, suggesting that interactions have played an important role in triggering the massive starbursts in the primary galaxies.  The close proximity of the companions suggests that the interactions were recent, and the similar velocities of both primary and secondary galaxies argues that these systems may end up gravitationally bound.  Since both primary systems are intense \\lya\\ emitters, these data support the interpretation that the ISM kinematics are an important mechanism that affects the escape of \\lya\\ photons from starburst regions.  Combined with the results on LBGs \\citep{shapley03}, these data have immediate implications for the use of the strength of \\lya\\ emission in determining star formation rates, since the results will be dependent on the geometry of the ISM and other factors (e.g., dust, star formation rate) and not on properties inherent to the starburst being considered.   Further \\HI\\ observations of \\lya-emitting galaxies (and, conversely, of starburst systems with no apparent \\lya\\ emission) are certainly warranted to further explore the role of the ISM in regulating the escape of \\lya\\ photons from starburst environments."
    },
    "0403/astro-ph0403472_arXiv.txt": {
        "abstract": "We investigate two important questions about the use of the nonextensive thermostatistics (NETS) formalism in the context of nonlinear galaxy clustering in the Universe. Firstly, we define a quantitative criterion for justifying nonextensivity at different physical scales. Then, we discuss the physics behind the ansatz of the entropic parameter $q(r)$. Our results suggest the approximate range where nonextensivity can be justified and, hence,  give some support to the applicability of NETS to the study of large scale structures. ",
        "introduction": "\\label{1}  The evolution of large scale structures in the Universe is one of the most important questions of modern cosmology. A considerable amount of work has been done on numerical and analytical approaches to characterise the clustering of matter at large scales. One of the difficulties inherent to this subject is that, in order to make progress in the understanding of general models, it is necessary to define methods for structure quantification.  Usually, quantitative studies of large scale structures are based on the two-point correlation function $\\xi(r)$ for the galaxy distribution. Estimates indicate this function is well approximated by $\\xi(r)=(r/r_0)^{-\\gamma}$ (with $\\gamma\\approx 1.8$), where $r_0$ is the correlation length, the scale marking the transition between linear and non-linear regimes. The usual range for $\\gamma$, found in the literature, is $1.5 < \\gamma < 1.97$ (see, e.g., \\cite{hawkins} \\cite{snethlage}), with $r_0 = 5{\\rm h}^{-1}$ Mpc for galaxies and $r_0 = 25{\\rm h}^{-1}$ Mpc for cluster (see, e.g., \\cite{peebles}). At larger scales, the structure of the Universe presents patterns like walls and filaments (with dimensions $\\sim 150 {\\rm h^{-1}}$ Mpc) seeemgly reaching homogeneization at Cosmic Microwave Background (CMB) scales ($\\simgreat 1000{\\rm h^{-1}}$ Mpc). Some authors have had some success describing the clustering properties of visible matter over this wide range of scales in terms of a multifractal phenomenon associated with density thresholds applied to multifractal sets (e.g. \\cite{martinez} \\cite{pan}). However, the relative success of the multifractal approach does not imply a better understanding of the physics behind this framework. Actually, it is not simple to find a dynamical connection between fractal sets and galaxy clustering. Recently, Ramos et al. \\cite{ramos1} (hereafter RWRR) shed some light on this question by putting forward a model based on the generalized thermodynamics (NETS) formalism \\cite{tsallis}. They show that applying the idea of nonextensivity, intrinsic to NETS, it is possible to derive an expression for the correlation function  \\begin{equation} 1+\\xi(r)={1\\over 3} D_2 r^{(D_2-3)}, \\end{equation}  \\noindent using a scale-dependent correlation dimension  \\begin{equation} D_2(r) = 3{{\\rm log}~ [2+a(1-q(r))]\\over {\\rm log}~ 2}, \\end{equation}  \\noindent where the entropic parameter $q(r)$ is given by the following ansatz  \\begin{equation} \\label{qq} r\\sim {1\\over (q-1)^\\beta} \\end{equation}  \\noindent with $a$ and $\\beta$ being free parameters of the model. This approach shows a smooth transition from a clustered Universe to large-scale homogeneity, with $D_2=3$. However, RWRR do not discuss two important questions concerning the conceptual basis of the model: the necessity of defining a criterion which allows us to assume nonextensivity in the context of galaxy clustering and the physics behind the ansatz for $q(r)$. In the next two sections these questions are further developed and some conclusions are drawn at the end. ",
        "conclusions": "\\label{4}  The basis for clustering statistics in cosmology is the study of galaxy distribution in the Universe. The analysis of an increasingly amount of astronomical data can provide the correct framework to explain the formation and evolution of large-scale structure in the Universe. In particular, correlation functions works suggest the possibility of describing the galaxy clustering properties in the context of the multifractal approach. Recently, RWRR presented a better physical interpretation to this approach by deriving multifractality from the NETS formalism. Now, extending their work, we define a quantitative criterium to use the NETS formalism based on deviations from strict extensivity. Assuming it should reach at least 10 times the expected deviation level within a cell containing only two typical galaxies, we obtain the approximate validity domain for the NETS models: $r <10$ Mpc for ${\\rm NETS}_{-}1$ model, $r < 30$ Mpc for ${\\rm NETS}_{-}2$, and at any scale for  ${\\rm NETS}_{-}3$. The choice of $q(r)$ has a physical meaning in the context of hydrodynamics turbulence and can be understood as a ``bottom-up'' fractal cascading,  in conceptual agreement with CDM structure formation models. Although the scenario we propose is not complete, we should keep in mind that there is not a unified framework to explain galaxy clustering in the Universe. Further numerical developments of NETS models will be presented in a forthcoming paper \\cite{dantas}.  This work was partially supported by FAPESP grant 00/06770-2. FMR also thanks the support given by CNPq through the research grant 300171/97-8. CAW thanks the partial support of CNPq through the research grant 300409/97-4-FA."
    },
    "0403/hep-ph0403001_arXiv.txt": {
        "abstract": "s{This rapporteur review covers selected results presented in the Parallel Session HEA2 (High Energy Astrophysics 2) of the {\\it 10th Marcel Grossmann Meeting on General Relativity}, held in Rio de Janeiro, Brazil, July 2003. The subtopics are: ultra high energy cosmic ray anisotropies, the possible connection of these energetic particles with powerful gamma ray bursts, and new exciting scenarios with a strong neutrino-nucleon interaction in the atmosphere.} ",
        "introduction": "Since the early 60's several ground-based experiments have observed extensive air showers, presumably triggered by ultra high energy cosmic rays (UHECRs) interacting in the upper atmosphere.\\cite{Anchordoqui:2002hs} The highest primary energy measured thus far is $E \\sim 10^{20.5}$~eV,\\cite{Bird:1994uy} corresponding to a center-of-mass energy $\\sqrt{s} = \\sqrt{2 m_p E} \\sim 750$~TeV, where $m_p$ is the proton mass. The interest in the origin of these particles is twofold: there is not only the intellectual curiosity about unknown properties of powerful astrophysical scenarios, but also the possibility to probe new physics at energies beyond the reach of any foreseeable man-made experiments.   Theoretically, one expects the CR spectrum to fall off somewhat above $10^{20}$~eV, because the particle's energy gets degraded through interactions with the cosmic microwave background (CMB), a phenomenon known as the Greisen-Zatsepin-Kuzmin (GZK) cutoff.\\cite{Greisen:1966jv} Unfortunately, as one can see in Fig.~\\ref{azukar}, the most recent measurements by the HiRes\\cite{Abu-Zayyad:2002sf} and AGASA\\cite{Takeda:2002at} experiments are apparently in conflict, if only statistical errors are taken into account, and the source of the difference remains unknown. However, if one takes the systematic uncertainties in the energy measurements into account, one finds that both data sets are mutually compatible on the $2\\,\\sigma$ level.\\cite{DeMarco:2003ig} Attempts to explain the AGASA data with a homogeneuos population of astrophysical sources that injects power-law distributions of CRs give unacceptable $\\chi^2$ (see, e.g., Refs.~[\\cite{Fodor:2003bn,Fodor:2003ph,Wick:2003ex}]).  On the other hand, an analysis\\cite{Bahcall:2002wi} of the combined data reported by the HiRes, the Fly's Eye, and the Yakutsk collaborations is supportive of the existence of the GZK cutoff at the $>5\\,\\sigma$ ($>3.7\\,\\sigma$, depending upon the extrapolated energy spectrum) level.\\footnote{This evidence disappears, however, if one assumes that UHECRs are protons and excludes nearby ($\\lwig\\, 50$~Mpc) sources from the otherwise homogeneous distribution.\\cite{Kachelriess:2003yy} In this case even the HiRes-1 data are incompatible with the GZK cutoff on the $3\\,\\sigma$ level.\\cite{Fodor:2003bn,Fodor:2003ph}} The deviation from GZK depends on the set of data used as a basis for power law extrapolation from lower energies. One caveat is a recent claim\\cite{Watson} that there may be technical problems with the Yakutsk data collection. In view of the low statistics at the end of the spectrum and the wide variety of uncertainties in these experiments, perhaps the rational thing to do is to wait for more data, conservatively arguing that the jury is still out.  \\begin{figure} [t] \\postscript{azukar.eps}{0.80} \\caption[...]{Data on the upper end of the cosmic ray energy spectrum with  their statistical error bars. (HiRes,\\cite{Abu-Zayyad:2002sf} AGASA,\\cite{Takeda:2002at} Fly's Eye,\\cite{Bird:wp} Haverah Park,\\cite{Ave:2001hq} and SUGAR.\\cite{Anchordoqui:2003gm}) \\label{azukar}} \\end{figure}   In this Parallel Session we saw many thorough reviews covering all the most interesting and timely topics in CR physics.  In this rapporteur summary we cannot do justice to all the presentations.\\footnote{A scenario in which UHECRs are able to break the GZK barrier was presented by She-Sheng Xue.\\cite{Xue:2002hq}} Priority will be given to two intriguing scenarios, which pose possible explanations of the data. ",
        "conclusions": "We have reviewed some of the interesting ideas currently invoked to explain the origin of UHECRs. We expect future CR experiments, such as the Pierre Auger Observatory,\\cite{Abraham} would clarify the confusing experimental situation and shed light on validity of models discussed in this report.  The Pierre Auger Observatory is designed to measure the energy and arrival direction of UHECRs with unprecedented precision. It will consist of two sites, one in the Northern hemisphere and one in the Southern, each covering an area $S \\approx 3000$~km$^2$.  The Southern site is currently under construction while the Northern site is pending. Once complete, these two sites together will provide the full sky coverage and well matched exposures which are crucial for anisotropy analyses. The base-line design of the detector includes a ground array consisting of 1600 water \\v{C}erenkov detectors overlooked by 4 fluorescence eyes. The angular and energy resolutions of the ground arrays are typically less than $1.5^\\circ$ and less than 20\\%, respectively. The detectors are designed to be fully efficient out to $\\theta_{\\rm max} = 60^\\circ$ beyond $10^{19}$~eV.  In 10~yr of running the two arrays will collect $\\approx 4000$ events above $10^{19.6}$~eV. Such statistics will enable us to solve the GZK puzzle."
    },
    "0403/astro-ph0403466_arXiv.txt": {
        "abstract": "s{ In this proceedings paper, I overview the current status of the X-ray luminosity function of AGNs in the soft (0.5-2 keV) band, extended using {\\it XMM-Newton} and {\\it Chandra} survey data. We found that the number density of low luminosity AGNs peaks later in the history of the universe ($z\\sim 1$) than the that of high luminosity AGNs ($z\\sim 1.7-3$). I also describe the basic results of a spectroscopic followup project of a complete {\\it HEAO-1} hard X-ray limited sample of AGNs using {\\it ASCA} and {\\it XMM-Newton} and present separate intrinsic hard X-ray luminosity functions for unabsorbed and absorbed AGNs. We found that the absorbed AGN XLF drops more rapidly at high luminosities, indicating a deficiency of absorbed luminous AGNs. } ",
        "introduction": "X-ray surveys are practically the most efficient means of finding active galactic nuclei (AGNs) over a wide range of luminosity and redshift. In order to construct an X-ray luminosity function (XLF) of AGNs, enormous efforts have been made to follow up X-ray sources with optical telescopes to establish their nature as AGNs and to measure their redshifts. Now that we have fairly complete samples of X-ray selected AGNs over 6 orders of magnitude in flux, from surveys ranging from all the high galactic latitude sky to the deepest pencil-beam fields. These enable us to construct and probe luminosity functions over cosmological timescales.  In this proceedings article, I overview the current progress of a few projects related to AGN XLF. Firstly, I overview the results of the soft X-ray (0.5-2 keV) luminosity function (SXLF), which is the continuation of our previous work with {\\it ROSAT} samples\\cite{mhs00}, extended with deep {\\it XMM-Newton} and {\\it Chandra} surveys.  While the soft X-ray surveys select against obscured AGNs, in the current situation, the number of available objects and area-flux coverage from extensive surveys make them useful for probing detailed behaviors of XLFs of the unabsorbed portion of the AGN activity. A complementary XLF in the hard band (2-10 keV) (HXLF) enables us to also look into obscured AGNs, and thus it provides most direct measure of the accretion onto supermassive blackholes (SMBHs). Ueda et al. in this volume covers our extensive recent work on HXLF (see also Ueda at al. 2003\\cite{ueda1} for a full description).  In this article, we also present the results of our {\\it XMM-Newton} and {\\it ASCA} spectroscopic followup of a complete hard X-ray flux-limited sample of bright AGNs selected from {\\it HEAO-1} catalogs\\cite{shino}, the basic results of which have been integrated in the Ueda et al's HXLF. ",
        "conclusions": ""
    },
    "0403/astro-ph0403185_arXiv.txt": {
        "abstract": "In this paper we examine the issue of characterising the transport associated with gravitational instabilities in relatively cold discs, discussing in particular the conditions under which it can be described within a local, viscous framework.  We present the results of global, three-dimensional, SPH simulations of self-gravitating accretion discs, in which the disc is cooled using a simple parametrisation for the cooling function. Our simulations show that the disc settles in a ``self-regulated'' state, where the axisymmetric stability parameter $Q\\approx 1$ and where transport and energy dissipation are dominated by self-gravity. We have computed the gravitational stress tensor and compared our results with expectations based on a local theory of transport. We find that, as long as the disc mass is smaller than $0.25M_{\\star}$ and the aspect ratio $H/R\\lesssim 0.1$, transport is determined locally, thus allowing for a viscous treatment of the disc evolution. ",
        "introduction": "One of the basic unknowns of accretion disc theory is the physical mechanism ultimately responsible for angular momentum transport and energy dissipation in the disc.  It is well known that classical hydrodynamical viscosity is not sufficient to provide accretion at the rates inferred from the observations in almost every astrophysical context where accretion discs play a role. The usual way to overcome this difficulty is to assume that transport is dominated by some ``anomalous'' viscous phenomenon, possibly related to collective instabilities in the disc, and to to give some {\\it ad hoc} parameterisations for the viscosity.  The most widely used of such parameterisation is the so-called $\\alpha$-prescription \\citep[see Section 2]{shakura73}.  It has been recently recognised that accretion discs threaded by a weak magnetic field are subject to MHD instabilities (see \\citealt{balbusreview} and references therein), that can induce turbulence in the disc, thereby being able to transport angular momentum and to promote the accretion process.  However, in many astrophysically interesting cases, such as the outer regions of protostellar discs, the ionisation level is expected to be low, reducing significantly the effects of magnetic fields in determining the dynamics of the disc, at least over limited radial ranges \\citep{gammie96}. A possible alternative source of transport in cold discs is provided by gravitational instabilities (\\citealt{lin87,laughlin94,armitage2001}). Moreover, in environments where the disc mass is a significant fraction of the central object mass, such as during the assembly of protostellar discs, it is inevitable that disc self-gravity will play a role.  The axisymmetric stability of a thin disc with respect to its own self-gravity is determined by the parameter $Q$ \\citep{toomre64}, defined as: \\begin{equation} Q=\\frac{c_s\\kappa}{\\pi G \\Sigma}, \\label{Q} \\end{equation} where $c_s$ is the effective thermal speed of the disc, $\\Sigma$ is the surface density, and $\\kappa$ is the epicyclic frequency (which, for a Keplerian disc, is equal to the angular velocity $\\Omega$). The disc is unstable if $Q$ is smaller than a threshold value $\\bar{Q}\\approx 1$. It has been long recognised, especially in the context of galactic dynamics \\citep{hohl71,bertinromeo88}, that the development of gravitational instabilities would lead to a self-regulation process: if the disc is initially cold, in the sense that $Q<\\bar{Q}$, then gravitational instabilities would heat it up on the fast dynamical timescale, bringing it toward stability; on the other hand, if the disc is hot enough to begin with, then radiative cooling is going to bring the value of $Q$ down toward an unstable configuration.  As a result of the presence of these two competing mechanisms, the ``switch'' associated with the onset of gravitational instabilities will act as a thermostat for the disc, which is therefore expected to be always close to marginal stability. A similar approach has been also suggested in the case of accretion discs \\citep{pacinski78,bertin97}.  From the observational point of view, there are already some clues that the disc self-gravity can be important both in the context of protostellar discs and in accretion discs around supermassive black holes in Active Galactic Nuclei (AGN). However, a detailed comparison with observations is limited by the lack of detailed models of self-gravitating discs and by an incomplete understanding of the basic physical processes involved.  The importance of the disc self-gravity in observations of protostellar discs was pointed out by \\citet{adams88}, who realized that a massive disc, with a flat rotation curve, could reproduce the observed flat long-wavelength Spectral Energy Distribution (SED) of many protostellar sources. However, in these early studies no detailed model of self-gravitating accretion discs was available, so that the required disc mass turned out to be unreasonably large. These ideas were recently revisited by \\citet{LB2001} in the light of a self-consistent disc model \\citep{BL99}, and they were able to model the observed SED of FU Orionis objects (a class of young stellar objects undergoing a phase of enhanced accretion), with a disc mass comparable to that of the central object. This model assumes that the outer disc (beyond a few au from the central star) is self-regulated. However, non-self-gravitating $\\alpha$-models of accretion discs would predict a rapidly declining radial profile of $Q$, which would eventually become unphysically small in the outer regions of the disc.  In the case of FU Orionis discs the disc is predicted to be marginally stable already at a distance of $\\approx 1 $ au from the central star \\citep{bellin94}.  The arguments at the basis of the self-regulation mechanisms would suggest that in the outer regions of the disc an additional heating source is required in order to keep the disc marginally stable.  \\citet{LB2001} argued that the difficulty with $\\alpha$-models arises because these viscous models assume that energy dissipation is determined locally, whereas gravitational instabilities would naturally act in a global way, leading to a modification of the standard estimates of the viscous dissipation power.  Analogous considerations also hold in the context of AGN discs. Here, \\citet{LB03a} have shown that the non-Keplerian rotation curve traced by water masers in the Seyfert galaxy NGC 1068 \\citep{greenhill96,greenhill97} can be reproduced by self-regulated disc models. \\citet{sirko03}, following the same arguments of \\citet{LB2001}, have modelled the long-wavelength SED of AGN discs, based on the requirement that $Q\\approx 1$. The latter authors also recognise the need for some additional heating, but, contrary to \\citet{LB2001}, attribute it to some ``external'' source, namely nuclear fusion in stars embedded in the disc (see also \\citealt{collin99}).  Actually, the issue of locality of transport in self-gravitating accretion discs still remains an open question.  \\citet{lin87} have suggested that the transport induced by self-gravity could be described within a viscous framework, and introduced a modified $\\alpha$-prescription, where $\\alpha\\sim Q^{-2}$.  In this way, a self-regulated disc would be characterised by a rather large effective viscosity, with $\\alpha\\approx 1$. On the other hand, \\citet{balbus99} have shown that the energy equation for self-gravitating discs cannot be put in the form of a diffusion equation, as required in a viscous scenario, and that the energy flux contains some extra terms, associated with global wave transport, that are ``anomalous'' from the point of view of viscosity. Similar ideas had also been suggested by \\citet{shu90}.  In this context, numerical simulations play a central role. \\citet{laughlin94} have performed global, smoothed particle hydrodynamics (SPH) simulations of massive isothermal discs, concluding that the density evolution of the disc could be reproduced in terms of simple $\\alpha$-models. \\citet{laughlin96}, using 2D grid based simulations, have shown that in order to reproduce the density evolution induced by gravitational instabilities an $\\alpha$ coefficient dependent on radius was needed.  However, these simulations did not include a detailed treatment of the heating and cooling processes in the disc, which have been shown to play a fundamental role in determining the outcome of the instability \\citep{pickett2000,nelson2000}.  \\citet{gammie01} has performed local, shearing-sheet, zero-thickness simulations of self-gravitating discs, including a simple cooling term, and has concluded that a local description is adequate in such ``razor-thin'' discs.  However, Gammie's simulations are not appropriate to test global effects, since locality is set up from the beginning, and they are only valid for infinitesimally thin discs, while the typical distance over which gravitational instabilities are expected to travel scales with the disc thickness. \\citet{rice03a,rice03b} have already shown using global, three-dimensional simulations, how global effects can be important in the dynamics of self-gravitating discs, in relation to the related issue of the fragmentation of a massive disc.  In this paper we present the results of global, three-dimensional, SPH simulations of massive, cooling, non-magnetised discs.  Our main purpose is to quantitatively determine the dissipation power provided by gravitational instabilities and to compare the results with the expectations based on a viscous theory of discs, in order to assess the extent to which the transport induced by gravitational instabilities can be regarded as a local process.  The paper is organised as follows.  In Section \\ref{sec:general} we summarise the basic transport properties of viscous and of self-gravitating discs, introducing the basic physical quantities involved. In Section \\ref{sec:setup} we describe the numerical setup of our simulations. In Section \\ref{sec:results} we show the results of our computations.  In Section \\ref{sec:discussion} we discuss our results in comparison with previous investigations and in relation to observed systems. In Section \\ref{sec:conclusions} we draw our conclusions. In the Appendix we discuss the more technical issue of transport induced by artificial viscosity in the numerical simulations presented. ",
        "conclusions": "\\label{sec:conclusions}  In this paper we have investigated the transport properties induced by disc self-gravity in relatively massive accretion discs. In particular, we have discussed the extent to which angular momentum transport and energy dissipation can be described within a viscous framework. To this goal, we have performed global, three-dimensional simulations using SPH.  Most of the recent work on the subject has focused on the issue of the conditions for disc fragmentation, especially in relation to the process of planet formation in protostellar discs \\citep{boss02,rice03b,mayer02}, or massive star formation in AGN discs \\citep{goodman03,levin03}. Here, we would like to check the behaviour of self-gravitating discs in the case where fragmentation does not occur and the disc evolves towards a quasi-steady state where cooling is balanced by heating from gravitational instabilities.  The issue of the transport properties induced by gravitational instabilities has been discussed analytically in the past by \\citet{lyndenbell72} and, more recently, by \\citet{balbus99}. There are two major aspects of the problem: {\\it (i)} to what extent is the shear stress $T_{\\mathrm{R\\phi}}$ at a given radius $R$ determined by local conditions?  This point directly calls into question the use of the $\\alpha$ formalism, which explicitly requires that $T_{\\mathrm{R\\phi}}$ is only dependent on the local values of density and thermal velocity (see also the discussion in Section \\ref{sec:self}). The second aspect is {\\it (ii)} whether energy transport is only determined by the gravitationally induced shear stress or whether instead global wave transport occurs, hence influencing energy dissipation in the disc, as argued by \\citet{balbus99}.  We have performed several simulations, with different values of the ratio $M_{\\mathrm{disc}}/M_{\\star}$, to test how the global transport properties of the disc change with increasing total disc mass. Our simulations appear to evolve to a quasi-steady state, characterised by an almost flat profile of the axisymmetric stability parameter $Q\\approx 1$, in which the heating provided by gravitational instabilities balances the imposed cooling rate. We have computed the torque induced by gravitational instabilities and the corresponding viscous dissipation rate, which we then compare to the actual dissipation rate in our simulated discs. We have found values of $\\alpha\\approx 0.05$, in reasonable agreement with the expectation from a viscous theory. Energy dissipation in our simulations is also fairly well described using a viscous approach. These results directly address aspect {\\it (ii)}, described above, and confirm the argument of \\citet{balbus99}, who had indeed argued that for ``self-regulated'' discs, in which $Q\\approx 1$, global wave transport of energy would not play a major role.  On the other hand, we have indeed noticed that more massive discs tend to be dominated by lower $m$ modes, leading to a more global pattern of the gravitational disturbances. We have also been able to determine the size of the region that mostly contributes to the torque at a given point in the disc. More than 80\\% of the torque is produced within a region of size $\\Delta\\approx 10H$, where $H$ is the thickness of the disc. These results directly address issue {\\it (i)}, described above. We can therefore conclude that a viscous description of the transport in self-gravitating discs is only appropriate when $H/R\\lesssim 0.1$. In systems like FU Orionis, where models predict disc thickness $\\approx 0.1$ \\citep{bellin94}, global effects could play a role and modify the dissipation rates in the outer disc.  The results of this work should be extended both toward a more thorough investigation of the parameter space (in particular, it is very important to test the transport properties of discs more massive than $0.25M_{\\star}$, and to explore the results obtained with different initial conditions and cooling timescale, which in the present work has been taken to be $t_{\\mathrm{cool}}=7.5\\Omega^{-1}$), and toward a more realistic description of the cooling function (see, for example, \\citealt{nelson2000} and \\citealt{gammie03})."
    },
    "0403/astro-ph0403699_arXiv.txt": {
        "abstract": "I show that advanced simulations of cosmological reionization are able to fit the observed data on the mean transmitted flux in the hydrogen Lyman-alpha line at $z\\sim6$. At the same time, posteriori models can be constructed that also produce a large value (20\\%) for the Thompson scattering optical depth, consistent with the WMAP measurements. Thus, it appears that a consistent picture emerges in which early reionization (as suggested by WMAP) is complete by $z\\sim6$ in accord with the SLOAN data. ",
        "introduction": "Are we knocking at the door of reionization? At least, it seems so, as the SLOAN quasar survey continues to find ever more distant quasars (Becker et al.\\ 2001; Fan et a.\\ 2002, 2003; White et al.\\ 2003). In fact, highest redshift quasars show essentially no transmitted flux just blueward of the quasar \\HII\\ region, which has been interpreted as the evidence that that reionization took place at a redshift $z\\sim6$ (Becker et al.\\ 2001; White et al.\\ 2003; a similar claim was also made by Djorgovski et al.\\ 2001 based on an extrapolation from a lower redshift observations).  However, one should be cautious before drawing such a conclusion. Indeed, the observed decrease in the mean transmitted flux at $z\\sim 6$ might simply indicate a decrease in the mean ionizing intensity rather then a real reionization of the universe. After all, a neutral fraction of only $2\\times10^{-4}$ is sufficient to absorb effectively all Lyman-alpha radiation at $z=6$ even in the gas that is underdense by a factor of 10. Thus, without an understanding of the evolution of the universe around the reionization epoch, the Ly-alpha absorption data cannot be used to constrain the epoch of reionization (unless damping wings are observed in the absorption profiles, Miralda-Escud\\'{e} 1998).  Fortunately, our theoretical understanding of the process of reionization, in part based on numerical simulations, is solid enough so that the evolution of the ionizing intensity at these redshifts can be predicted with a reasonable confidence level. Combining simulations with the observational data indeed allows one to come up with a consistent picture for the process of cosmological reionization. ",
        "conclusions": "I have shown that advanced self-consistent cosmological simulation with radiative transfer are able to reproduce the observed data on the evolution of the mean transmitted flux from the SLOAN survey when the values of cosmological parameters that best fit the WMAP data are adopted for the underlying cosmology. Thus, a consistent picture is emerging in which the universe reionizes at $6.1\\pm0.3$ after possibly a prolonged period of partial reionization."
    },
    "0403/astro-ph0403650_arXiv.txt": {
        "abstract": "{ We are using VIMOS to conduct a wide-field spectroscopic survey covering fields of $0.5 \\times 0.5\\,\\mathrm{deg^2}$ around the X-ray luminous clusters of galaxies Abell 1689 ($z=0.185$) and Abell 1835 ($z=0.25$). Here we describe the observations and first results on the redshift distribution of subsamples of cluster galaxies to $R\\simeq22$ for which we at present have obtained secure redshifts. These subsamples constitute $\\sim 40\\%$ of the total spectroscopic sample and contain 525 and 630 cluster members in Abell 1689 and Abell 1835, respectively, placing them amongst the largest redshift samples available for any cluster of galaxies.} ",
        "introduction": "\\label{sec:intro}  Imaging surveys of the infall regions of clusters of galaxies which use photometric redshift techniques to isolate cluster galaxies have made strong progress in recent years \\citep[e.~g.\\ ][]{Kodama2001, Gray2003}. While photometric surveys have the advantage of being complete down to well below $M^*$, only spectroscopic surveys can add a third dimension, velocity, to the galaxy distribution in and around clusters. The largest wide--field spectroscopic samples at present exist for near-by clusters, in particular as obtained by the CAIRNS project \\citep{Rines2003}. An example for a previous wide--field redshift survey at higher redshift is described by \\citet{Czoske2001, Czoske2002}, who present a catalogue of redshifts for 300 cluster members with $V \\la 22$ in Cl0024+1654 at $z=0.395$. In that case, it was only the redshift information extending to large cluster-centric distances which revealed the complex structure of what appeared in other observations to be a relaxed rich cluster.  The recent advent of high-multiplex spectrographs on 8--10 meter class telescopes has made it possible to obtain large numbers of high-quality spectra of galaxies and around clusters of galaxies in a short amount of time.  The data described by \\citet{Czoske2001} were obtained over the course of four years. Samples larger by a factor of $2 \\dots 3$ can now be obtained in $\\sim10\\,\\mathrm{hours}$ of observation time. Here I present the first results from a spectroscopic survey of the two X-ray luminous clusters Abell 1689 and Abell 1835 at redshifts $z=0.185$ and $z=0.25$ respectively. We use the VIsible imaging Multi-Object Spectrograph (VIMOS) on VLT UT3/Melipal. ",
        "conclusions": "\\label{sec:conclusions}  The project described in this contribution is still at an early stage. So far, secure redshifts have been obtained for only $\\sim 40\\%$ of the observed samples. Nevertheless, the subsamples presented here include 525 and 630 members galaxies of the clusters Abell 1689 and Abell 1835, respectively. Already, these samples are the largest available for any cluster with the exception of Coma \\citep{Rines2003}, and demonstrate the power of VIMOS to provide unprecedented redshift samples for clusters of galaxies at intermediate redshifts.  We have described the redshift distribution for the two clusters out to $\\ga 3\\,h^{-1}\\,\\mathrm{Mpc}$ ($\\ga 2r_{\\rm vir}$). The infall regions of the two clusters are markedly different, with Abell 1835 showing evidence for strong clustering in its outskirts whereas the outskirts of Abell 1689 seems much more homogeneous. Curiously, this is at contrast to the structure of the central parts of the clusters which is complex for Abell 1689 but seems relaxed for Abell 1835.  The spatially more concentrated galaxy distribution in Abell 1689 as compared to Abell 1835 agrees with results from weak lensing analyses by \\citet{Clowe-Schneider2001, Clowe-Schneider2002}, who find concentration parameters of $c=7.9$ for Abell 1689 and $c=2.9$ for Abell 1835.  The first goal of future work is to complete the redshift determination for the spectral samples and thus to extend the study of the cluster structures to fainter magnitude. After proper assessment of the completeness characteristics of the resulting galaxy catalogues it will be possible to provide mass profiles based on integration of the Jeans equation in the inner parts and analysis of the ``caustic'' structure in the outer parts, following \\citet{Diaferio1999}, which can then be compared to mass profiles derived from weak lensing \\citep{Clowe-Schneider2001, Clowe-Schneider2002, Bardeau2004}.  Further work will relate the spectral types of the galaxies to cluster-centric distance, the local galaxy density and surface mass density as determined from weak gravitational lensing.  \\setlength{\\bibsep}{0pt}"
    },
    "0403/astro-ph0403716_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}[htb]\\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}}  \\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 have analyzed archival {\\it HST/WFPC2} time-series data of the central parts of globular clusters M13, M30 and NGC~6712 in search of variable objects. Among a total of 21 identified variables there are 15 new discoveries. The sample includes nine RR~Lyr stars, two SX~Phe stars  and seven W~UMa-type contact binaries. One object is preliminarily classified as a detached eclipsing binary and another as an ellipsoidal variable. } {globular clusters: individual: M13, M30, NGC~6712 -- binaries: eclipsing -- stars: variables: RR Lyr} ",
        "introduction": "%  The {\\it Hubble Space Telescope} has an excellent resolving power and therefore offers an exeptional opportunity to explore the crowded central regions of Galactic globular clusters. The centers represent rich environments for the study of stellar evolution, dynamical processes in stellar systems as well as evolution of close binary systems.  Albrow \\etal (2001) used HST/WFPC2 time-series data to identify over one hundred variables in the central part of 47 Tuc. Pritzl \\etal (2003) analyzed the archival data on NGC 6441 and located  57 variables,  of which 38 were RR Lyr stars. Similarly, Pietrukowicz \\& Kaluzny (2003) used the archival WFPC2 images to detect 8 new variables  in the central region of globular cluster M22.  In this contribution we report indentification of several new variable stars in fields covering nuclear regions of globular clusters  M13, M30 and NGC~6712. ",
        "conclusions": "%  We have used the archival HST/WFPC2 data to identify 15 new variables located in fields covering central regions of three nearby globular clusters. Among others the sample includes seven W~UMa-type compact binaries, one detached eclipsing binary and two SX~Phe stars. We do not have any definitive information about membership status of these stars but note that investigated regions are strongly dominated by stars belonging to the respective clusters. Also the location of detected variables on the color-magnitude diagrams favoures the hypothesis that most of them are members of investigated  clusters. It is possible to obtain more firm conclusions about membership status of detected variables by obtaining second epoch HST/WFPC2 observations and performing suitable proper motion studies (\\eg Anderson et al. 2003).  Short time coverage and small number of available images limited sample of detectable variables to objects with period of the order of few hours. In fact, all but one of identified objects show continuous variability which makes their detection far easier than detection of variables with a short duty cycle, such as detached eclipsing binaries. Another factor limiting completness of derived sample of variables even for objects with short periods is a presence of many bright stars in the investigated fields. Images of these bright stars are severely saturated on the WFPC2 frames. Large number of dead/bad pixels on WFPC2 CCD's is yet another limiting factor.  \\Acknow{ We would like to thank Grzegorz Pojma\\'nski and Alex Schwarzenberg-Czerny for providing some useful software which was used in this project. We are grateful to Andrew Dolphin for helpful hints on the HSTphot package.  This work is based on observations with the NASA/ESA Hubble Space Telescope, obtained from the Data Archive at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS 5-26555. These observations are associated with programs \\#6121, \\#7379 and \\#8278.  JK was supported by the Polish KBN grant 5P03D00421. }"
    },
    "0403/hep-ph0403129_arXiv.txt": {
        "abstract": "The observational constraints on the primordial power spectrum have tightened considerably with the release of the first year analysis of the WMAP observations, especially when combined with the results from other CMB experiments and galaxy redshift surveys. These observations allow us to constrain the physics of cosmological inflation: \\begin{enumerate} \\item The data show that the Hubble distance is almost constant during inflation. While observable modes cross the Hubble scale, it changes by less than $3\\%$ during one e-folding: $\\dot{d}_{\\rm H} < 0.032$ at $2\\sigma$. The distance scale of inflation itself remains poorly constrained: $1.2 \\times 10^{-28} \\mbox{\\ cm} < d_{\\rm H} < 1 \\mbox{\\ cm}$. \\item We present a new classification of single-field inflationary scenarios (including scenarios beyond slow-roll inflation), based on physical criteria, namely the behaviour of the kinetic and total energy densities of the inflaton field. The current data show no preference for any of the scenarios. \\item For the first time the slow-roll assumption could be dropped from the data analysis and replaced by the more general assumption that the Hubble scale is (almost) constant during the observable part of inflation. We present simple analytic expressions for the scalar and tensor power spectra for this very general class of inflation models and test their accuracy. \\end{enumerate} ",
        "introduction": "With the release of the analysis of the first year WMAP data \\cite{Bennett:2003bz,Spergel:2003cb} and the first 3d SDSS power spectrum \\cite{Tegmark:2003uf}, we got splendid confirmation of the long-standing expectation that the primordial power spectrum of density fluctuations is almost scale-invariant, as anticipated already by Harrison and Zel'dovich in the 1970s. The simplest and most elegant mechanism known to produce primordial spectra with that property is cosmological inflation. Its beauty lies in the fact that we do not need to invent untested physical principles: quantum mechanics and general relativity in four space-time dimensions are good enough. What is needed is a cosmic substratum that gives rise to an epoch of accelerated expansion of the very early Universe. A positive energy density of the vacuum is the simplest example. Trying to figure out the underlying mechanism leads to the pressing question: {\\it What is the scale of cosmological inflation?} It is probably the most poorly constrained energy scale in cosmology; we could argue that cosmological inflation happens somewhere between the electroweak scale (requiring a mechanism to create baryons after the end of inflation) and the highest possible energy scale, which might be the scale of quantum gravity. This implies an uncertainty of some $16$ orders of magnitude.  Although we have no clue on how the world looks like at scales beyond the standard model of particle physics, we can make some generic predictions about the fluctuations in the energy density and space-time structure that are seeded by quantum fluctuations during cosmological inflation \\cite{Mukhanov:xt,Starobinsky:ty}.  We assume that the first prediction of cosmological inflation, spatial flatness of the Universe, is well established and take $\\Omega = 1$ (see \\cite{Benoit:2002mm,Spergel:2003cb} for a detailed discussion). Further, we assume that the observed perturbations are of isentropic nature, which is another prediction of cosmological inflation, unless there is some matter component in the Universe that never ever coupled to the radiation fluid. The most convincing evidence for the dominance of isentropic fluctuations comes from the recent detection of polarization in the cosmic microwave background (CMB) \\cite{Kovac:2002fg,Kogut:2003et}, which requires the presence of a quadrupole anisotropy at the moment of photon decoupling and therefore shows that large-scale fluctuations existed back then. In the following we focus on the power spectra of scalar and tensor fluctuations, which contain all the information if, again as predicted by inflation, the fluctuations are Gaussian (which is consistent with the data \\cite{Komatsu:2003fd}).  Two important inputs are required to predict the inflationary power spectra: the Hubble rate $H$ during cosmological inflation as a function of the scale factor $a$ (or, for later convenience, as a function of the logarithm of the scale factor, denoted by $N$) and the number of dynamical degrees of freedom that drive inflation. In the case of de Sitter space-time this number is zero, but in such a model inflation does not end; the vacuum energy thus must be dynamical. In the simplest scenarios with an exit from inflation a single scalar field $\\varphi$, usually called {\\it inflaton}, is responsible for the acceleration of the Universe. We will restrict the following considerations to that case, with equations of motion, \\begin{eqnarray} \\label{eq:KGphi} \\ddot{\\varphi} + 3H\\dot{\\varphi} + V'(\\varphi) = 0 \\, ,\\\\ \\label{eq:Feq} H^2 = \\frac{1}{3}\\left[\\frac{\\dot{\\varphi}^2}{2}+V(\\varphi)\\right]\\, . \\end{eqnarray} Here, $V(\\varphi)$ is the potential energy density of the inflaton field, the dot stands for a derivative with respect to cosmic time $t$, and the prime denotes a derivative with respect to the inflaton field. We use Planck units, $8\\pi G = c = \\hbar = 1$.  In principle, observations should allow the reconstruction of $H(N)$, but since only a finite interval of wave numbers $k$ is accessible to observations, we can only probe a finite and rather short interval $\\Delta N$. An efficient encoding of $H(N)$ is by the so-called horizon-flow functions, evaluated at a pivot point $N_*$ adapted to the experimental set-up. The horizon-flow functions are a generalization of the slow-roll parameters \\cite{Liddle:1994dx} and are defined recursively as the logarithmic derivatives of the Hubble scale with respect to the number of e-foldings $N$ \\cite{Schwarz:2001vv}: \\begin{equation} \\label{eq:hff} \\epsilon_{m+1} \\equiv {{\\rm d} \\ln |\\epsilon_m|\\over {\\rm d} N} \\quad \\forall m \\geq 0, \\quad \\epsilon_0 = {H(N_{\\rm i})\\over H(N)}\\, , \\end{equation} where $N_{\\rm i}$ denotes an arbitrary `initial' moment. The necessary condition for inflation to take place ($\\ddot{a} > 0$) becomes $\\epsilon_1<1$, and we assume here that the weak energy condition and null energy condition hold true, i.e.~$\\epsilon_1 \\geq 0$. We define the {\\it graceful exit} from inflation as the moment when $\\epsilon_1$ crosses unity. Specifying the set $\\{\\epsilon_m(N_*)\\}$ is equivalent to specifying $H(N)$. A truncation of the set $\\{\\epsilon_m(N_*)\\}$ corresponds to an incomplete knowledge of the evolution of the Hubble rate.  In single field slow-roll models of inflation, all the horizon-flow functions are typically small and, hence, equivalent to the slow-roll parameters. In terms of the inflaton potential $V(\\varphi)$ and its derivatives with respect to the inflaton field, the first two flow functions are given by \\cite{Leach:2002ar}: \\begin{equation} \\epsilon_1 \\approx \\frac{1}{2}\\left(\\frac{V^\\prime}{V}\\right)^2\\, , \\qquad \\epsilon_2 \\approx 2 \\left[\\left(\\frac{V^\\prime}{V}\\right)^2-\\frac{V^{\\prime\\prime}}{V}\\right]\\, . \\end{equation}  CMB experiments (the most recent results come from CBI \\cite{Pearson:2002tr,Readhead:2004gy}, Archeops \\cite{Benoit:2002mk}, ACBAR \\cite{Kuo:2002ua}, VSA \\cite{Grainge:2002da,Dickinson:2004yr} and WMAP \\cite{Bennett:2003bz}) and galaxy redshift surveys (especially 2dF \\cite{Percival:2001hw,Tegmark:2001jh} and SDSS \\cite{Tegmark:2003uf}) have measured the amplitude $A$ and the spectral index $n$ of density fluctuations and constrained the running\\footnote{Two recent works on new data from CBI \\cite{Readhead:2004gy} and VSA \\cite{Rebolo:2004vp} claim to see evidence for a large (compared to expectations from cosmological inflation) running of the spectral index with negative sign, but point out a calibration issue that weakens the evidence.} of the spectral index ${\\rm d}n/{\\rm d}\\ln k$, as well as the amount of gravitational waves, typically expressed as the tensor-to-scalar ratio $r$, that can contribute to the CMB signal. This set of observables $\\{A,n,{\\rm d}n/{\\rm d}\\ln k,r\\}$ could, in principle, provide a measurement of the inflationary parameters $\\{H,\\epsilon_1, \\epsilon_2,\\epsilon_3\\}$.  In a number of works, various combinations of experiments and methods have been used to determine a subset of these parameters, see especially Refs.~\\cite{Peiris:2003ff,Barger:2003ym,Kinney:2003uw,Leach:2003us, Tegmark:2003ud,Rebolo:2004vp}. Most interesting are the constraints in the $\\epsilon_1$--$\\epsilon_2$ plane; see Fig.~\\ref{fig:fig1}. The common finding is that models with almost scale-invariant spectra provide acceptable fits and $\\epsilon_1 < 0.032$ (at $2 \\sigma$) \\cite{Leach:2003us}. The corresponding likelihood contours from that work are shown in Fig.~\\ref{fig:fig1} where the dashed line denotes scale-invariance, up to third order corrections in the horizon-flow functions. Thus, observations show for the first time that $\\epsilon_1 \\ll 1$. The limits on $\\epsilon_2$ are much less restrictive; the situation $|\\epsilon_2| > \\epsilon_1$ covers a large part of the allowed parameter space. This is due to the so-called tensor degeneracy \\cite{Efstathiou:2001cv}, i.e. one can compensate an increase of the tensor contribution by making the scalar spectral index bluer. Since there exist only upper limits on the contribution from gravitational waves, the scale of inflation cannot be fixed by present observations. However, the combination of the upper limit on $\\epsilon_1$ and the measurement of the amplitude of scalar fluctuations also allows an upper limit on the energy scale of inflation $V^{1/4} < 0.01 (\\equiv 2.7 \\times 10^{16} \\mbox{\\ GeV})$ \\cite{Leach:2003us} or a lower limit on the distance scale during inflation $d_{\\rm H} > 1.2 \\times 10^{-28} \\mbox{\\ cm}$. A lower (upper) limit on the energy (distance) scale can be obtained from the facts that the Universe contains baryons and that there is no known mechanism to produce baryons below the electroweak scale ($100$ GeV, respectively $1$ cm) \\cite{Schwarz:2003du}. Thus observations put the scale of inflation at least two orders of magnitude below the Planck scale [this limit will improve proportionally to the limits on $\\sqrt{\\epsilon_1}$ and will be tightened as the value of the optical depth $\\tau$ is pinned down more precisely, since $A \\exp(-2\\tau)$ is the observed quantity].  \\begin{figure} \\label{fig:fig1} \\centerline{\\includegraphics[angle=270,width=0.95\\linewidth]{fig1.eps}} \\caption{Observational constraints, interpretation, and accuracy of various approximations in the $\\epsilon_1$--$\\epsilon_2$ plane: the likelihood curves have been provided by S.~M.~Leach \\cite{Leach:2003us} and enclose the $1\\sigma$-, $2\\sigma$- and $3\\sigma$-allowed regions. Models with a scale-invariant spectrum fall on the dashed line (HZ). An observation of $\\epsilon_2 > 0$ would favour simple models of inflation, in which the evolution toward a graceful exit is seen (toward-exit inflation). If $\\epsilon_2 \\leq 0$, a more contrived scenario must be realized, since in that case the mechanism of exit is hidden from observation (hidden-exit inflation). The region with $\\epsilon_2 > 0$ is further split up by the (full straight) line $\\epsilon_1 = \\epsilon_2/2$, which distinguishes models in which the kinetic energy density decreases (above the line) or increases with time. The latter requires a mechanism to ensure that the kinetic energy density is very small initially. The solid arc centered at the origin encloses the region in which the slow-roll approximation at second order (sra2) provides a prediction of the amplitude at the pivot scale better than $0.3\\%$. Below the hyperbolic lines the same accuracy is achieved for the constant-horizon approximation (cha2 at second order, cha3 at third order) and the growing horizon approximation (gha). The dotted lines represent the boundaries between cha3 and cha2, as well as between cha2 and gha.} \\end{figure}  Based on the observational constraints on the first two horizon-flow functions, we do not restrict our considerations to slow-roll models in this paper. We relax the slow-roll conditions ($\\epsilon_m \\ll 1$ for all $m$) to $\\epsilon_m \\leq 1$ for some $m$. In Section~\\ref{sec:hubble} we demonstrate that the condition $\\epsilon_1 \\ll 1$ is sufficient to guarantee that the Hubble scale is almost constant during the inflationary epoch of interest.  The physical meaning of the first and second horizon-flow functions is discussed in Section~\\ref{sec:classification} in the framework of single-field models, and a new physical classification of inflationary scenarios is introduced, based on the behaviour of the kinetic and total energy densities of the inflaton. To our surprise this approach closely resembles the small field/large field/hybrid classification of Dodelson et al.~\\cite{Dodelson:1997hr}, but is not restricted to slow-roll models.  In Section~4 we derive a procedure to self-consistently calculate higher-order corrections at any scale $k$, given an analytic expression of the power at the pivot scale $k_*$ as a function of the horizon-flow functions. A similar relation has been used in the literature to calculate the spectral index and its running so far, but, to our knowledge, it has never been justified rigorously.  Starting from the observation that $\\epsilon_1 \\ll 1$, we show that the slow-roll approximation is no longer needed in the analysis of the data and that it could be replaced by a more efficient approximation (constant-horizon approximation \\cite{Schwarz:2001vv}). The latter is more efficient in the sense that we need to know less horizon-flow functions in order to predict the power spectra to a given accuracy, as compared with the slow-roll approximation.  We map the classification of models from the $\\epsilon_1$--$\\epsilon_2$ plane to the space spanned by the tilt of the spectrum, $n-1$, and the ratio of tensor and scalar amplitudes $r$ in Section~5. At leading order in the horizon-flow functions this is a one-to-one mapping; including higher orders introduces an ambiguity, due to the running of the spectral index. We thus argue that the $\\epsilon_1$--$\\epsilon_2$ plane is more fundamental. For practical purposes, the $r$--tilt plane is nevertheless a useful tool.  In Section~6 we give examples to demonstrate the accuracy of the new method and discuss the applicability of the new method in the light of the recent data.  Our main results are summarized in Fig.~\\ref{fig:fig1}. We overlay our new classification of models on the likelihood contours as obtained by Leach and Liddle \\cite{Leach:2003us}. We also indicate which approximation to the power spectrum at the pivot scale is best suited for which region of parameter space. We assume an accuracy goal of $0.3\\%$ at the pivot, which puts the theoretical errors safely away from the systematic and statistical errors of WMAP and Planck. ",
        "conclusions": ""
    },
    "0403/astro-ph0403070_arXiv.txt": {
        "abstract": "We present a study of the long-term variability of Cyg X-1 using data from the RXTE/ASM and the RXTE/PCA during the time between the two soft states of 1996 and 2001/2002. This period has been characterized by many short ASM flaring episodes which we have identified as \"failed state transitions\". The 150\\,d period which has been seen before and shortly after the 1996 soft state is not obviously present in the ASM rate during most of this time. Applying selection criteria from our pointed RXTE/PCA observations to exclude the flaring episodes we show that the 150\\,d period can indeed still be significantly detected in the hard state. Furthermore, while the $\\sim$420\\,d timescale associated with the flaring is reduced in the selected hard state count rate, it is still pronounced in the temporal evolution of the corresponding hardness ratios. The Ryle radio flux is also consistent with the 150\\,d period being present but distorted during this time. ",
        "introduction": "Historically several different scales of long-term variability in Cyg~X-1 have been reported. Multiwavelength observations (optical, radio, hard and soft X-ray) between 1996 and 1998 established that a 150\\,d signature is the dominant variability timescale which is suggested to be caused by precession and/or radiative warping of the accretion disk \\citep{pooley:98a,brocksopp:99a}. Recently this period has also been confirmed with earlier Ginga data \\citep{kitamoto:00a}. From 1998 onwards, however, this signature is not obvious in the soft X-ray flux anymore (Fig.~\\ref{fig:asm}, upper panel): Although mainly in the hard state, Cyg~X-1 has been more active since before, i.e., more flaring episodes during which the source can be found in the intermediate state (=failed state transitions \\citep{pottschmidt:02a}) -- sometimes even reaching the soft state -- have been observed. This goes together with a sharp and significant change in the properties of the short-term variability in June 1998 \\citep{pottschmidt:02a}. \\citet{oezdemir:01a} suggested to avoid the distortion of the possible periodicities in the original data introduced by the high/soft state and the failed state transitions by truncating the RXTE/ASM lightcurve at 30\\,cps (threshold derived from visual inspection). They report a clearer detection of the 150\\,d period in the 1996 to 1999 ASM data. Based on the principal idea of selecting only hard state data, we:  \\begin{itemize} \\item perform a period search for RXTE/ASM data dominated by the flare-distorted phase from 1998-2001, \\item use physical selection criteria derived from monitoring \\mbox{Cyg~X-1} with RXTE/PCA since 1998, \\item use two statistics for unevenly sampled data (Lomb-Scargle PSD and epoch folding), \\item also study the variability of the RXTE/ASM colours. \\end{itemize}  Our selection criteria (lower panels of Fig.~\\ref{fig:asm}) are based on the photon index $\\Gamma$ of the PCA spectra and the time lags (between the 2--4 and 8--13\\,keV flux), since they have been shown to be good diagnostics for failed state transitions \\citep{pottschmidt:00a,zdz:02a}. $\\Gamma$ and lag are interpolated from the PCA sampling to the ASM sampling. In detail the states are defined as follows: (i) the canonical hard state is characterized by hard PCA spectra with $\\Gamma < 2.1$, (ii) spectra with $2.1 \\leq \\Gamma \\leq 2.4$ are typical for failed state transitions, and (iii) for the softest spectra with $\\Gamma > 2.4$, enhanced time lags $\\geq 5.3$\\,ms also point towards a failed state transition, while ``normal'' time lags $< 5.3$\\,ms characterize the soft state.  \\begin{figure*} \\centering \\includegraphics[width=1.45\\columnwidth]{benllochs_f1.ps} \\caption{Upper panel: RXTE/ASM 1.3--12.2\\,keV lightcurve of Cyg~X-1. A categorization of the ASM data points into different states has been performed using the parameters and selection criteria displayed in the two lower panels: the parameters have been derived from regular monitoring observations of Cyg~X-1 with the RXTE/PCA and consist of the Fourier time lag between the 2--4\\,keV and 8--13\\,keV high resolution lightcurves (middle panel) as well as of the spectral photon index $\\Gamma$ (lower panel). See text for an explanation of which selection criteria (horizontal lines) apply for which state.}\\label{fig:asm} \\end{figure*} ",
        "conclusions": ""
    },
    "0403/astro-ph0403246_arXiv.txt": {
        "abstract": "The wavelength dependence of interstellar polarization due to dust in M31 has been observed along four sightlines. Only one sightline had been measured previously. The globular clusters, S78, S150, S233 and Baade 327 were used as point sources to probe the interstellar dust in M31. The Serkowski law produces good fits for all the sightlines although the relationship between K and $\\lambda_{max}$ may be different from that found in the Galaxy. The results of this study imply that the slope K/$\\lambda_{max}$ may be significantly larger in M31. The Serkowski curves are significantly narrower than those of the same $\\lambda_{max}$\\ in the Galaxy and may require extreme modifications to the size distributions of silicate particles. The fits for the four sightlines reveal values of $\\lambda_{Max}$ ranging from 4800 to 5500 \\AA. These are consistent with average values of $\\lambda_{max}$ measured in the Galaxy and the Magellanic Clouds. The range measured for M31 corresponds to R$_V$ values of 2.7 to 3.1. The range in R$_V$ seen in the Galaxy is 2.5 to 5.5 implying, for this small sample, that the average size of interstellar grains in M31 is typically smaller than that seen for Galactic grains if the nature of the grains is the same. Also, the polarization efficiency for these sightlines is large although some bias is expected since sightlines known to have significant interstellar polarization were selected for the sample. ",
        "introduction": "There is an average Milky Way extinction relation,  depending on only one parameter, the ratio of total-to-selective extinction, R$_V$, which is applicable to a wide range of interstellar dust environments (Cardelli, Clayton, \\& Mathis 1989 (CCM)). Similarly, there is an interstellar polarization relation, correlated with a single parameter, $\\lambda_{max}$, the wavelength at which the interstellar polarization is a maximum (Serkowski, Mathewson, \\& Ford 1975; Clayton et al. 1995; Martin, Clayton, \\& Wolff 1999).  The parameters, R$_V$ and $\\lambda_{max}$, both involve integrals over the size distribution of interstellar dust grains and are well correlated with each other (Serkowski et al. 1975; Whittet \\& van Breda 1978; Clayton \\& Mathis 1988). The existence of these relations, valid over a large wavelength interval, suggests that the environmental processes that modify the grains are efficient and affect all grains. However, the CCM relation is not a universal law.  It does not fit the observed extinction in the sightlines observed in the Magellanic Clouds (e.g., Clayton et al. 1996; Gordon \\& Clayton 1998; Misselt, Clayton, \\& Gordon 1999; Gordon et al. 2003). The differences in observed extinction properties seen among the Local Group galaxies may be due to various environmental factors including metallicity. The metallicity of the LMC and SMC is lower than the Galaxy. Therefore, the nature of the dust in M31 is of great interest because its metallicity is similar to that in the Milky Way.  The nature of interstellar dust in M31 is summarized in Lequeux (2000). The UV extinction data for M31, which are based on only a few lightly reddened sightlines, show a CCM-like far-UV extinction and a weak 2175 A bump (Bianchi et al. 1996). Several studies have suggested that the extinction in M31 in the U-band differs from the Galaxy (Searle 1983; Iye \\& Richter 1985; Massey et al. 1995). It has also been suggested that there is a systematic change in the U-band extinction with distance from the center of M31 (Iye \\& Richter 1985; Freedman \\& Madore 1990). Various extimates of R$_V$ ranging from 2.5 to 3.8 have been made using individual stars and clusters in M31 (Kron \\& Mayall 1960; Sharov \\& Lyutyj 1989; Nedialkov \\& Veltchev 1999). Recently, Barmby et al. (2000) find that the U- to K-band extinction curve in M31, derived from a large number of globular clusters, is consistent with CCM with R$_V$=3.1.   The interstellar polarization in the LMC and SMC shows a similar wavelength dependence to that seen in the Milky Way and can be well fit by the Serkowski relation (Clayton, Martin, \\& Thompson 1983; Clayton et al. 1996; Rodrigues et al. 1997). Some interstellar polarization measurements in M31 have been made using globular clusters as point sources (Hiltner 1958; Martin \\& Shawl 1979, 1982). However,the wavelength dependence of interstellar polarization has been measured along only one line of sight (Martin \\& Shawl 1979).  This polarization measurement is consistent with the Serkowski relationship and a small value of $\\lambda_{max}$. We report further observations of the wavelength dependence of interstellar polarization in M31. ",
        "conclusions": "The interstellar polarization due to dust in M31 was first measured by Hiltner (1958) using globular clusters as point sources. He made unfiltered observations of about 20 clusters finding several with significant amounts of polarization. Filter polarimetry of a similar number of clusters was obtained by Martin \\& Shawl (1979,1982). Only one sightline in their sample, toward S78, had enough wavelength coverage for a Serkowski fit to be made. That sightline showed p$_{max}$ = 1.98$\\pm$0.07 \\%, and $\\lambda_{max}$ = 3500$\\pm$300 \\AA\\ using the Wilking et al. (1980) formulation of the Serkowski relation.  Our new data for S78 and three additional clusters cover approximately the same wavelength range (4100-8300 \\AA) as Martin \\& Shawl but with many more data bins and smaller uncertainties. We performed nonlinear least-squares fits to the foreground corrected data for all four clusters using the Serkowski law (Serkowski et al. 1975; Wilking et al. 1980; Wilking, Lebofsky, \\& Rieke 1982; Whittet et al. 1992). We found that the  K= 1.66$\\lambda_{Max}$ relationship, that holds for Galactic interstellar polarization, does not produce good fits for the M31 sightlines (Whittet et al. 1992). However, the Serkowski law with K as a free parameter does produce good fits for all of the sightlines. The three fit parameters are listed in Table 1. The relationship between K and $\\lambda_{max}$ is shown in Figure 2. Although, we do not have enough sightlines to define the K-$\\lambda_{max}$ relationship for M31, Figure 2 shows that the slope may be steeper than in the Galaxy.  The K-$\\lambda_{max}$ relationship for the M31 globulars is very interesting in that it produces distinctly narrower ``Serkowski curves'' than those of the same $\\lambda_{max}$\\ in the Galaxy.  As has been shown many times in the past (Wolff et al. 1994; Clayton et al. 1995; Martin et al. 1999), grain models have generally been able to reproduce observed Galactic interstellar polarization curves by small changes to the size distribution of aligned grains, alignment efficiency, upper integration limit, etc.  This is definitely not the case for the M31 curves in that similar changes to a distribution of astronomical silicates failed to produce K values of the observed amplitude.  Even more extreme modifications to the size distribution of silicate particles (use of ``flat'' distributions, very narrow size ranges, etc.) tended to produce K values that still remained 30-40\\% too low, at best.  As a result, it would appear that there is something intrinsically different (and interesting) about the grains producing the observed polarization in M31 from those in our Galaxy. We intend to investigate this behavior in much greater depth after our sample of M31 globulars has been increased.  Previously, only S78 had enough wavelength coverage for a Serkowski fit. Our new data, which have better wavelenth coverage and better S/N, show that $\\lambda_{Max}$$\\sim$4775 \\AA\\ for S78, very different from that found by Martin \\& Shawl (1982). Their value of 3500 \\AA\\ is at the extreme low end of $\\lambda_{max}$ values measured in the Galaxy (e.g., Serkowski et al. 1975). The unusual sightline toward HD 210121, which has one of the lowest R$_V$ values ever measured, has $\\lambda_{max}\\sim$3800 \\AA~(Larson, Whittet, \\& Hough 1996). None of the $\\sim$20 sightlines measured in the LMC show $\\lambda_{max}$ values below 4900 \\AA~(Clayton et al. 1983). Only a few sightlines have been observed in the SMC and the signal-to-noise is relatively low but $\\lambda_{max}$ lies in the normal Galactic range (Rodrigues et al. 1997).  The fits for the four sightlines in M31 reveal values of $\\lambda_{max}$ ranging from 4800 to 5500 \\AA. These are very average values of $\\lambda_{max}$ as measured in the Galaxy and the Magellanic Clouds. It has been found that R$_V\\sim$ 5.5$\\lambda_{max}$ (Serkowski et al. 1975, Whittet \\& van Breda 1978; Clayton \\& Mathis 1988). The range of $\\lambda_{max}$ measured in M31 corresponds to R$_V$ values of 2.7 to 3.1. The range in R$_V$ seen in the Galaxy is 2.5 to 5.5. So, at least for this very small sample, M31 does not show evidence for sightlines with greater than average grain size, if the nature of the grains is the same as for the Galaxy. This result is in contrast to the paucity of small grains inferred from IRAS colors (Xu \\& Helou 1994).  The polarizing efficiency of the M31 dust grains can be estimated using the ratio of peak linear polarization to color excess, p$_{max}$/E(B-V). While the quantity is often equated with the degree of grain alignment efficiency, it is also a function of the grain's non-spherical nature as well as the angle between the alignment axis and the line of sight.  However, for brevity, we simply refer to the value as the ``efficiency ratio.'' The maximum value of this ratio found in the Galaxy is 9\\% mag$^{-1}$ (Serkowski et al. 1975; Clayton \\& Cardelli 1988). We corrected the estimated reddenings for a Galactic foreground of E(B-V)=0.1 mag (Barmby et al. 2000). The values of p$_{max}$/E(B-V) for our sample are given in Table 1. They range from 7.1 to 15.3\\% mag$^{-1}$. These values generally support the results of Martin \\& Shawl (1982) and imply a similar polarization efficiency in M31 and the Galaxy. The values of p$_{max}$/E(B-V) for our sample are all near the maximum value but some bias is expected since sightlines known to have significant interstellar polarization were selected for inclusion in this sample. The polarization efficiency toward S233 seems very high, but the amount of reddening is fairly small and the uncertainty is a considerable fraction of the E(B-V) value itself.  Baade 327 is of interest as it is both the most luminous globular cluster known and also the most reddened globular cluster in M31 (Barmby, Perrett, \\& Bridges 2002). van den Bergh (1968) suggested that it is strange that the brightest globular cluster in M31 would also be the most reddened. This coincidence led to speculation that perhaps the value of R$_V$ might be small along this sightline resulting in less total extinction, A$_V$. Our result for Baade 327, $\\lambda_{max}\\sim$ 5450 \\AA\\, implies that the dust along this sightline has the Galactic average, R$_V\\sim$3.1. Even if R$_V$ = 2.5, as suggested by Kron \\& Mayall (1960), Baade 327 only drops to the second most luminous globular cluster (Barmby et al. 2002). They also point out that the most reddened globular clusters lie close to the center of M31 and those close to the center of M31 seem to be brighter. The polarization efficiency toward Baade 327 is tantalizingly high, particularly given its large color excess and its anticipated favorable viewing geometry (i.e., the angle between alignment and line of sight) (Martin \\& Shawl 1979, 1982).  Clearly, additional observations of M31 globular clusters would provide further insight into the dust environment in that galaxy."
    },
    "0403/astro-ph0403136_arXiv.txt": {
        "abstract": "Prompted by the discovery of the accreted Canis Major dwarf galaxy and its associated globular cluster (GC) system (Martin \\etal ), we investigate the contribution of accreted GCs to the Galactic system. The Canis Major GCs, and those associated with the Sagittarius dwarf galaxy, exhibit a range of galactocentric radii, prograde and retrograde motions, and horizontal branch morphologies, indicating that such properties are of limited use in identifying accreted GCs. By contrast, we find that the age-metallicity relation (AMR) of these dwarf galaxies is distinct from that of the main Galactic GC distribution at intermediate-to-high metallicities ([Fe/H] $\\ga -1.3$). The accretion of GCs with a distinct AMR would explain much of the apparent age spread in the Galactic GC system. The Canis Major and Sagittarius AMRs are similar to those of other Local Group dwarf galaxies and are consistent with a simple closed-box chemical enrichment model -- a further indication that these GCs formed outside of the Milky Way. The Canis Major GCs all have smaller-than-average sizes for their galactocentric distances, lending further support to their origin outside of the Milky Way. Our findings suggest that accretion of similar mass dwarfs does not appear to have played a major role in building the stellar mass of the thick disk or bulge of the Milky Way. ",
        "introduction": "The seminal paper of Searle \\& Zinn (1978) argued that the formation of the Milky Way galaxy was clumpy and rather chaotic. Furthermore, they suggested that some Galactic globular clusters (GCs) may not have formed {\\it in situ} but rather in proto-galactic fragments or dwarf galaxies that were later accreted.  Several workers have attempted to identify candidates for these accreted GCs. Methods have included the retrograde motion of some metal-poor GCs (Rodgers \\& Paltoglou 1984), the variation of horizontal branch (HB) morphology with metallicity in the outer halo (Zinn 1993), the Oosterhoff class (van den Bergh 1993), the size-perigalactic distance relation (van den Bergh 1995) and associations in phase-space (Lynden-Bell \\& Lynden-Bell 1995).  This subject was given impetus with the discovery of the accreted Sagittarius (Sgr) dwarf galaxy and its system of globular clusters (Ibata, Gilmore \\& Irwin 1995).  It has four clearly identified GCs (Terzan 7, Terzan 8, Arp 2 and M54), with M54 as a prime candidate for the nucleus of the disrupted dwarf (Layden \\& Sarajedini 2000). Comparing the phase-space distribution of GCs with the expected orbital path of the Sgr dwarf, Bellazzini \\etal~(2003a) confirmed the status of these four GCs and identified a number of other possible GCs. The closest two GCs to the orbital path are Pal 12 and NGC 4147. Bellazzini \\etal~(2003b) found evidence for Sgr dwarf `tidal debris' (i.e. stripped stars) around these GCs further supporting their association with the Sgr galaxy. For the third closest GC, Pal 5, Odenkirchen \\etal (2003) concluded that it was {\\it not} part of the Sgr dwarf as it is on a different orbital plane. Thus, we have a sample of six GCs which can be confidently associated with the Sgr dwarf. For a magnitude of $M_V = -13.8$, this corresponds to a specific frequency of S$_N$ = 18, comparable to that of Fornax dwarf (S$_N$ = 29; Forbes \\etal~2000), the highest S$_N$ galaxy in the Local Group.   Recently, Martin \\etal (2003) presented evidence for a second disrupted dwarf galaxy in the Milky Way (Martin \\etal 2003). The `Monoceros Ring', discovered earlier by Newberg \\etal (2002), is now thought to be the tidal stream of this disrupted dwarf (whose nucleus lies in the direction of Canis Major). Crane \\etal (2003) and Frinchaboy \\etal (2004) compared the phase-space distribution of Galactic clusters with M giant stars in the Monoceros ring. Assuming a typical thick disk velocity dispersion , they identified five GCs candidates of the disrupted dwarf. They are  NGC~2298, NGC~2808, NGC~5286, Pal~1 and BH~176 respectively.  In the Canis Major discovery paper, Martin \\etal~(2003) included an N-body simulation of the dwarf's orbit over the last 2 Gyrs. They simulated both a prograde and retrograde orbit. Although the prograde orbit is favoured, the retrograde orbit could not be ruled out. Further observational constraints and a non-static Galactic potential in the model should help to confirm which orbit is correct. Several GCs were found to have phase-space distributions in common with both models, which led Martin \\etal (2003) to conclude they were associated with the Canis Major dwarf. They are: NGC 1851, NGC 1904, NGC 2298 and NGC 2808. For an assumed magnitude similar to the Sgr dwarf (Martin \\etal~2003), the four GCs translate into a specific frequency of S$_N$ = 12. In a follow-up paper, Bellazzini \\etal (2003c) argued that the spatial position and stellar populations of the old open clusters AM-2 and Tom~2 imply that they too are associated with the Canis Major dwarf.  Here we derive an updated age-metallicity relation for the Sgr dwarf galaxy and compare it to the Canis Major and Monoceros ring cluster systems. We show that the Canis Major GCs reveal well-defined age-metallicity and size-galactocentric distance relations. Finally, we briefly discuss the possible thick disk/bulge association of the Canis Major GCs and assess the overall importance of accretion in building the stellar mass of the Milky Way. ",
        "conclusions": "The identification of globular clusters (GCs) with the accreted dwarf galaxies Sgr and Canis Major has highlighted the fact that GC properties such as the horizontal branch morphology, prograde or retrograde orbits, range in cluster magnitudes, metallicity and galactocentric distance do not provide a unique signature of an accreted GC.  We have examined the relation between half-mass size and galactocentric distance for the Canis Major GC system, finding that all of the associated GCs have smaller than average sizes. This suggests an origin outside of the Milky Way.   We have derived a simple closed-box age-metallicity relation that provides a good representation of the chemical enrichment history of the Sgr and Canis Major dwarf galaxies. This is consistent with formation of these GCs outside of the Milky Way. The model AMR deviates from the Milky Way GC distribution at intermediate-to-high metallicities, thus providing an alternative and fairly robust method of identifying accreted metal-rich GCs within the Milky Way GC system. Based on the literature age and metallicity measurements for the GC BH~176, we argue it is unlikely to be a former member of the Canis Major galaxy.  We support earlier suggestions that the younger, metal-poor Milky Way GCs are prime candidates for accreted GCs from as yet unidentified galaxies. This would reduce the GC age spread, at a given metallicity, to less than a Gyr and imply that any halo collapse was very rapid. As the age-metallicity relation of the known accreted GCs is distinct from that of other Milky Way GCs for [Fe/H] $\\ge -1.3$, it suggests that accretion was not a major factor in building the stellar mass of the thick disk or bulge. The majority of the Milky Way GC system and, by implication, the Galaxy itself formed {\\it in situ}."
    },
    "0403/astro-ph0403300_arXiv.txt": {
        "abstract": "Diffuse gamma rays probe the highest-energy processes at the largest scales. Here we derive model-independent constraints on the hadronic contribution to the Galactic and extragalactic $\\gamma$-ray spectra at in the energy range $50 \\ {\\rm MeV} \\la E_\\gamma \\la 10 \\ {\\rm GeV}$. The hadronic component is dominated by emission from neutral pions, with a characteristic spectrum symmetric about $m_{\\pi^0}/2$. We exploit the well-defined properties of the pion decay spectrum to quantify the maximum pionic fraction of the observed $\\gamma$-ray intensity. We find that the Galactic spectrum above 30 MeV can be at most about 50\\% pionic. The maximum pionic contribution to the extragalactic spectrum is energy dependent; it also depends on the redshift range over which the sources are distributed, ranging from as low as about 20\\% for pions generated very recently, to as much as 90\\% if the pions are generated around redshift 10. The implications of these constraints for models of $\\gamma$-ray and neutrino emission are briefly discussed. ",
        "introduction": "The prominence of diffuse emission in the $\\gamma$-ray sky above $\\ga 50$ MeV has been known since the earliest days of $\\gamma$-ray astronomy itself \\citep{fkh}. These diffuse photons carry unique and direct information about some of the most energetic sites and processes in nature. Diffuse $\\gamma$-ray observations thus provide a powerful tool both (1) to test specific models of known or postulated astrophysical sources, and (2) to constrain, in a model-independent way, known physical processes which might occur in one or more sources. We take the latter approach in this paper, focusing in particular on the $\\gamma$-ray spectrum and the constraints it places on the contribution of hadronic interactions to the overall diffuse background.  The diffuse $\\gamma$-ray sky is dominated by emission from the Galactic plane \\citep{hunter}, but the presence of emission even at the Galactic poles already suggests that an extragalactic component is present as well \\citep{sreekumar}. The spectra of these two components are each remarkable both for what they show and what they do not show. Namely, in neither spectrum is there a strong indication of hadronic interactions, which are dominated by proton collisions with interstellar matter, which yield $\\gamma$-rays predominantly through pion production and decay: $pp \\rightarrow pp\\pi^0 \\rightarrow \\gamma\\gamma$. The pionic spectrum is symmetric about a peak at $m_\\pi/2$. This feature, the ``pion bump,'' is notably inconspicuous in the $\\gamma$-ray data.   As we will see in detail below, the Galactic spectrum is well-described by a simple broken power law, with a break at $\\sim 0.77$ GeV. No strong pion bump is observed. \\citet{hunter} do note that there is as a $\\sim 2\\sigma$ deviation in the $60-70$ MeV energy bin, but this region in the spectrum is otherwise well-fit by a smooth power law. If real, this feature is remarkably narrow. Intriguingly, detailed models of known Galactic processes run into difficulties explaining this spectrum (and its simplicity). The model of \\citet{smr00} includes a sophisticated 2-D model of the cosmic-ray, gas, and photon fields in the Galaxy, and includes hadronic interactions, electron bremsstrahlung and inverse Compton scattering of starlight. However, when using only known cosmic ray populations and spectra, this model is unable to account for the observed $\\gamma$-ray spectrum. The spectrum above about 1 GeV is flatter than the prediction of pionic emission, so other sources seem to be required as well.  Proposed explanations for this ``GeV excess'' include modifications to the proton spectrum, and additional inverse Compton radiation due to an extended halo of cosmic ray electrons \\citep{smr00}. One of the main goals of this paper is to quantify the portion that can be pionic.  Information about the extragalactic component of diffuse $\\gamma$-rays is more difficult to obtain, as one must first subtract the Galactic foreground, which is large at low--and possibly even high--Galactic latitudes. As we will see, the nature of the extragalactic spectrum depends on the method used to subtract the Galactic foreground. Different techniques have recently emerged, leading to different results for the shape and amplitude of the spectrum. \\citet{sreekumar} find a single power-law, while \\citet{smr03} find a smaller but ``convex'' spectrum. In either case, no pion bump is seen.  Many astrophysical sites have been proposed to explain the extragalactic emission. These necessarily include ``guaranteed'' sources, namely, active \\citep{ss,mc} and normal \\citep{pf} galaxies. These are the classes of objects which have been directly detected in nearby objects, but which would be unresolved when at large distances. Also, there is a growing consensus that the formation of large scale structures leads to shocks in the baryonic gas, and thus to particle acceleration. The resulting ``cosmological cosmic rays'' have recently become the subject of intense interest \\citep{mrkjco,lw,miniati,ti,kwlsh,bd,fl}. These can also contribute to the diffuse $\\gamma$-background, and would have emission from both hadronic and inverse Compton processes.  Here we wish to find model-independent constraints on hadronic and thus pionic emission mechanisms. We choose to focus on this component because its detection would finally confirm observationally the theoretical expectation that the same astrophysical acceleration processes which give rise to non-thermal electrons (and associated synchrotron radiation) also give rise to non-thermal ions. Also, we wish to exploit the well-defined nature of the pion decay spectrum which allows us to make a roughly model-independent comparison with observations. Finally, since the same hadronic processes that produce neutral pions also produced charged pions and hence neutrinos, our limits will have implications for neutrino production as well. ",
        "conclusions": "We have presented model-independent upper limits on hadronic $\\gamma$-ray emission based on the shape of observed spectra and their lack of a pion bump. Above 100 MeV, one might expect that gamma-rays from $\\pi^0$ decay should dominate the Galactic spectrum. However, we find that they can make only about 50\\% of total Galactic gamma-ray flux. From the shape of the residual function that is plotted in Fig.\\ref{fig:mw} we can see that due to the break at 0.77 GeV, additional gamma-ray components are needed both above and below the $\\sim 250$ MeV regime at which the pionic contribution can be maximized. The residual gamma-rays below this scale can possibly be accounted by bremsstrahlung and inverse Compton scattering \\citep{smr00}. However, if the pionic component is near its maximum, there still is the need for at one more component  of gamma-rays above $\\sim 300$ MeV to account for the EGRET data; alternatively, a single process may be at work, but with pionic emission being sub-dominant at all energies. This model-independent result is in good agreement with the sophisticated analysis of different models by \\citet{smr00}. We also stress again that this analysis was based on the assumption that the pion bump is not observed. If a future $\\gamma$-ray mission such as GLAST were to identify this feature that would allow us to set a more definite and stronger limit on the pionic fraction of diffuse gamma-rays.  The maximum pionic fraction of extragalactic $\\gamma$-rays can be seen in Fig.\\ref{fig:fraction} for different methods of foreground subtraction. It can go from $20-90\\%$, depending on the assumed redshift of cosmic ray origin, but also on method used in subtracting the Galactic foreground to obtain the EGRB. Namely, the \\citet{smr03} fit gives the fraction of  20\\% for cosmic rays that originated at the present, up to about 70\\% for $z=10$, and in both cases there is still a factor of $\\ga 2$ deficit of high-energy photons. On the other hand, the \\citet{sreekumar} fit implies a 20\\% pionic fraction for recent cosmic rays and about 90\\% for redshift 10 cosmic rays, with no significant deficit at high energies.. This large variation underscores the need for a robust procedure for determining the EGRB, and also emphasizes the power of a firm EGRB spectrum to constrain emission processes.  Our limits can be compared to the results of specific models. \\citet{miniati} finds that the pionic component contributes about 30\\% of the total emission from structure-formation cosmic rays, with the balance arising from electron synchrotron emission.  He in turn finds that the entire cosmic-ray component itself can be $\\sim 20-30\\%$ of the total \\citep{sreekumar} observed background.  Thus, the \\citet{miniati} model finds that the pionic contribution is $\\sim 6-10\\%$ of the observed level, a fraction which would be larger if the smaller \\citet{smr03} background were used. The work of \\citet{kwlsh} neglects pionic emission entirely, arguing that the synchrotron emission should dominate.  Of course, our extragalactic constraints reflect our simple single-redshift approximation for the origin of cosmic rays. More precise analysis should include some kind of averaging over the redshifts. Further progress first requires a detailed knowledge of the redshift evolution of sources and targets. On the other hand, if the ``pionic bump'' was observed in the EGRB spectrum by GLAST or other future experiments then the position of its peak could immediately tell us something about the mean cosmic-ray flux. This information would then give us a better handle on the star-formation rate as a function of redshift.  Finally, we note that our constraints on photons of hadronic origin also have implications for neutrinos produced in the same processes, a connection which has been emphasized by \\citet{wb}. This is of particular interest in the case of extragalactic neutrino emission, which may lead to high-energy ($E_\\nu \\ga 1$ TeV) events observable by ICECUBE \\citep{icecube}. \\citet{wb} use the energy density of ultra-high-energy ($\\ga 10^{19}$ eV) cosmic rays to derive limits on the high-energy cosmic neutrino flux. Furthermore, they find that the EGRB (in the power-law form of \\citet{sreekumar}) implies a neutrino flux above $\\sim $ TeV which violates this limit by a factor of up to $\\sim 100$, if the EGRB is of hadronic origin. Our limits on the pionic fraction of the EGRB are complementary to the \\citet{wb} result. Our constraints on the hadronic origin of the EGRB are weaker than their $\\sim 1\\%$ fraction, but are derived independently, based on the EGRB spectrum itself."
    },
    "0403/astro-ph0403015_arXiv.txt": {
        "abstract": "{The polarization of the optical afterglow (AG)  of Gamma-Ray Bursts (GRBs) has only been measured in a few instances at various times after the GRB. In all cases except the best measured one (GRB 030329) the observed polarization and its evolution are simple and easy to explain in the most naive version of the ``Cannonball'' model of GRBs:  the ``intrinsic'' AG polarization is small  and the observations reflect the ``foreground'' effects of the host galaxy and ours. The polarization observed in GRB 030329 behaves chaotically, its understanding requires reasonable but ad-hoc ingredients. The polarization of the $\\gamma$-rays of a GRB has only been measured in the case of GRB 021206. The result is debated, but similar measurements would be crucial to the determination of the GRB-generating mechanism. ",
        "introduction": "Spectropolarimetric measurements of radiations from astronomical objects are an important diagnostic tool of their production mechanism. Gamma Ray Bursts (GRBs) are not an exception. For these phenomena one must distinguish between two observables: the polarization of the ``prompt'' $\\gamma$-rays of the GRB itself, which has been measured in just one instance (GRB 021206; Coburn and Boggs 2003) and the polarization of the GRB afterglows (AGs), observed at optical frequencies in a handful of cases (GRB 990510: Wijers et al.~1999; Covino et al.~1999; GRB 990712: Rol et al.~2000; GRB 010222: Bjornsson et al.~2001; GRB 011211: Covino et al.~2002; GRB 020405: Bersier et al.~2002; Masetti et al.~2003; Covino et al.~2003a; GRB 020813: Barth et al.~2003; Covino et al.~2003b; Gorosabel et al.~2003; GRB 021004: Rol et al.~2003; Wang et al.~2003; GRB 030329: Efimov et al.~2003;  Magalhaes et al.~2003; Covino et al.~2003c; Greiner et al. 2003).   In this paper we are primarily concerned with the polarization of optical AGs, though we first comment on that of a GRB itself. We shall conclude that, while (difficult) convincing measurements of the polarization of GRBs would be decisive in establishing the {\\it mechanism} generating GRBs, the measurement of the polarization of AGs is unlikely (at least in the CB model) to shed much light on their understanding. ",
        "conclusions": "Examining the results shown in Figs.~(\\ref{POL0208133}) to (\\ref{329b}), we conclude that, for GRBs 020405, 020813 and 021004, the fits ---which did not include an {\\it intrinsic} polarization of the light emanating from the CBs--- are good enough to conclude that this contribution vanishes within errors. As required in that case, the polarization and angle tend at late times to the Galactic foreground values. This is particularly convincing for GRB 020813. The most naive expectation ---that the observed polarization simply reflects the foreground effects of the host galaxy and ours--- is vindicated.  The case of GRB 030329, for which the data are particularly precise and abundant, brings havoc to the previous clear and simple conclusion. This is the case even if one neglects the data after day $\\sim 6$, probably contaminated by the parent SN, which is particularly prominent in this instance.  In the CB-model, the explanation for the ``humps\" in the AGs of some GRBs, such as GRBs 970508 and 000301c (Dado et al.~2002a) and 030329 itself (Dado et al.~2004), is simple: the optical fluence $F_\\nu(t)$ is a direct and {\\it quasi-local} tracer of the density of the ISM through which a CB travels: spatial changes in density translate into temporal changes in fluence. This statement is not  as sterile as it sounds, for it permits the extraction of the circumburst density and its radial profile from the time-dependence of the early AGs and, very satisfactorily, the magnitude and $\\sim 1/r^2$ profile of the result are those expected from observations of ``winds'' of ``pre-supernova'' massive stars (Dado et al.~2003b).  In the case of GRB 030329, the observed deviations of the AG light curves from smoothly-varying functions are attributed to density fluctuations encountered by the CBs at the time they exit from the superbubble in which the explosion took place. The size and shape of these density fluctuations can be explicitly extracted from the data: they are a series of density jumps followed by $\\propto 1/r^2$ declines, as befits the remnants of the explosions that created the superbubble (Dado et al.~2004). Their magnitude and shape (as functions of time) are shown in Fig.~(\\ref{overdensity}).  Fluctuations of the host's ISM density though which a CB travels may also cause the polarization fluctuations observed in GRB 030329. This may be a foreground ``integral'' effect (induced by the varying amounts and field directions of magnetized dust in the complicated density profile along the line of sight to the observer). It  may also be a local ``intrinsic\" effect: the magnetic field within a CB may not be, at the few percent level, totally chaotic. It may be influenced by fluctuations in the density of the ISM particles that impinge into it and generate its magnetic structure. There appears to be no clearly convincing correlation between the fluctuations in the fluence and those in the polarization, though the periods of rapid variability coincide: compare Figs.~(\\ref{329a}) and (\\ref{329b}) to Fig.~(\\ref{overdensity}). This inclines the balance somewhat in favour of an integral foreground effect, as if indeed there were no intrinsic polarization in the radiation emanating from the CBs, as expected for synchrotron radiation in a thoroughly disordered magnetic field.  Admittedly, the considerations of the previous paragraph are not simple and robust. They drive the conclusion that, in the CB model, no much is to be learned from the polarization of optical AGs. This is in contrast to the polarization of the $\\gamma$-rays of a GRB, which, we contend, is crucial for deciding what the GRB-generating mechanism is: inverse Compton scattering if the polarization is measurably large.  \\vskip .3cm   {\\bf Acknowledgments}  The authors are grateful to J. Greiner for providing tabulated data on the polarization of the AG of GRB030329. S. Dado and A. Dar are indebted to the theory Division of CERN for hospitality. A. De R\\'ujula is indebted to the Physics Department and Space Research Institute of the Technion for its hospitality.  This research was supported in part by the Helen Asher Space Research Fund for research at the Technion."
    },
    "0403/astro-ph0403709_arXiv.txt": {
        "abstract": "To first order, the initial cluster luminosity function appears to be universal. This means that the brightest young cluster in a galaxy can be predicted from the total number of young clusters based purely on statistics. This suggests that the physical processes responsible for the formation of clusters are similar in a wide variety of galaxies, from mergers to quiescent spirals. One possibility is that conditions for making young massive clusters are globally present in mergers while only locally present in spirals (i.e., in the spiral arms). However, understanding the destruction of clusters and the accompanying survival rates is more important for understanding cluster demographics than understanding their formation. This is because only about 1 in 1,000 clusters with mass greater than 10$^4$ \\mdot\\/ will survive to become an old globular cluster.  In this paper we briefly review this basic framework and then develop a toy model that allows us to begin to address several fundamental questions.  In particular, we demonstrate that young clusters in the Antennae Galaxies have a high ``infant mortality'' rate, with roughly 90 \\% of the clusters being destroyed each decade of log(time). We also advocate the use of an objective classification system for clusters, with the three parameters being mass, age, and size. ",
        "introduction": "Historically, star clusters have been divided into three types in the Milky Way; globular clusters, open clusters, and associations. This has conditioned us to think in terms of three distinct modes of cluster formation. However, when we look at external galaxies we see a power law continuum of young cluster masses. How can these two outlooks be reconciled?  There is growing evidence  that all star clusters may form from a universal initial cluster mass function, which is then modified by a variety of destruction mechanisms. These destruction mechanisms include infant mortality (e.g., unbound clusters that freely expand in the first $\\approx$ 10 Myr), environmental effects (e.g., disk and bulge shocking, dynamical friction), internal effects (e.g., evaporation from 2-body relaxation ) and stellar evolution (e.g., mass loss). By convolving these formation and destruction rates with the star formation history of a galaxy, and after taking into effect various observational artifacts and selection effects, we should be able to predict the current distribution of young and old star clusters we see in a particular galaxy.  In this contribution we first outline some of the results which have lead to the development of this basic framework, and then develop a ``toy'' model designed to help answer three fundamental questions.    \\subsection{Is the Initial Cluster Mass Function Continuous or Modal ?}  Essentially all studies of the luminosity functions of young clusters have found them to be power laws, $\\phi(L)dL \\propto L^{\\alpha} dL$ with a value of $\\alpha$ $\\approx$ --2 (e.g., Whitmore et al. 1999, see Whitmore 2003 for a review [originally appearing in 2000 as astro-ph/0012546]). Mass functions have been determined for only a small subset of these galaxies, but a similar power law relationship is generally found, again with an index --2 (e.g., Zhang and Fall 1999). On the other hand, the luminosity and mass functions of old globular clusters are peaked, with a mean $V$ magnitude M$_v$ $\\approx$  -7.4, a mean mass $\\approx$ 2 $\\times$ 10$^5$ \\mdot, and a width $\\sigma \\approx 1.4$ mag (e.g., Whitmore 2003). Several theoretical studies (e.g., Vesperini 1998, Fall and Zhang 2001) have suggested that a natural way to explain this apparent discrepancy is the destruction of the fainter, less massive clusters, due to effects such as 2-body relaxation, tidal shocks, and stellar mass loss.  Other theoretical studies have also highlighted the important role that cluster destruction may play in determining the demographics of clusters (e.g., Fall \\& Rees 1977, Gnedin \\& Ostriker 1997). The current paper adopts a similar framework, but broadens it to include all star clusters; young and old, near and far. The basic question is: {\\it ``Can all observations of star cluster demographics be explained by a universal initial cluster mass function followed by the destruction of a subset of the clusters which carves away regions of parameter space ?''}     \\subsection{Is the Initial Cluster Mass Function Universal ?}   On the surface, the demographics of star clusters appear to differ dramatically in different galaxies.  For example, merging galaxies have large numbers of very bright, very massive young clusters, often called super star clusters. On the other hand, in relatively quiescent galaxies, such as the Milky Way, the young clusters appear to be primarily faint and low mass. We might conclude that certain types of clusters can only be formed in certain types of galaxies. This will be referred to as ``special creation'' in this paper.     An alternative to ``special creation'' is suggested by the work of Elmegreen and Efremov (1997), Whitmore (2003), and Larsen (2002), who all suggest that the initial cluster mass function may be universal. Whitmore (2003) examined the luminosity functions of eight galaxies as part of his review of the formation of star clusters. Part of the motivation was to look for a truncation at the bright end of the luminosity function  which might indicate that only the major mergers can produce the very massive clusters (i.e., supporting ``special creation''). However, he found that the cluster luminosity functions  for mergers, starbursts, and barred galaxies were all power laws. The main difference is the normalization of the power law, with young active mergers having thousands of young clusters, while other systems  have hundreds of young clusters.  The large difference in numbers might indicate that conditions for making clusters are {\\it globally} present in mergers but only {\\it locally} present (e.g., in spiral arms) in more quiescent galaxies.   Since the number of galaxies with sufficient data to form a meaningful luminosity function was limited, Whitmore (2003) went on to use a poor-man's version of the diagram, shown in Figure 1, to enlarge the sample.  This includes the sample of spiral galaxies from Larsen \\& Richtler (2000) and plots only the luminosity of the brightest cluster vs. the log of the number of clusters in a galaxy that are brighter than M$_V$ = --8. Somewhat surprisingly, this shows no    \\clearpage  \\begin{figure}[hb]\t% \\plotfiddle{whitmore_fig1.ps}{5.5 cm}{0}{35}{35}{-100}{-80} \\vspace{2 cm} \\caption{Magnitude of the brightest cluster vs. log of the number of clusters in 34 galaxies. Filled circles are spiral galaxies from Larsen \\& Richtler (2000), open circles are mergers, stars are starbursting galaxies, and the half-filled square is a barred galaxy. The solid line is a best fit (excluding NGC 1569) while the dashed line is the prediction from a universal power law luminosity function with index $\\alpha$ = --2. See Whitmore (2003) for details.} \\end{figure}      \\noindent evidence of a truncation in the ability of quiescent spiral galaxies to form bright star clusters. The observed slope is consistent with nearly all of the galaxies having the same universal luminosity function. This suggests that active mergers have the brightest clusters only because they have the most clusters.  {\\it It appears to be a matter of simple statistics rather than a difference in physical formation mechanisms.}    Larsen (2002) has further developed this basic idea by performing Monte-Carlo simulations that show that the spread in the scatter of the relation can also be explained by statistics. Several other recent studies have also begun to consider this ``size-of-sample'' effect   (e.g., Hunter et al., 2003, and Billet, Hunter, \\& Elmegreen 2002).     \\subsection{What Fraction of Stars are Formed in Clusters and What Fraction are Formed in the Field ?}  Understanding the destruction of star clusters is more important for understanding their demographics than understanding how they form. For every cluster that survives to an age of 10 Gyr, roughly a thousand were created and have been destroyed, their stars being dispersed into the field. As a result of this, it now appears that the majority of stars are created in clusters rather than in the field.  Whitmore (2003) points out that if we assume that the Antennae has been making clusters at the same rate for the past 200 Myr (a rather uncertain assumption), it is possible to use Figure 2 from Zhang and Fall (1999, an updated figure is also shown as Figure 2c below),  to show that for every 20 clusters in the age range 0 - 10 Myr, only one will survive to an age of $\\approx$ 100 Myr (i.e., there are roughly the same number of clusters in the 0 - 10 Myr age bin as in the 20 - 200 Myr age bin). While this crude calculation is probably not justifiable for a single galaxy, larger samples are now becoming available (e.g., M83 - Harris et al. 2001, LMC - Hunter et al. 2003, NGC 6745 - de Grijs et al. 2003, M51 and M101 - Chandar et al. 2004) which show similar demographics in nearly all star forming galaxies that have been observed. ",
        "conclusions": ""
    },
    "0403/astro-ph0403223_arXiv.txt": {
        "abstract": "We determine the abundance of $i'$-band drop-outs in the recently-released {\\em HST}/ACS Hubble Ultra Deep Field (UDF). Since the majority of these sources are likely to be $z\\approx 6$ galaxies whose flux decrement between the F775W $i'$-band and F850LP $z'$-band arises from Lyman-alpha absorption, the number of detected candidates provides a valuable upper limit to the unextincted star formation rate at this redshift. We demonstrate that the increased depth of UDF enables us to reach an $8\\,\\sigma$ limiting magnitude of $z'_{AB}=28.5$ (equivalent to $1.5\\,h^{-2}_{70}\\,M_{\\odot}\\,{\\rm yr}^{-1}$ at $z=6.1$, or $0.1\\,L^{*}_{UV}$ for the $z\\approx 3$ $U$-drop population), permitting us to address earlier ambiguities arising from the unobserved form of the luminosity function. We identify 54 galaxies (and only one star) at $z'_{AB}<28.5$ with $(i'-z')_{AB}>1.3$ over the deepest 11\\,arcmin$^{2}$ portion of the UDF field. The characteristic luminosity ($L^*$) is consistent with values observed at $z\\approx 3$. The faint end slope ($\\alpha$) is less well constrained, but is consistent with only modest evolution. The main change appears to be in the number density ($\\Phi^{*}$). Specifically, and regardless of possible contamination from cool stars and lower redshift sources, the UDF data support our previous result that the star formation rate at $z\\approx 6$ was approximately $\\times 6$ {\\em less} than at $z\\approx 3$ (Stanway, Bunker \\& McMahon 2003). This declining comoving star formation rate ($0.005\\,h_{70}\\,M_{\\odot}\\,{\\rm yr}^{-1}\\,{\\rm Mpc}^{-3}$ at $z\\approx 6$ at $L_{UV}>0.1\\,L^{*}$ for a salpeter IMF) poses an interesting challenge for models which suggest that $L_{UV}>0.1\\,L^{*}$ star forming galaxies at $z\\simeq 6$ reionized the universe. The short-fall in ionizing photons might be alleviated by galaxies fainter than our limit, or a radically different IMF. Alternatively, the bulk of reionization might have occurred at $z\\gg6$. ",
        "introduction": "\\label{sec:intro}  There has been considerable progress over the past decade in locating galaxies and QSOs at high redshifts. These sources have enabled us to probe the Universe at early epochs where its physical characteristics are fundamentally different from those at the present epoch. Observations of the most distant $z>6.2$ QSOs (Becker et al.\\ 2001, Fan et al.\\ 2002) show near-complete absorption at wavelengths shortward of Lyman-$\\alpha$ (Gunn \\& Peterson 1965), suggesting an optical depth in this line that implies a smooth neutral hydrogen fraction which is increasing rapidly with redshift at this epoch. Temperature-polarization cross-correlations in the cosmic microwave background from  {\\em WMAP} indicate that the Universe was completely neutral at redshifts of $z> 10$ (Kogut et al.\\ 2003).  Although there is a growing consensus that cosmic reionization occurred in the redshift interval $6<z<15$, a second key question is the nature of the sources responsible for this landmark event. Optical and X-ray studies to $z\\simeq 6$ suggest the abundance of active galactic nuclei (AGN) at early epochs is insufficient when account is taken of the relevant unresolved backgrounds (Barger et al.\\ 2003). A more promising source is star-forming galaxies whose early ancestors may be small and numerous. Along with the escape fraction for the ionizing photons from the massive and short-lived OB stars in such sources, a major observational quest in this respect is a determination of the global star formation rate at early epochs.  In previous papers, our group has extended the {\\it Lyman-break} technique (Steidel, Pettini \\& Hamilton 1995; Steidel et al.\\ 1996) to address this question. Using the Advanced Camera for Surveys (ACS, Ford et al.\\ 2002) on the Hubble Space Telescope ({\\em HST}) with the sharp-sided SDSS F775W ($i'$) and F850LP ($z'$) filters, we located ``$i$-drop'' candidates with $z_{AB}'<$25.6 at $z\\simeq 6$ for further study. In a series of papers, we have shown that this selection technique can effectively locate $z>5.7$ galaxies using ACS images from the {\\em HST} Treasury ``Great Observatory Origins Deep Survey'' (GOODS; Giavalisco \\& Dickinson 2002). On the basis of GOODS-South photometric catalogues published by Stanway, Bunker \\& McMahon (2003, hereafter Paper I), spectroscopic follow-up using Keck/DEIMOS and Gemini/GMOS field demonstrated our ability to find high redshift galaxies (Bunker et al.\\ 2003, hereafter Paper II; Stanway et al.\\ 2004a). To address potential cosmic variance issues, we performed a similar analysis in the GOODS-North field, which yielded a consistent estimate of the surface density of $z\\simeq 6$ star forming sources (Stanway et al.\\ 2004b, hereafter Paper III).  Although our initial study (Papers I-III) revealed the importance of ascertaining the difficult spectroscopic verifications, and highlighted the problem of contamination from Galactic stars, we nonetheless determined that the abundance of confirmed star forming galaxies at $z\\sim 6$ must be less than that expected on the basis of no evolution from the well-studied $z\\sim 3-4$ Lyman break population (Steidel et al.\\ 1999).  Working at the robustly-detected bright end of the luminosity function, in Paper~I we showed that the comoving star formation density in galaxies with $z'_{AB}<25.6$ is $\\approx 6\\times$ {\\em less} at $z\\approx 6$ than at $z\\approx 3$.  Our $z'_{AB}<25.6$ flux limit corresponds to $>15\\,h^{-2}_{70}\\,M_{\\odot}\\,{\\rm yr}^{-1}$ at $z=5.9$, equivalent to $L^{*}_{UV}$ at $z\\approx 3$. In Papers I--III we restricted our analysis to luminous galaxies (where we take ``luminous'' to mean $L>L*$ for the rest-UV). Other groups have claimed less dramatic evolution or even no evolution in the volume-averaged star formation rate, based on the same fields (Giavalisco et al.\\ 2004; Dickinson et al.\\ 2004) and similar {\\em HST}/ACS data sets (Bouwens et al.\\ 2003; Yan, Windhorst \\& Cohen 2003), but these groups work closer to the detection limit of the images and introduce large completeness corrections for the faint source counts. The major uncertainty in converting the abundance of our spectroscopically-confirmed sample in the GOODS fields into a $z\\simeq 6$ comoving star formation rate is the form of the luminosity function for faint, unobserved, sources. As discussed in Paper~III, if the faint end of the luminosity function at $z\\simeq 6$ was steeper than that at lower redshift, or if $L^{*}$ was significantly fainter, a non-evolving star formation history could perhaps still be retrieved.  The public availability of the Hubble Ultra Deep Field (UDF; Beckwith, Somerville \\& Stiavelli 2003) enables us to address this outstanding uncertainty. By pushing the counts and the inferred luminosity function of $i'$-band drop-outs at $z\\approx 6$ to a limiting lower luminosity equivalent to one well below $L^{*}_{3}$ for the $z\\approx 3$ population, it is possible to refine the integrated star formation rate at $z\\approx 6$. In this paper we set out to undertake the first photometric analysis of $i'$-drops in the UDF. Our primary goal is to understand the abundance of fainter objects with characteristics equivalent to those of $z\\simeq 6$ sources and address uncertainties in the global star formation rate at this redshift.  The structure of the paper is as follows. In Section~\\ref{sec:obs} we describe the imaging data, the construction of our catalogues and our $i'$-drop selection. In Section~\\ref{sec:discuss} we discuss the luminosity function of star-forming sources, likely contamination on the basis of earlier spectroscopic work, and estimate the density of star formation at $z\\approx 6$. Our conclusions are presented in Section~\\ref{sec:concs}. Throughout we adopt the standard ``concordance'' cosmology of $\\Omega_M=0.3$, $\\Omega_{\\Lambda}=0.7$, and use $h_{70}=H_0/70\\,{\\rm km\\,s^{-1}\\,Mpc^{-1}}$. All magnitudes are on the $AB$ system (Oke \\& Gunn 1983). ",
        "conclusions": "\\label{sec:concs}  We summarize our main conclusions as follows:  \\begin{enumerate}  \\item{We present an $i'$-dropout catalogue of $z\\approx 6$ star forming galaxy candidates in the Ultra Deep Field (UDF) to a limiting flux ($8\\,\\sigma$) of $z'_{AB}<28.5$. This represents a substantial advance over the depths achieved in the GOODS catalogues and enables us, for the first time, to address questions concerning the contribution of the faint end of the luminosity function.}  \\item{We detect 54 resolved sources with $(i'-z')_{AB}>1.3$ in the deepest 11\\,arcmin$^{2}$ portion of the UDF and consider this to be an upper limit to the abundance of star forming galaxies at $z\\approx 6$.}   \\item{Using simulations based on lower redshift data, we deduce that, regardless of contamination by foreground interlopers, the abundance of $i'$-dropouts detected is significantly less than predicted on the basis of no evolution in the comoving star formation rate from $z=3$ to $z=6$ (integrating to the same luminosity level). The UDF data supports our previous suggestions that the star formation rate at $z\\approx 6$ was about $\\times 6$ {\\em less} than at $z\\approx 3$ (Stanway, Bunker \\& McMahon 2003).}  \\item{The inferred comoving star formation rate of $0.005\\,h_{70}\\,M_{\\odot}\\,{\\rm yr}^{-1}\\,{\\rm Mpc}^{-3}$ from $L>0.1\\,L^{*}_{UV}$ galaxies at $z\\approx 6$ may poses a significant challenge for models which require that luminous star forming galaxies in the redshift range 6$<z<$10 are responsible for reionizing the Universe.}  \\item{The contamination of our $i'$-drop sample of candidate $z\\approx 6$ galaxies by cool Galactic stars appears to be minimal at $z'_{AB}>26$, possibly because we are seeing beyond the Galactic disk at the faint magnitudes probed by the UDF.}  \\end{enumerate}  \\subsection*{Note Added in Proof} A recent preprint by Yan \\& Windhorst (astro-ph/0407493) independently repeats our selection of candidate $z\\approx 6$ galaxies in the Hubble Ultra Deep Field  with $(i'-z')_{AB}> 1.3$ (astro-ph/0403223 and this paper). The Yan \\& Windhorst catalogue also pushes to fainter magnitudes than our $z'_{AB}<28.5$ limit, where the completeness corrections become significant. This subsequent independent analysis recovers almost all of our original $i'$-band drop-out galaxies, and the catalogues agree at the 98\\% level (one discrepant object out of 50). In astro-ph/0407562 (Bunker \\& Stanway 2004) we present a matched catalogue of these $i$-band dropouts in the Hubble Ultra Deep Field.  \\subsection*{Acknowledgements}  We thank Steve Beckwith and colleagues at the Space Telescope Science Institute for making the UDF data available as a public database, on schedule and in a manner suitable for immediate analysis. ERS acknowledges a Particle Physics and Astronomy Research Council (PPARC) studentship supporting this study. Based on observations made with the NASA/ESA Hubble Space Telescope, obtained from the Data Archive at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS 5-26555. These observations are associated with program \\#9978. We thanks the anonymous referee for some helpful suggestions."
    },
    "0403/astro-ph0403479_arXiv.txt": {
        "abstract": "IGR J$16320-4751$ was re--discovered by IBIS/ISGRI on board \\emph{INTEGRAL} in early February 2003 during the observation of the black hole candidate 4U$1630-47$ (PI Tomsick). This source, already observed by ASCA and BeppoSAX, belongs to the class of heavily absorbed objects ($N_H > 10^{23}$~cm$^{-2}$) that populate some arms of the Galaxy. Soon after the rediscovery by IBIS/ISGRI, the source was observed by \\emph{XMM-Newton}: the arcsec position found with XMM allowed ones to find the most likely infrared counterpart. We present here the reanalysis of the high energy emission from IGR J$16320-4751$ detected by IBIS/ISGRI, including the spectral and temporal characteristics. We also present a reanalysis of the \\emph{XMM-Newton} and optical/IR data. ",
        "introduction": "IGR J$16320-4751$ was serendipitously discovered on February 1.4, 2003 UT, with the IBIS/ISGRI detector (Ubertini et al. 2003, Lebrun et al. 2003) on board the \\emph{INTEGRAL} satellite during the AO1 observation of the black hole candidate 4U $1630-47$ (PI Tomsick). The coordinates (J2000) were $\\alpha=16^{\\rm h}:32^{\\rm m}{\\!}.0$ and $\\delta=-47^{\\circ}:51^{\\rm m}{\\!}.0$, with an uncertainty of $2'$ (Tomsick et al. 2003a). A quick archival research allowed ones to find immediately the X--ray counterpart in the ASCA catalog, as the source AX J$1631.9-4752$ (Sugizaki et al. 2001). IGR J$16320-4751$ is therefore apparently located in the Norma Arm of the Galaxy, an active star forming region distant from the Sun about $5$ kpc (Georgelin \\& Georgelin 1976, Russeil 2003).  A Target of Opportunity request (ToO) with \\emph{XMM-Newton} was immediately activated and it found one single source inside the error circle of ISGRI. The coordinates (J2000) were $\\alpha=16^{\\rm h}:32^{\\rm m}:01^{\\rm s}{\\!}.9$ and $\\delta=-47^{\\circ}:52^{\\rm m}:29^{\\rm s}$ with an uncertainty of $4''$ (Rodriguez et al. 2003a). In this reduced error circle it was possible to identify two reliable infrared counterparts in the 2MASS catalog (Tomsick et al. 2003b).  A preliminary analysis of the \\emph{XMM-Newton} observation, together with the identification of the infrared counterpart was presented in Rodriguez et al. (2003b). Here we present a reanalysis of the \\emph{INTEGRAL} AO1 observation, together with a reanalysis of the \\emph{XMM-Newton} ToO data.  \\begin{figure*} \\centering \\includegraphics[scale=0.5]{isgri_mosa.ps} \\caption{IBIS/ISGRI mosaic of the Norma region centered on the BH 4U$1630-47$. (\\emph{left}) Energy band $20-40$ keV, (\\emph{right}) Energy band $40-60$ keV, with a total exposure of $297$ ks. Equatorial coordinates are superimposed. \\label{fig:isgri_mosaic}} \\end{figure*} ",
        "conclusions": "In the present work we confirmed and extended the results obtained by Rodriguez et al. (2003b). The nature of IGR~J$16320-4751$ is not clear yet and, even though the most probable hypothesis is a Galactic X--ray binary, the possibility of an extragalactic object cannot be completely ruled out. The source appears to be intrinsically absorbed, with an $N_{H}$ one order of magnitude greater than the Galactic absorption along the line of sight. Also the infrared observation suggests an excess due to circumstellar plasma or dust along the line of sight. The X--ray spectrum could support the hot plasma solution, but the confidence level is not enough (only $91\\%$) and should be confirmed with further observations with much more statistics. Therefore, the picture of a high--mass X--ray binary (HMXRB) with a compact object accreting through winds appears to be very likely for IGR~J$16320-4751$. The winds could also be the source of obscuration, wrapping up the compact object.  The nature of the accreting object is not clear as well: the strong differences in the photon index in the energy bands $0.5-10$~keV and $20-60$~keV could be due to changes in spectral states, thus suggesting a black hole. On the other hand, the drop of emission for energies greater than $60$~keV is common in neutron stars. Although, the \\emph{INTEGRAL} and \\emph{XMM-Newton} are not simultaneous, we tried a simultaneous fit (Fig.~\\ref{fig:allspec}). The best fit model ($\\chi^2=344.8$, $dof=344$) is an absorbed power law with exponential cutoff (\\texttt{cutoffpl} model in \\texttt{xspec}), typical of accreting neutron stars (see, e.g. White et al. 1983). The photon index is $\\Gamma=0.7_{-0.3}^{+0.2}$, the column density $N_H=(2.0\\pm 0.2)\\times 10^{23}$~cm$^{-2}$, and the cutoff energy $10\\pm 3$~keV. The scaling constant applied to the ISGRI spectrum is $5_{-2}^{+3}$. We stress that the observations were \\emph{not} simultaneous and the source is strongly variable. Therefore, this fit should be taken with extreme care.  Another possibility, given the low luminosity in both bands ($8\\times 10^{34}$~erg/s in the $0.2-12$~keV energy band, and $3\\times 10^{35}$~erg/s in the $20-60$~keV band) if the source is located in the Norma Arm ($5$~kpc), is that we are observing the emission from a jet. In this case, we could make the hypothesis that the change in the photon index could be due to a break, because of a change in the mechanism of cooling of electrons.  The key question in the study of this source is therefore the spectral variability. An approved coordinated \\emph{INTEGRAL} and \\emph{XMM-Newton} observation to be performed by next August should give us sufficient data to disentangle these hypotheses."
    },
    "0403/astro-ph0403153_arXiv.txt": {
        "abstract": "Spectroscopy of the Crab nebula along different slit directions reveals the 3 dimensional structure of the optical nebula. On the basis of the linear radial expansion result first discovered by Trimble~\\cite{c1}, we make a 3D model of the optical emission. Results from a limited number of slit directions suggest that optical lines originate from a complicated array of wisps that are located in a rather thin shell, pierced by a jet. The jet is certainly not prominent in optical emission lines, but the direction of the piercing is consistent with the direction of the X-ray and radio jet. The shell's effective radius is $\\approx 79$ seconds of arc, its thickness about a third of the radius and it is moving out with an average velocity $1160\\,\\mathrm{km\\,s^{-1}}$. ",
        "introduction": "The Crab nebula was an object of  intense interest even before its pulsar was discovered in radio by Staelin and Reifenstein~\\cite{c14} and in the optical by Cocke, Disney, and Taylor \\cite{c2}. Trimble \\cite{c1} first measured proper motions of nebular components and, combining them with spectroscopic data taken by Guido M\\\"unch~\\cite{c15}, discovered that they all point to a very small, almost point-like origin in the past.  Velusamy, Roshi, \\& Venugopal \\cite{c3} obtained a similar result from radio images of the nebula. Both observations are consistent with an outward pointing velocity field with the proper motion magnitude proportional to the distance from the origin. Trimble \\cite{c1}, Wyckoff and Murray \\cite{c4}, Caraveo and Mignani \\cite{c5} pointed out that the position of this origin is displaced with respect to that of the pulsar. Furthermore the proper motions of nebular components are about 10\\% faster than expected if the nebular debris were expanding freely. In radio images the linear expansion is not as obvious and shows some anomalies further away from the pulsar. More recently, attention has focused to phenomena connected to Crab's high-energy emission, its jet and to the very rapid expansion of inner wisps discovered optically by Hester et al. \\cite{c6} and in radio by Bietenholz et al. \\cite{c7}. It has become increasingly clear that the Crab pulsar and its nebula are connected through a multi layer structure of different predominant energy ranges. To understand the mutual interrelationship of these layers, it is useful to have a three dimensional picture of at least part of the structure.  The first attempts to understand the 3D structure of the Crab nebula go back to Chevalier and Gull \\cite{c13} who studied the nebula in the light of all 13 important visible emission lines by using narrow band filter imaging. They found a structure \"consisting of dense filaments from which thin clumpy sheets of gas fan out radially\". Clark et al. \\cite{c19} studied spectroscopically the radial velocity field across the nebula to investigate the nebula 3D properties and find the line emitting region to be a thick hollow shell composed of bright inner and faint outer components. They also see the filaments as generally circumferential, but radial \"spokes\", connecting the inner and outer shells. They conclude that the nebular synchrotron emission is confined within the inner brighter shell. A comprehensive attempt to understand the geometry, composition and mass of the Crab nebula has been done by MacAlpine et al. \\cite{c20} and \\cite{c21}. They combine interference filter imaging and long-slit spectroscopy to obtain information on composition and distribution of gas in the nebula. On the basis of a few selected spectra and using model calculations MacAlpine et al. \\cite{c20} deduce that collisional contributions to observed exceptionally strong helium lines are insignificant. Using further model assumptions they estimate the combined mass in the filaments to be in the range 6-9 $M_\\odot$. On the basis of interference filter and Fabry-P\\' erot imaging Lawrence et al. \\cite{c11} interpret the [OIII]$\\lambda 5007$ emission to have a spatial relationship to a \"high-helium torus\".  The possibility to make a 3D representation of the line emitting regions in the Crab nebula is offered by the proper motion result of Trimble and others. On the basis of this result one is led to speculate that all the three components of the velocity vector field describing the motion of nebular components point in the outward radial direction since the motion is ballistic. According to this assumption the speed of this motion is proportional to the distance from the origin, i.e. we assume: \\begin{equation} \\label{eq:trimble} \\mathbf{v} = \\mathbf{r}/T \\quad , \\end{equation} where $\\mathbf{v}$ is the velocity vector at the position $\\mathbf{r}$ measured from the center of the nebula, and $T$ is time elapsed. If the ballistic assumption was fully satisfied, $T$ would be the time since the explosion in 1054, but, as Trimble's proper motion results show the possibility of some acceleration of the debris in the past, we consider $T$ as a model parameter that may differ somewhat from the time since the explosion. Observing spectroscopically a small patch of the nebula at the angular position $(\\xi,\\,\\eta)$ from the pulsar, one expects to find multiple Doppler shifted emission lines, each Doppler shift corresponding to a different radial velocity component ($v_z$) of the cloud. According to (\\ref{eq:trimble}), such a component corresponds to a line of sight coordinate\\- $z = ({v_z} - {v_{zc}}) T$, while the other two components of the vector $\\mathbf{r}$ are $x = (\\xi - \\xi_c) D$ and $y = (\\eta - \\eta_c) D$, where $ v_{zc}$ and $(\\eta_c, \\xi_c)$  are the velocity and (angular) coordinates of the center of the cloud, while $D$ is the distance to the nebula. Locating line emitting regions in this way, one can construct a tentative 3D image of (line emitting) nebular components. The only parameter critical for the shape of such a reconstructed image is the ratio $T/D$. We choose to consider it as a free parameter and construct 3D models of the nebula for different values of it. We find it remarkable that it is possible to find a value for this parameter such that the 3D model of the nebula in  the radiation from the strongest optical emission lines ([NII], [SII], and  H$_\\alpha $) appears as embedded inside a relatively thin, pierced spherical shell with a well defined inner radius. We consider this particular ratio $T/D$ of some interest, yet, we would like to emphasize that its implications for the distance to the nebula depend strongly on the ballistic assumption. ",
        "conclusions": "We obtained multiple long slit spectra of the Crab nebula in the light of the optical emission lines of $\\mathrm{H}_\\alpha$, [NII]$\\lambda$6548,6583, and [SII]$\\lambda$6716,6730. The simple assumption of ballistic expansion, expressed by (\\ref{eq:trimble}) and suggested by the work of Trimble~\\cite{c1} allows a 3D representation of the nebula, although through an arbitrary scale in the line of sight. Such a 3D representation suggests that line emitting regions are organized in a relatively thin shell, regardless of the line of sight scale. Therefore, we make the assumption that the shell is likely to be quasi spherical and adjust the unknown scale parameter accordingly. The resulting 3D image indeed appears to fill a spherical shell with an inner radius of $\\approx 55\\,\\mathrm{arc\\,seconds}$, as already shown by results of Lawrence et al. \\cite{c11}. The shell is pierced by the jet in the SE direction and at about $75^\\circ$ with respect to the line of sight. The center of this shell is well aligned with the two-dimensional projection of the center of the optical and X-ray continuum emissions. It is clearly displaced from the pulsar  but appears to be connected to it through the north-western extension of the jet. The 3D distribution of the optical continuum and the X-ray emission can not be determined, but the observed 2D projections of these are very similar and consistent with the idea that the optical and X-ray continuum come from the same, rather flat, oblate structure centered in the spherical shell. The optical continuum emission drops regularly with distance from the center of the nebula (Fig.~\\ref{Continuum0036}), becoming weak when reaching the spherical shell.  Based on the absence of the [NII]5755 line, we estimate that the effective gas temperature in the shell is below 7500K. The filling of the $55\\,\\mathrm{arc\\,seconds}$ shell seems patchy. We note that some bright spots are only arc seconds wide (cf. Fig.~\\ref{Oba036x}) and possibly less than $16\\,\\mathrm{km\\,s^{-1}}$ ``deep'' (Fig.~\\ref{OzkiInSiroki}), so that we expect  that a more detailed high resolution 3D view would show that the shell is woven with delicate threads. Note that the small size of line emitting regions indicates high emission coefficients in the bright regions; in particular, we estimate that the region shown in Figs.~\\ref{Oba036x} and \\ref{OzkiInSiroki} is only about 0.02 ly thick which leads to the H$_\\alpha$ emission coefficient of the order $j_{\\mathrm{H}_\\alpha}\\rho ~\\approx ~10^{-21}\\mathrm{\\,erg\\,cm^{-3}\\,sec^{-1}}$, which is an order of magnitude higher than the average emission coefficient at these wavelengths of the whole nebula including the synchrotron cloud. Hints of threads are suggested in the superimposition of the radio image by Velusamy et al. \\cite{c3} on the optical image and are clearly seen in the newer radio image by Bietenholz et al.~\\cite{c7} shown in Fig.~\\ref{FourImages}. Note that optical filaments closely map the smoother radio loops. The jet, which is so prominent in X-rays and in radio, is marked also by optical continuum emission and seems to be threaded also by weak line emission (see Fig.~\\ref{Grobi036}, in particular at the $144^\\circ$ orientation and Fig.~\\ref{144Sat}). Line emission from regions where the jet hits the $55\\,\\mathrm{arc\\,seconds}$ shell is particularly strong - the upper left brightest spot in the reduced spectrum $\\psi_s=144^\\circ$ and also $\\psi_s=162^\\circ$ of Fig.~\\ref{Grobi036}.  In concluding, we would like to call attention to two observations: The roundness and average regularity of the spherical shell as expressed by the sharpness and simplicity of the radial distribution of line emitting regions in Fig.~\\ref{RadDist} is quite remarkable. The delicate radiating tubes or blobs filling this relatively thin shell closely follow radio loops on large scale,  even if, as emphasized by Sankrit et al.~\\cite{c34} and Hester et al.~\\cite{c33}, on short scale they are dominated by the magnetic Rayleigh-Taylor instabilities. The 3D image, even if it still lacks in detail, suggests that the filaments are wound on the shell along long smooth threads, possibly following the magnetic field lines."
    },
    "0403/astro-ph0403365_arXiv.txt": {
        "abstract": "The origin of the luminosity dependence of the equivalent width (EW) of broad emission lines in AGN (the Baldwin effect) is not firmly established yet. We explore this question for the broad \\C~$\\lambda 1549$ line using the Boroson \\& Green sample of the 87 $z\\le 0.5$ Bright Quasar Survey (BQS) quasars. Useful UV spectra of the \\C\\ region are available for 81 of the objects, which are used to explore the dependence of the \\C~EW on various emission properties. We confirm earlier results on the strong correlations of the \\C~EW with some of the emission parameters which define the Boroson \\& Green Eigenvector 1, and with the optical to X-ray slope $\\alpha_{ox}$. In addition, we find a strong correlation of the \\C~EW with the relative accretion rate, \\Ledd. Since \\Ledd\\ drives some of the Eigenvector 1 correlations, it may be the primary physical parameter which drives the Baldwin effect for \\C. ",
        "introduction": "The inverse correlation of the broad emission line equivalent width (EW) with luminosity in AGN, discovered by Baldwin (1977) for the \\C~$\\lambda 1549$ line, was intensively explored over the past 20 years (see a comprehensive review by Osmer \\& Shields 1999; and more recent studies by Wilkes et al. 1999; Green et al. 2001; Croom et al. 2002; Dietrich et al. 2002; Kuraszkiewicz et al. 2002; and Shang et al. 2003). The physical origin for this effect is not clearly established yet, but a plausible explanation is softening of the ionizing continuum shape, and increasing gas metalicity with increasing luminosity (Korista, Baldwin \\& Ferland 1998). The softening of the ionizing continuum shape with luminosity is supported by: 1. some observed correlations of the emission line strength with the ionizing spectral shape (Green 1998; Wang, Lu \\& Zhou 1998), 2. predictions of simple accretion disk models for the ionizing spectral shape dependence on luminosity (Netzer, Laor \\& Gondhalekar 1992), and 3. the dependence of the slope of the Baldwin effect on the ionization potential of the emitting ions (Espey \\& Andreadis 1999; Green et al. 2001; Dietrich et al. 2002; Kuraszkiewicz et al. 2002). Non isotropic continuum emission (Netzer 1987), and the intrinsic Baldwin effect (Pogge \\& Peterson 1992; also review in Osmer \\& Shields 1999, \\S 5) can produce some of the observed scatter in the Baldwin effect for the various lines.  Significant correlations also exist among various optical emission line properties in AGN, including the [O III] and Fe~II strength, and the \\Hb~FWHM, which form part of Eigenvector 1 (EV1) in the comprehensive study of Boroson \\& Green (1992, hereafter BG92). Boroson, Persson \\& Oke (1985) speculated that the underlying physical parameter which drives these correlation is the relative accretion rate, \\Ledd, a speculation which received strong support through recent observations (see review in Laor 2000). Furthermore, Wills et al. (1999) found that various UV emission properties, including the \\C~EW, are correlated with the optical EV1 parameters, which suggests that the \\C~EW may also be largely driven by \\Ledd. The Baldwin effect may then just be a secondary correlation induced by the tendency of more luminous AGN to have a higher \\Ledd. The possible role of \\Ledd\\ in the Baldwin effect was already mentioned by Brotherton \\& Francis (1999). Later, Wilkes et al. (1999) noted that narrow line AGN are outliers to the Baldwin effect, but the meaning of this result was not interpreted at that time (see \\S 4 here).  Most recently Shang et al. (2003) suggested an indirect indication, based on spectral principle component analysis, that \\Ledd\\ may contribute to the scatter in the Baldwin effect.  Recent studies established that reasonably accurate estimates of the black hole mass ({\\Mbh}) can be obtained in AGN based on the continuum luminosity and {\\Hb} FWHM (\\Mbh$\\propto L^{1/2}$(\\Hb~FWHM)$^2$, and thus \\Ledd$\\propto L^{1/2}$(\\Hb~FWHM)$^{-2}$, e.g. Laor 1998). This opens up the possibility to test directly whether the Baldwin effect is driven by {\\Ledd}, which is the main point of this paper. The data set and measurement procedure are described in \\S 2, the correlation analysis is presented in \\S 3, and the main results are summarized in \\S 4. ",
        "conclusions": "The purpose of this paper is to obtain some indications for the physical parameters which drive the Baldwin effect in \\C\\ through correlation analysis. We use archival {\\it HST} and {\\it IUE} \\C\\ spectra of sufficient quality available for 81 of the 87 BQS quasars, together with the optical emission line parameters from BG92, and $\\alpha_{\\rm ox}$ from Brandt et al. (2000).  We find that a major source of scatter in the Baldwin effect is the \\Hb~FWHM. Its inclusion as a second parameter to $L$, in the form $L^{1/2}\\times$(\\Hb~FWHM)$^{-2}$, which is proportional to \\Ledd, leads to a significant increase in the correlation strength (from $r_s=-0.154$ to $r_s=-0.581$). This indicates that the Baldwin effect is a secondary relation, which is induced by the stronger relation of the \\C~EW and \\Ledd, and the tendency of high $L$ AGN to have a higher \\Ledd. This also explains why NLS1s, which are high \\Ledd\\ AGN at a low $L$, have much lower \\C~EW than expected for their $L$ (as found by Wilkes et al. 1999).  We do not know the physical mechanism responsible for the reduction in the \\C~EW with increasing \\Ledd. However, we find that the scatter in the Baldwin effect can be further reduced by including either the \\Ox~EW, the relative \\Fe\\ strength, or $\\alpha_{\\rm ox}$, as a third parameter. The \\Fe\\ strength may be an indicator of the BLR optical depth, and $\\alpha_{\\rm ox}$ may be an indicator of the ionizing continuum shape, which are known theoretically to affect the \\C~EW. The physical mechanism implied by the correlation with the \\Ox~EW remains to be understood. At least some of the remaining scatter can be produced by a non-isotropic, or a time variable continuum source.  The results presented here can be tested with larger samples of AGN. In particular, it will be interesting to explore the extension to high $z$ AGN, where the \\Hb\\ region is observable in the IR (e.g. McIntosh et al. 1999; Yuan \\& Wills, 2003; Shemmer et al. 2004). In addition, different lines show different slopes for the Baldwin effect, and it will be interesting to explore what are the primary physical parameters which drive other emission lines."
    },
    "0403/astro-ph0403686_arXiv.txt": {
        "abstract": "High precision astrometric data have enabled us to determine the trajectories through the Galactic potential for a growing number of pulsars. This has resulted in the discovery of a pulsar pair (B2020+28 and B2021+51) that has a common origin in the Cygnus Superbubble or in one of its related OB associations at an epoch which is comparable with the spin-down ages of the pulsars. Analysis of the Galactic orbits indicates that the progenitors of the pulsars had similar masses and were in a binary system, which was disrupted after the second supernova explosion. The implied pulsar birth velocities are consistent with the high velocities of neutron stars in general. The initial spin period of the pulsar that was formed in the second supernova explosion was $\\sim 200$~ms. A further increase in astrometric accuracy will allow us to more tightly constrain the birth velocities and the kick velocities that were imparted by the two respective supernova explosions. Two additional pulsars in the sample of 24 with parallax measurements may also have originated in the Cygnus Superbubble. ",
        "introduction": "The velocity distribution of pulsars has long ago been shown to extend to much larger velocities then those of their progenitors \\citep{GO70}. Several possible scenarios have been suggested as causes for high pulsar velocities. These include asymmetry in the birth supernovae (SN) of pulsars \\citep{S70, DC87} and the disruption of binaries by mass-loss from the SN explosion \\citep{B61,GGO70}. Considerable evidence exists that both effects take place.  Exactly how an asymmetry in the core collapse of a SN explosion is converted to a kick imparted on the newly formed neutron star (NS) is still unclear; possibilities include hydrodynamic or convective instabilities \\citep{BH96, JM96, LG00} and mechanisms such as asymmetric neutrino emission in the presence of strong magnetic fields \\citep{AL99}.  Binarity is particularly common among the early type O and B stars that are expected to produce Type II SN and subsequently NS; approximately 70\\% of these stars exist in binary systems \\citep{B67}. In about 50\\% of the binaries, both stars are massive enough to become SN; most of these binaries are disrupted when the first SN occurs, while only a small percentage evolves into NS-NS binaries \\citep{BB98}. Until now 6 of such NS-NS binaries have been detected \\citep{burg03}.  While most massive star binaries are disrupted at either the first or the second SN, many individual pulsars originate from such binaries. Identification of these pulsars will allow us to put constraints on kick velocities and the amount of asymmetry in the SN explosions.  Previously, several pairs pulsars where hypothesized to come from the same disrupted binary based on their positional information alone \\citep{GGO70, W79}.  Recently, high precision VLBI astrometry has produced accurate measurements of parallaxes and proper motions of about two dozen nearby pulsars, which allows much more accurate tests for pulsar pair identifications. Using the astrometric data, we have analyzed possible trajectories for the pulsars that take into account the Galactic gravitational potential and plausible values of the unknown radial velocities.  We find that two of the pulsars (B2020+28 and B2021+51) have Galactic orbits that cross at an epoch consistent with their spin-down ages, which makes them prime candidates for pulsars that have been ejected from the same binary system. In this paper we first discuss the identification of the pulsar pairs and describe the Galactic potential used to calculate the pulsar orbits (\\S~\\ref{sec2}).  The results of the orbit calculations are presented in \\S~\\ref{sec3}. The birth location of the pulsars and several possible birth scenarios are discussed in \\S~\\ref{sec4}, followed by our conclusions in \\S~\\ref{concl}. ",
        "conclusions": "\\label{concl}  After examining the available sample of pulsars that have precision astrometric data, we have found that the pulsars B2020+28 and B2021+51 likely share their origin in a disrupted binary even though they are presently separated by $\\sim 23^\\circ$ on the plane of the sky. The Galactic orbits of the pulsars, traced back through a 3 component St\\\"ackel potential that satisfies the current Galactic parameters, are found to cross approximately 1.9 Myr ago at 1.9 kpc distance from the Sun in the region known as the Cygnus superbubble. Simulations of pulsar samples, and an analysis of the likelihood functions, show that it is unlikely that our pulsar pair identification is the result from a chance alignment.  The most likely scenario for the creation of the runaway pulsar pair is the disruption of a binary system in the second SN explosion. As no signs of accretion are present, the progenitor stars must have had similar masses and have been in a wide binary orbit. The second SN must have followed the first by only a short time ($<0.5$~Myr).  The velocities and direction angle of the pulsars at birth can be explained if the second (asymmetric) SN imparted a kick velocity of $\\approx 200$~\\kms. The velocity analysis indicates that B2020+28 was created in the second SN.  The system was disrupted $\\approx~1.9$~Myr ago, which must correspond to the true age of at least one of the pulsars. The spin-down ages of both pulsars are very similar: B2020+28 has a spin-down age of $2.88$~Myr, while B2021+51 has a spin-down age of $2.75$~Myr. As the relationship of spin-down age to true age is largely uncertain, the similarity between our derived age and the spin-down ages is remarkable.  An age of 1.9 Myr implies that the initial period of B2020+28 was 201~ms, assuming a constant braking index of $n=3$. If the SN that created B2021+51 was the source of the binary disruption, the initial pulsar period will have been 295~ms, also assuming a constant $n=3$. In both cases, a constant braking index of $n>4$ can be ruled out.  The most likely three dimensional space velocities of the pulsars at birth are $\\approx 150$~\\kms\\ for B2021+51 and $\\approx 500$~\\kms\\ for B2020+28. Recent analyses of pulsar velocities have found that a two component birth velocity distribution seems to best fit the current data \\citep{CC98, ACC02}. Using the sample of pulsars with accurate astrometric data, Chatterjee et al. (2004) find that the pulsar velocities are best described by a 2-component birth velocity distribution with characteristic velocities of 86~\\kms\\ (70\\%) and 292~\\kms\\ (30\\%). The probability of finding pulsars with a three dimensional space velocity $\\geq 500$~\\kms\\ ranges from $\\sim 10\\%$ using the Chatterjee et al. (2004) distribution, which has probable selection effects against the highest velocity pulsars, to $\\sim 45\\%$ for the \\citet{ACC02} distribution, which was based on a much larger sample if pulsars, but used a different methodology that was based largely on DM-distances using the \\citet{TC93} model. The birth velocities determined for B2020+28 and B2021+51 then indicate that B2021+51, likely created in the first SN explosion, is part of the lower velocity population, while B2020+28, created in the SN explosion that disrupted the binary system, is part of the high velocity component.  The current astrometric data are completely consistent with the common binary origin of the two pulsars and an improvement of astrometric precision is unlikely to invalidate our conclusions. However, as the accuracies of the proper motion and parallax measurements are increased, we will be able to more precisely determine the current radial velocities of the pulsars. This will lead to a significant improvement on the determination of the birth velocities and as a result, we will be able to obtain better constraints for the kick velocities of the SN explosions as well as other characteristics of the progenitor binary system.  As the sample of pulsars with accurate proper motions and parallaxes increases, the identification of additional related pulsar pairs will allow us to put constraints on kick velocities and the amount of asymmetry in the SN explosions. Unfortunately, due to pulsar beaming, we can only detect approximately one out of five pulsars. This ratio is somewhat better for young pulsars, which, as seen in \\S~\\ref{chance}, are also the pulsars for which a positive pair identification is the most likely. For a compete sample of nearby, young pulsars with precision astrometric data, we expect to be able to identify related pulsars for at least $10\\%$ of them."
    },
    "0403/astro-ph0403403_arXiv.txt": {
        "abstract": "We present the results of an investigation on the main structural properties derived from {\\em VRI} and \\Ha\\ surface photometry of galaxies hosting nuclear emission-line regions (including Seyfert~1, Seyfert~2, LINER and starburst galaxies) as compared with normal galaxies. Our original sample comprises 22 active galaxies, 4 starbursts and 1 normal galaxy and has been extended with several samples obtained from the literature. Bulge and disc parameters, along with $B/D$ relation, have been derived applying an iterative procedure. The resulting parameters have been combined with additional data in order to reach a statistically significant sample. We find some differences in the bulge distribution across the different nuclear types that could imply familes of bulges with different physical properties. Bulge and disc characteristic colours have been defined and derived for our sample and compared with a control sample of early type objects. The results suggest that bulge and disc stellar populations are comparable in normal and active galaxies. ",
        "introduction": "Global AGN host galaxy properties, including mass, luminosity concentration, morphological type, bulge-to-disc relationship, metal abundance and total luminosity, may influence nuclear activity and perhaps in turn be affected by it. It has been argued that there can exist a connection between nuclear activity and intense star formation episodes. While there are some observational indications that star formation is enhanced in circumnuclear regions of Seyfert galaxies (see for instance Hunt \\& Giovanardi 1992), other observations point in the opposite sense (e.g. Carone 1992). An outstanding, still controversial question is related to the fuel supply involving the transport of large quantities of mass from reservoirs located in the galactic disc (see for instance Shlosman et al. 1989); in this process the matter has to lose almost all of its angular momentum. This can be accomplished by gravitational torques produced by non-axisymmetric perturbations in the gravitational potential. These are associated either with interactions with companion galaxies or with structural features like bars, rings, oval distortions etc.  Current research on structural properties of galaxies hosting nuclear activity is now largely concentrated in  the near infrared (NIR) bands. There are several advantages in the use of the NIR for studying the host galaxy properties of AGNs: on the one hand, a normal   stellar population is dominated by old red giants and therefore peaks around 1$\\mu$m. On the other, dust extinction is substantially reduced at these wavelengths. Therefore, the NIR is ideal to look for bars and morphological distortions that can be hidden by dust in optical wavelengths (e.g. Mulchaey \\& Regan 1997). Moreover, the contrast between stellar and non-thermal contributions is maximised, since the latter has a local minimum around 1$\\mu$m.  Hunt et al. \\shortcite{Hunt1997} performed NIR $J$, $H$, $K$ broadband images and colours in a sample of 26 galaxies with $ z \\geq 0.015$ and found that the average colours, both of the outer and the inner discs, are independent of the activity class and equal to those of normal spiral galaxies. Outer disc colour gradients are also consistent with those of normal spirals. On the other hand, on a scale of a few kiloparsecs from the centre, starburst galaxies show ongoing star formation; in Seyfert~2 galaxies, they found evidence of ``fossil'' star formation activity a few hundred Myr old; finally, no evidence of star formation, either ongoing or fossil, exceeding that of a normal spiral, is found in Seyfert~1 galaxies. They infer an evolutionary starburst-Seyfert connection, although not so global and not so tight in time as previous studies implied. NIR structural parameter analysis from Hunt et al. \\shortcite{Hunt1999} shows that bulges of Seyfert and starburst galaxies follow very closely the trend defined by normal spiral bulges and elliptical galaxies. On the other hand, for a given disc effective radius, the mean surface brightness of the Seyfert discs is higher (by $\\sim$ 0.9 K mag arcsec$^{-2}$) than those of early-type spirals.  M\\' arquez et al \\shortcite{Marquez2000} performed NIR $J$ and $K'$ broadband imaging and colours on nearby, isolated spiral galaxies either with an active nucleus (18 objects, either Seyfert~1 or Seyfert~2 according to V\\' eron-Cetty \\& V\\' eron 1993) or without it (11 galaxies). They state that global properties are similar in both active and non active galaxies, since both samples define the same bulge Kormendy relation between $\\mu_e$ and $r_e$\\footnote{Nevertheless, there is a difference in the slope of the Kormendy relation of active and non active galaxies at $\\sim 2 \\sigma$ level that is not considered as relevant.} and disc components also share the same properties (contradicting the result from Hunt 1999); they also found that bulge and disc scale lengths are correlated, with $r_{bulge} \\approx 0.2 r_{disc}$. On the other hand, NIR central colours of active galaxies are found to be redder than those of non active spirals, most probably due to the AGN light re-emitted by the hot dust and/or to the presence of active star formation in circumnuclear regions.  In spite of the advantages of the NIR region for such studies, optical wavelengths can still provide useful data; the use of the optical region has its own advantages: on the one hand, observers can benefit from the large size of current optical detectors, providing a convenient sampling of the galaxy disc; on the other, optical studies can rely on the large quantity of published data. Moreover, optical colour maps can be used as excellent tracers of the line emission regions present in AGNs and starburst galaxies (see, for instance Terlevich et al. 1991)  Early work from Yee \\shortcite{ApJ272:Yee} using a SIT-vidicon camera has been followed by a large number of studies using CCD detectors. Mediavilla et al. (1989) and  Pastoriza et al. \\shortcite{Pastoriza} performed {\\em VI} broadband surface photometry on a sample of 13 Seyfert galaxies; they consider two classes: ``bulge galaxies'' and ``disc galaxies'' and found that the \\mbox{\\em V-I} colour has a constant, equivalent value in both galaxy groups in the 3~Kpc aperture, indicating similar stellar populations, but outside the 5~Kpc aperture colours become bluer for disc galaxies and redder for bulge galaxies.  Mediavilla et al. \\shortcite{ApSS157:Mediavilla} use the same sample to obtain bulge and disc parameters from the surface brightness radial profiles, finding that structural parameters are comparable to those of normal spirals and ellipticals.  MacKenty \\shortcite{ApJS72:MacKenty} performed Kron-Cousins {\\em BVR} CCD surface photometry of 51 Seyfert galaxies from the NGC and Markarian catalogs finding that in almost all cases, a mechanism to transport material to the nucleus is present (either bars or interactions); about one-fifth of the sample is composed by amorphous galaxies; some of them can be a remnant of past interactions. Colours and disc morphological parametrs are in general similar to those found in normal spiral galaxies.  Xanthopoulos \\shortcite{MNRAS280:Xanthopoulos} performed {\\em VRI} CCD surface photometry of 27 Seyfert galaxies belonging to type 1, 2 and intermediate categories, finding that: {\\em (a)} generally, Seyfert galaxies are perturbed objects (either by bars or by interactions); {\\em (b)} nearly half of the sample galaxies are barred; moreover, almost all of them present evidence of the existence of a mechanism to transfer material to the nuclear region;  {\\em (c)} all non-barred galaxies are either S0 or Sa; {\\em (d)} Seyfert~1 and Seyfert~2 galaxies follow a very similar Hubble type distribution; Seyfert~2 galaxies tend to have redder nuclei; {\\em (e)} Seyfert~1 and Seyfert~2 galaxies present the same disc colours. {\\em (f)} colour gradients remain constant along  both the bulge and the disc of the galaxy for both Seyfert types; {\\em (g)} both types of Seyfert galaxy show the same range and dispersion of bulge and disc morphological parameters, and also similar bulge-to-disc luminosity ratios; the wide range of scale lengths observed in Seyfert galaxies may support their classification as early-type spirals. Moreover, the bulge and disc parameters of both Seyfert types are the same as those of normal early-type spirals.  Virani et al. \\shortcite{Virani} performed {\\em R} surface photometry on a sample of 32 Seyfert and 45 non active galaxies\\footnote{From the control sample, at least two galaxies have been found to be active and at least five objects are nuclear starbursts according to Ho et al. \\shortcite{Ho1997}.} from the CfA survey and derived bulge and disc morphological parameters from the azimuthally averaged surface brightness profiles. No statistically significant differences between the Seyfert and control samples are found in either morphological parameters or in light asymmetries (bars, rings, isophotal twisting, etc.), or even in the companion galaxies properties (magnitude, separation from the host, position angle relative to the host, magnitude difference  between the companion and the host, strength of the tidal parameter). They conclude that the nearby environment of Seyfert galaxies is not significantly different from nonactive galaxies with the same morphological distribution.  Boris et al. \\shortcite{Boris} performed {\\em BVI} surface photometry on a sample of 10 Seyfert~1 galaxies along with narrow-band \\Ha\\ images of a subsample of 6 objects. They found that morphologies are confined to early-type galaxies; half of them can be considered as compact; bars are found in only two cases; average \\mbox{\\em B-V} colours are found bluer than those expected for this morphological type (perhaps due to the contribution of the active nucleus and/or the disc to the total luminosity). Evidence of tidal interactions is found in six objects of the sample. They obtain morphological parameters for the nucleus (gaussian), bulge and disc components finding that in most cases (six out of eight), the disc is best described by a truncated  exponential profile. The cutoff is generally associated with the existence of reddened regions (probably dust in the inner few Kpc of the galaxy). Only one galaxy shows disc emission in \\Ha .  This paper presents a new optical study on the structural parameters of galaxies hosting nuclear emission-line regions. In order to overcome the usual weakness of this kind of studies (relatively small and frequently biased samples, uncertainties in the nuclear classification, specially in the control samples), whenever possible, we have tried to supplement our original sample with data from the literature, in particular morphological parameters, and to use an homogeneous, reliable nuclear classification from optical spectroscopy. When data from the literature is used, it is clearly indicated throughout.  The paper is structured as follows: first, a description of the original sample and the adopted nuclear classification is outlined; second, a description of the bulge and disc decomposition procedure is given, followed by the results obtained; the discussion section then proceeds, including a thorough analysis of the structural parameters of galaxies hosting nuclear emission-line regions. Our analysis is closed with a discussion on the radial colour distribution of our sample objects. Finally, our conclusions are outlined in section \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  We have investigated the main structural properties derived from {\\em VRI} and \\Ha\\ surface photometry of galaxies hosting nuclear emission-line regions (including Seyfert~1, Seyfert~2, LINER and starburst galaxies) as compared with normal galaxies. Our original sample is thoroughly described in Paper I and comprises 21 active galaxies, 3 starbursts and 1 normal galaxy. For our investigation on bulge and disc parameter distributions, we have extended the sample up to 261 objects hosting all levels of nuclear activity (ranging from normal to Seyfert~1 galaxies). Our study has found several differences between the structural parameters of active and non-active galaxies. From the statistical analysis performed in section \\ref{disc_bulge_parameters}, we cannot claim that the individual morphological parameters of active and non-active galaxies follow the same statistical distribution. On the other hand, we have studied the distribution of morphological parameters in the ($\\mu_0$, $\\log r_0$) and ( $\\mu_e$, $\\log r_e$) planes. Disc parameters are found strongly clustered (perhaps a selection effect) but the linear correlation is weak. On the other hand, there is a strong linear correlation of bulge parameters. We have studied the distributions accross the different nuclear types, finding that starbursts constitute a lower luminosity bulge class. LINERs and normal galaxies follow very similar distributions. Seyfert~1, Seyfert~2 and LINERs follow different distributions at statistically significant level. These differences could perhaps imply families of bulges with different physical properties. Type II discs are analysed in section \\ref{TypeII_discs}. We don't find a correlation between type II discs and nuclear type but with Hubble type. $B/D$ relation is studied in section \\ref{b_d_relation}. We find that the distribution of $(B/D)_{median}$ seems to be generally independent of the nuclear type, except for the earliest ($T < 0$) Hubble types. Finally, bulge and disc characteristic colours derived from radial profiles have been investigated in section \\ref{radial_colour_distribution}. We find that the range of characteristic bulge and disc colours is comparable in all nuclear types, though mean bulge and disc colours are bluer in Seyfert galaxies than in LINERs. Nevertheless, this result is uncertain since bulge colours can be affected by line emission contamination arising from the NLR and/or circumnuclear star-forming regions. When comparing our characteristic bulge and disc colours with those of the control sample of early-type galaxies,we find a very similar colour range that may suggest that the bulge and disc stellar populations are comparable in normal and active galaxies. On the other hand, we find that bulges and disc characteristic colours derived from profiles are very similar. We don't detect a correlation between bulge colours and bulge or total galaxy luminosities."
    },
    "0403/astro-ph0403129_arXiv.txt": {
        "abstract": "We discuss HI and optical evidence for ongoing ICM-ISM interactions in 6 HI-deficient Virgo cluster spiral galaxies. One of the clearest cases is the highly inclined Virgo galaxy NGC~4522, which has a normal stellar disk but a truncated gas disk, and lots of extraplanar gas right next to the gas truncation radius in the disk. Unusually strong HI, H$\\alpha$ and radio continuum emission are all detected from the extraplanar gas. The radio continuum polarized flux and spectral index peak on the side opposite the extraplanar gas, suggesting ongoing pressure by the ICM. Four other HI-deficient edge-on Virgo spirals show evidence of extraplanar ISM gas or exhibit asymmetries in their disk HI distributions, but contain much less extraplanar HI than NGC~4522. Comparison with recent simulations suggests this difference may be evolutionary, with large surface densities of extraplanar gas observed only in early phases of an ICM-ISM interaction. In NGC~4569, the H$\\alpha$ image shows 2 effects of ICM pressure on the galaxy ISM. An anomalous arm of HII regions, possibly extraplanar, emerges from the edge of a truncated H$\\alpha$ disk. This resembles the arms seen in simulations which are formed by the combined effects of wind pressure plus rotation. An extended nebulosity near the minor axis, also in the NW, is interpreted as a starburst outflow bubble disturbed by ICM wind pressure. ",
        "introduction": "Interactions between the interstellar medium (ISM) of a galaxy and the gas in the intracluster medium (ICM) are believed to be one of the main processes which drive galaxy evolution in clusters (Gunn \\& Gott 1972; Poggianti et al 1999; Koopmann \\& Kenney 2003; van Gorkom 2003). Gas stripping is likely responsible for the transformation of many cluster spirals into  S0's and Sa's, and thus may help cause both the morphology-density relation (Dressler et al 1997) and the Butcher-Oemler effect (Butcher \\& Oemler 1978). ISM stripping is also significant for the evolution of the ICM, as gas removed from the galaxies enriches the ICM with heavy elements.  There is widespread evidence for ICM-ISM stripping in cluster spiral galaxies, from single dish HI studies showing that many cluster galaxies are HI-deficient (e.g., Solanes et al 2002) , and interferometer HI studies showing many cluster spirals with truncated gas disks (Cayatte et al 1990; Bravo-Alfaro et al 2001). Since a galaxy will be deficient in gas and have a truncated gas disk long after an interaction, it is not clear whether stripping is active or occurred in the past for most cluster galaxies. Here we discuss HI and optical evidence for ongoing ICM-ISM interactions in 6 HI-deficient cluster spiral galaxies in the nearby Virgo cluster (16 Mpc). With detailed observations of galaxies experiencing ICM-ISM stripping we can begin to learn what actually happens in an ICM-ISM interaction, and the consequences for galaxy evolution. ",
        "conclusions": ""
    },
    "0403/astro-ph0403635_arXiv.txt": {
        "abstract": "We critically examine the theory of disk locking, which assumes that the angular momentum deposited on an accreting protostar is exactly removed by torques carried along magnetic field lines connecting the star to the disk.  In this letter, we consider that the differential rotation between the star and disk naturally leads to an opening (i.e., disconnecting) of the magnetic field between the two.  We find that this significantly reduces the spin-down torque on the star by the disk.  Thus, disk-locking cannot account for the slow rotation ($\\sim$10\\% of breakup speed) observed in several systems and for which the model was originally developed. ",
        "introduction": "}      We now know that most protostars undergo a phase in which they accrete mass from a disk and that a stellar magnetic field dominates this process near the star.  However, the transport of angular momentum, while fairly well-understood within the disk, remains a significant challenge for models of stellar spin \\citep[for a review, see][]{bodenheimer95}.  Classical T Tauri stars (CTTS's) comprise a class of accreting protostars for which there is an abundance of observational data.  For typical parameters in these systems, the angular momentum deposited by accretion can spin the star up to breakup speed in $\\sim 10^5$ years, assuming the star hadn't already formed at breakup speed.  Since the accretion lifetime is often greater than $10^6$ yr, the stars must somehow rid themselves of this excess angular momentum.  Furthermore, a significant number of CTTS's (the so called ``slow rotators'') have spin rates as slow as $\\sim$10\\% of breakup speed \\citep[e.g.,][]{bouvierea93}, requiring significant torques to maintain.  \\citet[][hereafter K91]{konigl91} proposed the currently accepted view that the slow rotators could be explained with a ``disk-locking'' (DL) mechanism \\citep[adapted from][]{ghoshlamb79}, for which the star needs a kilogauss-strength dipole magnetic field.  Indeed, CTTS's are now known to possess kG-strength fields \\citep*[e.g.,][]{johnskrull3ea99}, and magnetic star-disk interaction models have also explained several observed spectral features of CTTS's (e.g., K91).  However, observations by \\citet{stassunea99} and \\citet{johnskrullea99} and theoretical considerations of \\citet{safier98} have called the standard DL scenario into question. According to DL theory, magnetic field lines causally connect the star to the disk (acting as ``lever arms'') and carry torques that oppose and balance the angular momentum deposited by accretion.  Although not usually considered in DL theory, \\citet*[][hereafter UKL]{uzdensky3ea02} demonstrated that differential rotation between the star and disk leads to an opening of the field, drastically reducing the magnetic connection between the two.  In this letter, we consider the consequences of this reduced connection on DL.  We find that the magnetic torque is significantly reduced, and show that the DL model therefore cannot account for the angular momentum loss of the slow rotators. ",
        "conclusions": "\\label{sec_summary}     We have shown that a large portion of the magnetic field connecting the star to the disk will open up due to the differential rotation between the two, resulting in a spin-down torque on the star by the disk that is much less than if the field is assumed to remain closed. The predicted disk-locked spin rate is therefore much faster, so that the DL scenario cannot explain the angular momentum loss of the slow rotators.   In addition, there are at least three, completely independent issues raised by other authors: 1) \\citet{stassunea99} found no correlation between accretion parameters and spin rates of TTS in Orion.  2) CTTS's apparently do not have strong dipole fields \\citep[e.g.,][]{safier98, johnskrullea99}, which are required for DL at slow spin rates.  3) Stellar winds are expected to open field lines that would otherwise connect to the disk \\citep{safier98}.  A disk wind could have a similar effect.       We conclude that, in order for accreting protostars to spin as slowly as 10\\% of breakup speed, there must be spin-down torques acting on the star other than those carried by magnetic field lines connecting the star to the disk.  The presence of open stellar field lines implies that excess angular momentum may be carried by a stellar wind \\citep[e.g.,][]{toutpringle92}, but this remains an open question.  We are investigating the role of stellar winds, since then $\\spineq$ will depend mostly upon the stellar wind parameters, thereby explaining the lack of correlation between spin periods and accretion parameters."
    },
    "0403/astro-ph0403545_arXiv.txt": {
        "abstract": "{ In order to investigate the properties of large-scale structures of galaxies in the universe, we  present an analysis of their spatial distribution at z$<$0.033. We used the LEDA extragalactic database containing over 1 million of galaxies covering the all-sky and the SDSS data included in the public release DR1, yielding to a sample of around 134,000 galaxies having a measured redshift in two survey areas representing 690 sq. degrees. The results of the study are 2D, 3D maps and magnitude number counts of galaxies, drawn from B-band samples. ",
        "introduction": "Angular and radial maps of the Local Universe are the first step for astronomers to investigate the nature of large scale structures observed in the galaxy distribution. An early discovery of a large scale structure was the \"great wall\" of galaxies, in the CfA1 redshift survey  \\cite{CFA} and \\cite{CFA1}.  Building on this promising technique for tracing 3-D structures, deeper redshift surveys including: CfA2 \\cite{CFA2}, IRAS, LCRS, 2dF \\cite{2DF}, and SDSS \\cite{DR1}, all show new more extensive and complex structures such as walls, chains, voids, and superclusters.  A related approach to quantify the complexity of the galaxy distribution is to construct the galaxy number counts in mag, N(m), and in distance, N(r), using statistical indicators: spatial and angular two-point correlation functions.  Combining approaches, physical interpretation of these counts leads to the understanding of the nature of large scale structure, and to its cosmological implications for the models of formation of galaxies.  Much attention is focused on the determination of the scale of homogeneity in the universe.  It is not a trivial concern, some statistical indicators are interpreted as showing the scale of homogeneity to be as small as 15 Mpc; while the new redshift surveys show extended structures larger than 100 Mpc, dwarfing even cluster sized stucture. In this present paper, we will not directly address this issue, but instead focus on the nearby Universe, creating a global picture for the redshift and number count distribution of galaxies.    The paper is organized as follows : section 2 presents the data, the completness studies, the number counts and the construction of volume limited subsamples; section 3 shows the 2D and 3D maps of the Local Universe in various coordinates systems and preliminary comparison with SDSS DR1 and section 4 gives our conclusions. ",
        "conclusions": "With the astonishing large scale structure observed by 2MASS, SDSS, 2dF and through the LEDA archive, it is increasingly clear that the cosmological models must account for structures ten times larger than previous convention, upwards of 300 Mpc. The second idea is that we cannot simply estimate the scale at which one is looking at the structures, just by looking at the maps : we are seeing similar structures at different scales. This could be another argument in favor of using general tools that can describe both homogeneous and scale invariant distributions, to analyse the galaxy point distribution."
    },
    "0403/astro-ph0403259_arXiv.txt": {
        "abstract": "We discuss spatial fluctuations in the gravitational wave background arising from unresolved Galactic binary sources, such as close white dwarf binaries, due to the fact the galactic binary source distribution is anisotropic. We introduce a correlation analysis of the two data streams of the Laser Interferometer Space Antenna (LISA) to extract spherical harmonic coefficients, in an independent manner, of the hexadecapole moment ($l=4$) related to the projected two-dimensional density distribution of the binary source population. The proposed technique complements and improves over previous suggestions in the literature to measure  the gravitational wave background anisotropy based on the time modulation of data as LISA orbits around the Sun. Such techniques, however, are restricted  only to certain combinations of spherical harmonic coefficients of the galaxy with no ability to separate them individually. With LISA, $m=2,3$ and $4$ coefficients of the hexadecapole ($l=4$) can be measured with signal-to-noise ratios at the level of 10 and above in a certain coordinate system. In addition to the hexadecapole coefficients, when combined with the time modulation analysis, the correlation study can also be used, in principle, to measure quadrupole coefficients of the binary distribution. ",
        "introduction": "In addition to gravitational waves from massive black holes at cosmological distances, the Laser Interferometer Space Antenna (LISA; \\cite{lisa}) is expected to detect the galactic binary source background, such as close white dwarf binaries, with gravitational waves at frequencies between 0.1 mHz to 100 mHz \\cite{webb}. At frequencies above 3 mHz, LISA will resolve a large number of galactic binaries \\cite{Neletal01}, while at the low frequency end, the unresolved population will form a confusion background of gravitational waves. This background is expected to be anisotropic given the fact that galactic binaries trace the density distribution of our galaxy such that the cumulative gravitational wave flux  is expected to be concentrated towards the highest density regimes such as the disk and the bulge of the Milky Way. As LISA orbits around the Sun, its response will be sensitive to different regions on the sky and one expects the data stream to modulate as a function of time. This modulation, in return, can be used to extract information related to the density distribution and to reconstruct the anisotropy of the gravitational wave background \\cite{gia,Cornish:2001hg,Ungarelli:2001xu}. Note that the binary background can be considered as both a source of noise or a signal; For example, by treating the background as a signal, we can attempt to obtain certain information related to the Galactic structure through anisotropies of the background \\cite{gia,Cornish:2001hg,Ungarelli:2001xu}. On the other hand, if the background is treated as a source of noise, we might also discuss how anisotropies affect the data analysis and consider certain strategies to minimize any impact \\cite{Seto04}.  Here, we treat the background as a signal and consider the presence of anisotropies and their potential measurement using a correlation between the two data streams of LISA. For this purpose, following Ref. \\cite{Seto04}, we introduce a spherical harmonic moment analysis of LISA data in terms of spatial inhomogeneities in the two-dimensional projected density distribution when integrated over the observer to source distance. We also emphasize the subtle use of a freedom related to data combinations that also corresponds to an effective rotation angle of LISA detectors in the detector plane. The technique we discuss here compliments, and improves, studies based on the time modulation of the data stream and allows one to extract individual coefficients of multipole moments related to the two-dimensional binary source distribution; Prior techniques were only restricted to combinations of coefficients related to the monopole ($l=0$), quadrupole ($l=2$), and the hexadecapole ($l=4$). Considering LISA observational parameters, we show that certain coefficients of the hexadecapole multipole moments can be extracted with signal-to-noise ratios at the level of 10 and above.  We also study the complimentary nature of the two methods involving time modulation and the correlation of data streams and suggest that, in combination, one can also, in principle, extract information related to coefficients of the quadrupole distribution. The only multipole moment coefficients that will remain individually undetermined are ones of $l=0$,$m=0$ mode (monopole) and the $l=2, m=0$ mode related to the quadrupole. These two, however, can be determined, in combination, as a sum. The correlation analysis would be one of the crucial elements for LISA's follow-on missions aiming to  detect weak gravitational  wave background from the early universe \\cite{Ungarelli:2000jp,Allen:1996vm}. While it is expected that the  basic design of  any potential missions in the long term is likely to be largely different from that of LISA, the correlation analysis with  LISA data streams may present an important first step for any planned studies in the future that attempt to use similar procedures to both extract signals or to remove foreground noise.  The discussion is organized as follows: in the next Section, we outline the spherical harmonic formulation of the anisotropies in the gravitational wave background and how the low order moments, $l=0,2$ and 4, are related to the LISA data streams. In Section~3, we discuss the correlation between data streams as a way to extract information related to these anisotropies beginning with a brief discussion of prior proposed techniques related to modulations in the data streams. In Section~4, We conclude with a summary of our main results. For background information related to the technique presented in this paper, we refer the reader to Ref.~\\cite{Seto04}. ",
        "conclusions": "In this paper, we have discussed spatial fluctuations in the gravitational wave background arising from unresolved Galactic binary sources, such as close white dwarf binaries, due to the fact the this galactic source distribution, when projected on the sky, is anisotropic. We introduce a spherical harmonic analysis of this anisotropy  based on a correlation study of the two  data streams of LISA and propose the measurement of individual spherical harmonic coefficients related to the hexadecapole moment ($l=4$), in addition to coefficients of the quadrupole ($l=2$).  The proposed technique complements and improves over previous suggestions in the literature to measure  the gravitational wave background anisotropy based on the time modulation of the single data stream ({\\it e.g.} $S^B_{AA}(t)$) as LISA orbits around the Sun and restricted  only to combinations of spherical harmonic coefficients with no ability to separate them. The correlation analysis, by making use of a single freedom related to the rotation angle of detectors in the detector plane or similarly how data are combined,  provides a method to independently extract multipole coefficients of the hexadecapole moment of the background. The information related to the quadrupole comes from the correlation analysis in combination with the time modulation. The only independently undetermined coefficients turn out to be that of the monopole and the $m=0$ mode of the quadrupole, though the two can be determined as a sum.  With LISA, we have shown that certain coefficients of the hexadecapole can be measured with signal-to-noise ratios at the level of 10 and above. While LISA will present a first step towards a correlation analysis of the gravitational wave background, future missions are likely to exploit this in detail both as a source of signal, to extract spatial information, and noise, to decrease noise to detect almost isotropic backgrounds such as the one related to primordial gravitational waves. This would certainly be the case at low frequencies as the foreground is likely to be the dominant source of anisotropy. At few mHz frequencies, the fluctuation amplitude related to the extragalactic component is expected to be at the level of a few percent, but restricted to arcminute angular scales where galaxy density field is observed to vary. At $\\mu$Hz frequencies, where the merging massive black holes are likely to be the domiant foreground source, the anisotropy will be dominated by the shot-noise related to the finite number density of such binaries on the sky. On the other hand, the anisotropy of the primordial gravitational wave background is expected to have a pattern close to that of the cosmic microwave background, with a fluctuation level of $\\delta \\sim 10^{-5}$, at large angular scales. A confirmation of their similarity in anisotropy can be used as a proof on the primordial nature of the background. For a direct detection of spatial fluctuations in the gravitational wave background, in  the optimistic case where the detector noise is dominated by the monopole mode of the background, one is forced to gravitational wave frequencies at the level of $f\\sim \\Delta f\\sim 10^4 ({\\rm SN}/10)^2 (\\delta/10^{-5})^{-2}(T/3{\\rm yr})$Hz. Thus, for a detection of the anisotropy at a level $\\delta\\sim 10^{-5}$ over a realistic observational period with significant signal to  noise ratio, the relevant frequencies are at $f\\gsim 10$kHz. Interestingly, at this high end of the gravitational wave spectrum, there are no viable astrophysical sources and suggests that, in future, one may in fact be able to extract anisotropy associated with the primordial gravitational wave background."
    },
    "0403/astro-ph0403284_arXiv.txt": {
        "abstract": "We present spectropolarimetric observations of the FK~Com star HD~199178 obtained between 1998 December and 2003 August at the T\\'elescope Bernard Lyot (Observatoire du Pic du Midi, France). We report the detection of a photospheric magnetic field and reconstruct its distribution by means of Zeeman-Doppler Imaging. We observe large regions where the magnetic field is mainly azimuthal, suggesting that the dynamo processes generating the magnetic activity of \\hd\\ may be active very close to the stellar surface. We investigate the rapid evolution of surface brightness and magnetic structures from a continuous monitoring of the star over several weeks in 2002 and 2003. We report that significant changes occur in the distribution of cool spots and magnetic regions on typical timescales of the order of 2 weeks. Our spectropolarimetric observations also suggest that the surface of \\hd\\ is sheared by differential rotation, with a difference in rotation rate between equatorial and polar regions of the order of 1.5 times that of the Sun. ",
        "introduction": "\\label{sect:introduction}  Our understanding of the Solar dynamo has benefited from important progresses during the last decade, thanks to the combined breakthroughs of observational and modeling techniques. Helioseismological studies are now able to reveal the internal velocity field of the Sun (Schou et al. 1998) and recent MHD simulations give credence to the idea that the dynamo processes at the origin of the large-scale solar field are mainly operating within a thin interface layer separating the radiative core from the convective envelope of the Sun~: the tachocline. This region allows an efficient transformation of a seed poloidal field into a toroidal component as well as a storage of the field on a timescale of the order of the solar magnetic cycle (Rempel et al. 2000). The poloidal component of the field is believed to be regenerated (with opposite polarity) through the action of the Coriolis force, by means of a still mostly unknown mechanism called ``$\\alpha$-effect'', which may also be mostly efficient within the tachocline (Dikpati \\& Gilman 2001).  The magnetic field observed at the photospheric level of the Sun was first detected in sunspots (Hale 1908) which are the largest magnetic structures of the solar surface (though only covering about $10^{-4}$ of the photosphere at solar maximum), featuring mostly radially oriented field lines. Magnetic elements are now observed at much smaller scales in the so-called quiet photosphere (see Solanki 2001 for a review of the properties of such magnetic regions). The interplay, during the solar cycle, between the small magnetic structures and the large-scale field has not yet been understood, whereas small magnetic elements store a significant part of the magnetic energy of the photosphere, and despite the fact the distribution of some of these structures has a temporal evolution correlated to the solar cycle. Some studies also suggest that a turbulent dynamo, disconnected from the large-scale solar dynamo, may efficiently generate small-scale magnetic structures (Cattaneo 1999).  Surprisingly, these properties of the solar photospheric field are far from the general features observed on other active stars, at least among the fast rotator group. The surface magnetic field of several fast rotating stars has been mapped by means of Zeeman-Doppler Imaging (thereafter ZDI, Semel 1989). For these extremely active objects, a field is generally observed at basically all locations of the photosphere (see e.g. Donati et al. 2003a), in the form of discrete structures that present an homogeneous orientation of field lines and cover a significant fraction of the stellar surface. On some objects, the surface field is dominated by an azimuthal component distributed in arcs or even in complete rings encircling the rotation axis at different latitudes. These observations suggest that the dynamo operating within these objects is very different from that of the Sun, since the azimuthal component of the solar field is supposed to be deeply buried at the base of the convective envelope. Its presence at the photospheric level on fast rotators therefore suggests that their dynamo may be active very close to the surface.  The observing effort engaged during the last decade to investigate, by means of high-resolution spectropolarimetry, the magnetic activity of active late-type fast rotators has been concentrated on several pre-main sequence stars and on the evolved close binary system HR~1099. The present study focuses on \\hd, a member of the small FK~Com group. Objects belonging to this class are single fast rotating late-type giants displaying signs of a strong magnetic activity. Their short rotation periods (of order of a few days) suggest that FK~Com stars have undergone very unusual mechanisms during their evolution. In particular, it was proposed by Bopp \\& Stencel (1981) that they may result from the coalescence of contact binaries. Long-term photometric studies of FK~Com itself (Jetsu et al. 1991, Korhonen et al. 2002) report that its cool spot distribution is asymmetric (a so-called ``active longitude'' dominating the spot activity), with occasional 180\\degr\\ shifts of the active longitude (a behavior dubbed ``flip-flop'') resulting in a 6.5 years cycle. The distribution of cool spots of \\hd\\ has also been investigated by means of Doppler imaging (Strassmeier et al. 1999, Hackman et al. 2001).  The aim of the present study is to enrich our knowledge of this object by reconstructing its magnetic topology altogether with its spot distribution, in a twofold purpose. We first plan to compare the general properties of the photospheric magnetic field of this evolved star with the magnetic topologies mapped for younger objects. We also intend to check whether active longitudes reported by Hackman et al. (2001) are also observed in our own maps of the cool spot distribution and test whether they have a counterpart in the magnetic images.  We first summarize the series of spectropolarimetric observations used for this study and describe the modeling procedures employed to reconstruct the magnetic topology of the star at every observing epoch. We then present the main characteristics of its surface brightness and magnetic structures and study their evolution on different timescales (from a few weeks to several years), including an analysis of the surface \\drot. We finally summarize the main results and outline the informations they provide on the nature of the underlying dynamo processes. ",
        "conclusions": "\\label{conclusion}  This study reports the detection of a photospheric magnetic field on the FK~Com giant \\hd. The spatial distribution of the field is reconstructed by means of ZDI. The magnetic topology of \\hd\\ presents at all observing epochs large regions in which field lines are mostly azimuthally oriented. This observation, together with similar ones reported for other fast rotators (Donati et al. 2003a), suggests that the dynamo processes generating the magnetic activity of \\hd\\ may be active very close to the stellar surface.  Local short-term evolution of surface brightness and magnetic structures are observed on timescales as short as a couple of weeks. A polarity reversal is reported during summer 2002 for high-latitude magnetic regions hosting a radial field. A slow increase of the fraction of magnetic energy stored at low-latitude in the large-scale azimuthal component is also observed. Beside these local evolutions, the largest features of the brightness and magnetic topologies remain stable over several years. Although the observed surface changes remain mostly limited to small structures on our 4.5~yr monitoring, future observations may tell us whether the magnetic field of \\hd\\ undergoes a cyclical evolution similar to that observed on the Sun.  It is finally suggested that the surface of \\hd\\ is sheared by a solar-like differential rotation. The difference in rotation rate between equatorial and polar regions is reported to be about 1.5 times that of the Sun.  The observations reported here are at the limit of the capacity of the MuSiCoS spectropolarimeter, in terms of \\sn\\ and spectral resolution. Much more accurate observations are therefore expected for the same star when the new generation of spectropolarimeters becomes available, with ESPaDOnS at the Canada-France-Hawaii Telescope (Donati 2003) and NARVAL at TBL (Auri\\`ere 2003). Better data may in particular allow us to probe more discreet photospheric effects, like possible secular fluctuations in the amount of differential rotation, as already reported for other active stars (Donati et al. 2003b)."
    },
    "0403/astro-ph0403551_arXiv.txt": {
        "abstract": "Observations show that the slope of the inner part of the H$\\alpha$/[NII] rotation curves of disk galaxies is depressed by extinction: at fixed luminosity, the observed slope is in fact seen to depend on the disk inclination to the line of sight. Using a simple extinction model, we are able to reproduce well the observed trends. The model assumes an exponential distribution, both in the radial and $z$ directions, identical for star--forming regions and dust. Fits to the data are optimized by varying the scale height and scale length of absorbers, as well as the disk's central optical depth $\\tau_\\circ$, as seen face--on. The observations indicate that disk central opacity increases with total luminosity. Model fits for the most luminous galaxies (brighter than $M-5\\log h=-21.4$ in the I band) yield $\\tau_\\circ \\simeq 3.5^{+4.0}_{-2.0}$, near $\\lambda=0.66$ $\\mu$m. The large uncertainty on the inferred value of $\\tau_\\circ$ is due to the poorly known characteristics of the distribution of absorbers in the disk, as well as to the likelihood of strong departures from an exponential radial distribution near the galaxy centers. ",
        "introduction": "Scattering and absorption by interstellar dust depresses the observed flux of galaxy disks at optical and infrared wavelengths. While the magnitude of the effect has been disputed, there is no doubt that the effect is strongly dependent on the inclination of a disk to the line of sight. A number of techniques have been used in observationally deriving extinction laws (see for example the volume edited by Davies and Burstein 1995, the review by Calzetti 2001 and Masters, Giovanelli \\& Haynes 2003 for details). Goad and Roberts (1981) first discussed a kinematical technique which we briefly describe as follows.  Let $V(r)$ be the rotational velocity of an axially symmetric disk at the distance $r$ from its center and $(x,y)$ be a set of Cartesian coordinates in the plane of the disk. If the disk is thin and its symmetry axis is inclined by an angle $i$ to the line of sight, the observed component of velocity which intercepts the disk at $(x,y)$ is \\be V_\\parallel = V(r) {x\\over \\sqrt{x^2+y^2}} \\sin i + V_{nc}, \\ee where $V_{nc}$ accounts for non--circular motions and the $x$ axis is oriented along the disk's apparent major axis. The observed rotation curve, as derived for example from a long--slit \\hal ~spectrum positioned along the major axis, will depart from the prescription given by Eqn. 1 as due to seeing, instrumental resolution, averaging across the slit  width, the finite thickness of the disk and extinction occurring within the disk itself. As a realistically thick disk approaches the edge--on perspective, lines--of--sight along the major axis sample regions of increasingly broad range in $y$, yielding a velocity distribution with a peak velocity contributed by parcels of gas asymmetrically distributed with respect to $y=0$: if extinction is important, only foreground parts of the disk contribute to the emission and the factor $(x/ \\sqrt{x^2+y^2})<1$ depresses the velocity distribution observed at $r=x$. Goad \\& Roberts noted how this tapering effect may, in opaque edge--on disks, yield observed rotation curves resembling solid--body behavior, independently of the true shape of $V(r)$. Bosma \\etal (1992) tested this technique on two edge--on systems, NGC 100 and NGC 891. They compared HI synthesis and \\hal ~observations and concluded that the disk of NGC 100 is transparent, while in the case of NGC 891 they could not exclude the possibility of extinction in the inner parts of the disk. Their conclusions were however strongly influenced by the limited spatial resolution of the HI data. Prada \\etal (1994) compared long--slit spectra of NGC 2146 in the optical and near IR and reported evidence for extinction in the inner parts of the galaxy. More recently, Giovanelli \\& Haynes (2002, hereafter GH02) applied the same technique in a statistical fashion, detecting clear evidence for the effect of extinction at 0.66 \\m. Using a sample of more than 2000 \\hal ~rotation curves, they found that the opacity in the inner disks of spiral galaxies is luminosity dependent, in a manner previously found by purely photometric means (Giovanelli \\etal 1995; Tully \\etal 1998). The observational data of GH02 are corroborated by the detailed radiative transfer calculations of Baes et al.(2004), which also obtain dramatic increases in the apparent inner scale length of rotations curves as disks approach the edge--on perspective.  In this report, we use a simple model for the dust gas distribution in spiral disks, in order to reproduce the findings of GH02 and in so doing to obtain quantitative inferences on the characteristics of the disk opacity. In Section 2 we summarize the observational results that are to be modelled. In Section 3 we describe the disk model to be fitted to the data, while the best fit parameters are discussed in Section 4. In Section 5 we present our conclusions. Throughout this work, distance dependent quantities are scaled according to a Hubble parameter $H_\\circ=100 h$ \\kmsm.    \\section {Rotation Curve Shape and Disk Inclination}  GH02 fitted \\hal/[NII] rotation curves of spiral galaxies with the parametric model \\be V_{pe}=V_\\circ (1-e^{-r/ r_{pe}})(1+\\alpha r/ r_{pe}) \\label{vpe} \\ee where $V_\\circ$ scales the amplitude of the rotation curve, $r_{pe}$ yields a scale length for the inner steep rise, and $\\alpha$ corresponds to the slope at large $r$. In their Table  1, for different luminosity classes as measured in the I--band they list mean values of $\\alpha$, $h r_{pe}$ and of the ratio between $hr_{pe}$ and the scale length of the disk light for relatively face--on systems. In their Figure 1, for each luminosity class the dependence of $h r_{pe}$ on disk inclination $i$ is shown. For a given luminosity class, with increasing $i$, $<h r_{pe}>$ remains at first constant; as $i$ approaches $70^\\circ$, $<h r_{pe}>$ starts increasing, the rate of increase being larger for more luminous systems. In the latter, $<h r_{pe}>$ at high inclination has values more than double the values observed at low inclination.  The inclination dependence was interpreted as due to the increase in the opacity of spiral disks with inclination, while the luminosity dependence was attributed to a combination of increased disk size and higher interstellar medium metallicities in the more luminous systems. By contrast, the outer slope $<\\alpha>$ is unaffected by disk inclination. ",
        "conclusions": "A simple model, whereby dust and HII regions are distributed exponentially in both the radial and $z$ directions, with equal scale height and scale length, can reproduce the observed variation in the inner slope of H$\\alpha$/[NII] rotation curves of disk galaxies. The effect is due to internal extinction. For reasonable values of the ratio between scale height and scale length, the model can constrain central disk opacities for the most luminous galaxies, in which the effect of extinction is the strongest. However, pinning down the exact value of central opticall depth of the disk at 0.66 $\\mu$m is difficult, due to uncertainty in the form of the distribution of absorbers. If \\hal ~and the absorbers have scale length about 1.9 times that of the I band starlight and scale height to scale length ratio of $\\sim 0.05$, we infer central optical depth of disk, in the face--on perspective of $\\tau_\\circ\\simeq 3.5$, for galaxies brighter than $M-5\\log h=-21.4$ in the I  band. However, the uncertainty in the radial and $z$--distributions of absorbers constrains the inferred values of $\\tau_\\circ$ to an uncertainty no better than a factor of 2. For less luminous spiral galaxies, extinction becomes progressively less important, and the model cannot usefully constrain the disk opacity."
    },
    "0403/astro-ph0403598_arXiv.txt": {
        "abstract": "Our recent progress on time-dependent modeling of the multiwavelength spectra and variability of blazars with leptonic and hadronic jet models is reviewed. Special emphasis is placed on X-ray spectral variability of low-frequency peaked (LBLs) and intermediate BL Lac objects (IBLs). As an example, recent observational and modeling results of an extensive multiwavelength campaign of BL~Lacertae in 2000 are presented. It is demonstrated how combined spectral and variability modeling of LBLs and IBLs can significantly constrain emission models and potential variability scenarios. In the case of BL Lacertae, the variability appears to be driven primarily by fluctuations in the spectral index of the non-thermal, ultrarelativistic electron population in the jet. Such constraints allow us to refine predictions of the intrinsic GeV $\\gamma$-ray emission and the dominant electron cooling mechanism in these objects. ",
        "introduction": "Blazars are the most extreme class of active galaxies known to date. They have been observed in all wavelength bands --- from radio through very-high energy (VHE) $\\gamma$-ray frequencies. 66 blazars have been identified as sources of $> 100$~MeV emission detected by the EGRET telescope on board the {\\it Compton Gamma-Ray Observatory} \\citep[CGRO;][]{hartman99}, and 6 blazars (Mrk~421: \\citet{punch92,petry96}; Mrk~501: \\citet{quinn96,bradbury97}; PKS 2155-314: \\citet{chadwick99}; 1ES~2344+514: \\citet{catanese98}; 1H~1426+428: \\citet{horan02,aharonian02}; 1ES~1959+650: \\citet{nishiyama99,holder03,aharonian03}) have now been detected at VHE $\\gamma$-rays ($> 350$~GeV) by ground-based air \\v Cerenkov telescopes. Many of the EGRET-detected $\\gamma$-ray blazars appear to emit --- at least temporarily --- the bulk of their bolometric luminosity at $\\gamma$-ray energies. Blazars exhibit variability at all wavelengths on time scales --- in some cases --- down to less than an hour. VLBI radio observations and monitoring often reveal one-sided kpc-scale jet structures, exhibiting apparent superluminal motion. The radio through optical emission from blazars often shows linear polarization, pointing towards a synchrotron origin.  \\begin{figure*} \\includegraphics[width=\\linewidth]{blazarfits.eps} \\caption{Broadband spectra of the FSRQ 3C~279 \\citep[left panel;][]{hartman01,collmar04}, the LBL BL~Lacertae \\citep[middle panel;][]{boettcher03}, and the HBL Mrk~501 \\citep{petry00}. For each object, two simultaneous SEDs at two different epochs are shown, illustrating the range of long-term spectral variability seen in these objects. For 3C~279, the results of recent multiwavelength observations in 2003, including INTEGRAL \\citep{collmar04} are also shown.} \\label{bbspectra} \\end{figure*}  \\subsection{\\label{obs_spectra}Blazar Spectra}  The broadband continuum spectra of blazars are dominated by non-thermal emission and consist of at least two clearly distinct, broad spectral components. A sequence of sub-classes of blazars can be defined through the peak frequencies and relative $\\nu F_{\\nu}$ peak fluxes of those components, which also appear to be correlated with the overall bolometric luminosity of the sources \\citep{fossati98}: In the case of flat-spectrum radio quasars (FSRQs), the low-frequency (synchrotron) component extends from radio to optical/UV frequencies, with a peak frequency generally in the mm or IR band; the high-frequency component extends from X-rays through GeV $\\gamma$-ray energies, with a $\\nu F_{\\nu}$ peak frequency corresponding to $\\sim 10$~MeV --- 1~GeV. No FSRQ has so far been detected by ground-based air \\v Cerenkov telescope facilities at energies $> 100$~GeV, although in flaring states the $\\gamma$-ray $\\nu F_{\\nu}$ peak flux of FSRQs dominates over the low-frequency emission by up to $\\sim 1$ order of magnitude. In the case of high-frequency peaked BL~Lac objects (HBLs), the low-frequency component often extends far into the X-rays, with peak frequencies ranging from the UV/soft X-ray to the hard X-ray regime \\citep{pian98}, depending on the source and its state of activity; the high-energy component of HBLs extends from hard X-rays far into the VHE $\\gamma$-ray regime. All blazars detected at VHE $\\gamma$-ray energies to date are HBLs. In spite of extending to extremely high photon energies, the $\\nu F_{\\nu}$ peak flux of the $\\gamma$-ray component of HBLs is generally at most comparable to the spectral output in the low-frequency component. In terms of their overall bolometric luminosity, FSRQs appear to be several orders of magnitude more powerful sources than HBLs.  Apparently intermediate between the FSRQs and the HBLs are the low-frequency peaked BL~Lac objects (LBLs). The peak of their low-frequency component is typically located at IR or optical wavelengths, their high-frequency component peaks typically at $\\sim$~several GeV, and the $\\gamma$-ray output is of the order of or slightly higher than the level of the low-frequency emission. Fig. \\ref{bbspectra} shows a compilation of broadband spectra of a typical FSRQ (3C279), an LBL (BL~Lacertae), and an HBL (Mrk~501), respectively. Each panel shows the respective object at two or three different observing epochs, in two different activity states.  \\subsection{\\label{variability}Blazar Variability}  Fig. \\ref{bbspectra} already illustrates that in particular the high-energy emission from blazars can easily vary by more than an order of magnitude between different observing epochs. However, variability has been observed on much shorter time scales, in some extreme cases less than an hour \\citep{gaidos96}. Fig. \\ref{bllac_lc} shows examples of light curves of the LBL BL~Lacertae, taken during a broadband observing campaign carried out in the second half of 2000 \\citep{boettcher03,villata02}. While the radio emission of blazars is generally variable on time scales of weeks -- months (Fig. \\ref{bllac_lc}a), the optical light curve shows significant variability on time scales down to $\\sim 1.5$~hr (Fig. \\ref{bllac_lc}b).  \\begin{figure*} \\includegraphics[width=\\linewidth]{bllac_lightcurves.eps} \\caption{Radio and optical light curves of the LBL BL~Lacertae during the multiwavelength campaign of 2000. The dark-shaded areas in panel a) indicate the time segments of BeppoSAX observations.} \\label{bllac_lc} \\end{figure*}  Often, both the optical and X-ray emission show characteristic hardness-intensity correlations. Fig. \\ref{spectralvar} illustrates this for BL~Lacertae in 2000. Some HBLs (e.g., Mrk~421 and PKS~2155-304) have been observed to exhibit characteristic, clockwise loop structures \\citep[``spectral hysteresis'';][]{takahashi96,kataoka00}, which can be interpreted as the synchrotron radiation signature of gradual injection and/or acceleration of ultrarelativistic electrons into the emitting region, and subsequent radiative cooling \\citep{kirk98,gm98,kataoka00,kusunose00,lk00}.  In LBLs, the soft X-ray emission is also sometimes dominated by the high-energy end of the synchrotron emission component, so similar spectral hysteresis phenomena should in principle be observable. However, those objects are generally much fainter at X-ray energies than their high-frequency peaked counterparts, making the extraction of time-dependent spectral information an observationally very challenging task. Fig. \\ref{spectralvar} clearly illustrates that the {\\it Beppo}SAX observations of BL~Lacertae in 2000 revealed evidence for spectral variability, but lacked the sensitivity to clearly establish or rule out spectral hysteresis. Such a measurement might require the new generation of X-ray telescopes such as {\\em Chandra} or {\\em XMM-Newton}. ",
        "conclusions": ""
    },
    "0403/astro-ph0403621_arXiv.txt": {
        "abstract": "We report on a very large set of simulations of collisions between two main sequence (MS) stars. These computations were done with the ``Smoothed Particle Hydrodynamics'' method. Realistic stellar structure models for evolved MS stars were used. In order to sample an extended domain of initial parameters space (masses of the stars, relative velocity and impact parameter), more than 14\\,000 simulations were carried out. We considered stellar masses ranging between 0.1 and $75\\,\\Msun$ and relative velocities up to a few thousands {\\kms}. To limit the computational burden, a resolution of 1000--32\\,000 particles per star was used. The primary goal of this study was to build a complete database from which the result of any collision can be interpolated. This allows us to incorporate the effects of stellar collisions with an unprecedented level of realism into dynamical simulations of galactic nuclei and other dense stellar clusters. We make the data describing the initial condition and outcome (mass and energy loss, angle of deflection) of all our simulations available on the Internet. We find that the outcome of collisions depends sensitively on the stellar structure and that, in most cases, using polytropic models is inappropriate.  Published fitting formulae for the collision outcomes, established from a limited set of collisions, prove of limited use because they do not allow robust extrapolation to other stellar structures or relative velocities. ",
        "introduction": "\\label{sec:intro}   \\subsection{Stellar collisions in galactic nuclei} \\label{sec:StellCollInNat}  \\begin{figure*} \\begin{center} \\resizebox{0.9\\hsize}{!}{% \\includegraphics{TcollTrel_NoyGal_simple.eps} } \\end{center} \\caption{ Relaxation and collision times at 0.1\\,pc from a massive black hole in the centre of a galactic nucleus (inspired from a similar diagram by \\citealt{Arab97}). We plot curves of iso-$T_{\\mathrm{relax}}$ (dashes, blue in colour version) and iso-$T_{\\mathrm{coll}}$ (solid lines, magenta in colour version) in a plane parameterised by the stellar density at 0.1\\,pc and the mass of the central black hole. The right ordinate scale indicates the stellar velocity dispersion (Keplerian value, $\\sigma_{V}=\\sqrt{GM_{\\rm BH}/r}$).  All stars are assumed to be of solar type. The left shaded sector (cyan in colour version) corresponds to conditions for which both $T_{\\mathrm{relax}}$ and $T_{\\mathrm{coll}}$ are larger than Hubble time, so that such nuclei are not expected to show signs of secular evolution. In the shaded lower right corner, the black hole does not dominate the kinematics and the effect of the cluster's self gravitation should be included in the computation of the characteristic time-scales. \\newline Large black dots show the estimated conditions for observed galactic nuclei. In most cases, the estimation of the stellar density at 0.1\\,pc requires important extrapolation from the data, as such a small radius is resolved only for a few galaxies of the local group (the Milky Way, M\\,31 and M\\,32). For this extrapolation, we use a power-law cusp of the form $\\rho(r)=\\rho_0(r_0/r)^\\gamma$. The values of $\\rho_0$, $r_0$ and $\\gamma$ are taken from \\citet{GebhardtEtAl03} or \\citet{Faber97}. The densities for M\\,31 and M\\,32 are from \\citet{Lauer98} and the Milky Way's value from \\protect\\citet{GTKKTG96}. The values of $M_{\\mathrm{BH}}$ are estimates by \\protect\\citet{Mag98}, \\protect\\citet{vdM99c} or better constrained values gathered by Kormendy (\\citeyear{Kormendy04}, see \\url{http://chandra.as.utexas.edu/~kormendy/bhsearch.html} for these data and a list of original references). In some cases, $\\rho_0$ and $r_0$ have already been extrapolated from larger radii! Cases with an horizontal line connected to a second smaller dot are nuclei for which the slope $\\gamma$ is observational compatible with $0$, according to \\protect\\citet{Faber97}. The second point indicates the density value at 0.1\\,pc if constant $\\rho$ is assumed up to $r_0$.  } \\label{fig:TrelTcollNG} \\end{figure*}  In the recent years, the study of stellar collisions has received renewed interest from researchers studying the dynamics of dense stellar systems, either open/globular clusters or galactic central regions (see the contributions in \\citealt{Shara02}). Our own motivation is to perform simulations of the long-term evolution of dense stellar systems, particularly galactic nuclei, with a new Monte Carlo stellar dynamics code which incorporate collisions as ``micro-physics'' \\citep{FB01a,FB02b}.  Before going into a brief description of the astrophysical motivations of these works, some clarification about the notion of ``stellar collision'' is called for. In this paper we shall use this term to refer to a process during which two stars, previously unbound to each other, get so close that not only gravitational forces but also hydrodynamical ones come into play to determine the outcome of the interaction. So, strictly speaking, collisions are not only contact encounters but also comprise tidal interactions. However, for reasons to be exposed in Sec.~\\ref{subsec:tid_bin}, we restrict here to events leading to physical contact at first periastron passage.  To assess the importance of collisions in a given astrophysical context, the key quantity to monitor is the collision time, which we define here as the average time needed for ``test-star'' 1 to experience a collision with any ``field-star'' 2, \\begin{eqnarray} T_\\mathrm{coll} &=& \\left( n_2 S v_\\mathrm{rel} \\right)^{-1} \\label{eq:Tcoll}\\\\ \\mbox{with\\ } S &=& \\pi d_\\mathrm{coll}^2\\left( 1 + \\frac{2G(M_1+M_2)}{d_\\mathrm{coll}v_\\mathrm{rel}^2} \\right), \\label{eq:Scoll} \\end{eqnarray} where $n_2$ is the number density of the field stars, $v_\\mathrm{rel}$ the relative velocity and $d_\\mathrm{coll}$ the pericentre distance leading to a collision ($d_\\mathrm{coll} = R_1 + R_2$ for contact at first passage, neglecting tidal deformation). $S$ is the collisional cross-section. At low relative velocity, it is greatly enhanced over the geometric value by gravitational attraction. This effect, dubbed ``focusing'' is expressed by the second term in the brackets of Eq.~\\ref{eq:Scoll}. In most astronomical contexts, the velocity dispersion is much smaller than the stellar escape velocity $V_{\\star}=\\sqrt{2GM_{\\star}/R_{\\star}}$ ($\\simeq 620$\\,{\\kms} for sun-like stars) and gravitational focusing dominates. In these cases, integrating over a Maxwellian distribution for relative velocities yields \\cite[Eq.~8--125]{BT87}: \\begin{eqnarray} T_\\mathrm{coll} &\\simeq& 7\\times10^{14}\\,\\mathrm{yrs}\\,\\times \\\\ && \\nonumber \\left(\\frac{R_{\\star}}{{\\Rsun}}\\right)^{-1} \\left(\\frac{M_{\\star}}{{\\Msun}}\\right)^{-1} \\left(\\frac{n}{\\mathrm{pc}^{-3}}\\right)^{-1} \\left(\\frac{\\sigma_v}{\\mathrm{km}\\,\\mathrm{s}^{-1}}\\right). \\end{eqnarray}  In systems with $T_\\mathrm{coll}$ smaller than typical stellar ages, collisions have expectedly imprinted not only the stellar population but also the global dynamical structure. Very high densities are necessary for such situations to take place but even when collisions occur at frequencies too low to be of dynamical relevance, they still can be of great astrophysical interest {\\it per se} because they are suspected to lead to the formation of unusual individual stellar objects, such as blue stragglers or stripped giants \\citep[and references therein]{Davies96,Shara99}. Collisions are unimportant in the bulk of a galaxy; the probability for the sun to suffer a collision during its 10~Gyr main sequence life, amounts to no more than $10^{-7}$! Only in stellar clusters and galactic nuclei, is there a non-vanishing probability for at least some stars to experience collisions. For reviews about the role of collisions in various environments, we refer to the various papers in \\citet{Shara02}.   Among known stellar environments, galactic nuclei are those in which the most extreme values of the stellar density and velocity dispersion are attained. The best known case is our own Galaxy. Inside a sphere of radius 0.4\\,pc at the Galactic centre, the stellar density exceeds $4\\times10^6\\,{\\Msun}\\mathrm{pc}^{-3}$ and a velocity dispersion of order 500\\,km\\,s$^{-1}$ has been reported at a distance of 0.01\\,pc of the $2-3\\times 10^6\\,{\\Msun}$ central black hole \\citep{GTKKTG96,GPEGO00}. Most other galactic nuclei are not resolved yet so we can only produce very uncertain estimates of $T_{\\mathrm{coll}}$ for these systems.  Some of them are indicated in Fig.~\\ref{fig:TrelTcollNG}. As a bias toward our own interests, we treat only the situation of a massive black hole dominating the kinematics of the surrounding stars.  From this diagram, we see that there are very few galaxies for which we can be certain that collisions played a important dynamical role. Using Nuker model fits to represent the density profiles and the empirical relation between the mass of the central object and the velocity dispersion in the spheroidal component \\citep{TremaineEtAl02} as a proxy for the BH's mass, \\citet{Yu03} estimated the collision times for a series of observed galactic nuclei. She found only a few cases with $T_{\\mathrm{coll}}$ possibly shorter than the Hubble time and that, in present-day nuclei, collisions do not produce observable colour gradients in the stellar populations. It may be that the importance of these processes has been somewhat overestimated in the past \\citep{vdBergh65,Sanders70a}.  The centre of the Galaxy is a particularly complex and fascinating environment \\citep{GenzelEtAl03,GhezEtAl03,SchoedelEtAl03}. The ``SO'' stars orbiting the $3-4\\times 10^6\\,\\Msun$ BH {\\SgrA} at distances smaller than 0.04\\,pc seem to be on the MS with masses of at least 10\\,{\\Msun} \\citep{GhezEtAl03b}. Recent stellar formation at this place seems impossible and scenarios to bring them from a few pc away in less than their short lifetime require considerable fine tuning \\citep[][and references therein]{KFM04}. Consequently, it is tempting to hypothesise they were created in a sequence of mergers of older, lighter MS stars \\citep{GenzelEtAl03}.  Using simple Fokker-Planck modelling (not including a central BH), \\citet{Lee94,Lee96} concluded that mergers can not account for the formation of the massive stars found near the centre. On the other hand, whether collisions are responsible for the observed relative depletion of red giants at the Galactic centre is still a debated issue \\citep{Gerhard94,DBBS98,Alexander99,BD99}. Clearly, more detailed stellar dynamical models, that take into account the presence of the central BH and include a realistic treatment of collisions and stellar evolution are called for to establish the role of collisions in the MW central cluster.  \\defcitealias{SS66}{SS66}  There are however strong theoretical motivations to believe that stellar encounters may have taken place in large numbers in the past evolution in many galactic centres with sufficiently high stellar densities. The main reason is the presence of massive compact dark objects in the centre of many, if not most, galaxies. These mass concentrations are most probably supermassive black holes (SMBHs) with masses $10^5-5\\times 10^9\\,{\\Msun}$ (\\citealt{KR95,Barth04}; \\citealt*{BHS04}; \\citealt{FF04,Kormendy04,Pinkney03}). From a series of works published in the 70's \\cite[among others]{Peebles72,SL76,BW76,BW77,DO77a,DO77b,Young77a}, it is known that a SMBH-surrounding stellar system whose long-term evolution is driven by 2-body gravitational encounters will relax to a density profile, close to a power-law $\\rho(r)\\propto r^{-\\gamma}$, which yields a constant flux of stars toward the centre where they are destroyed either by tidal disruptions or energetic collisions. If all stars have the same mass, the exponent is $\\gamma=1.75$. In the innermost regions of such a cusp, a high collision rate is expected. But the collisions themselves could act as a feed-back mechanism on the evolution of the stellar system and the growth of the black hole so that the actual formation of relaxational cusp is questionable. From analytical considerations, \\citet{Frank78} concludes that collisions in the cusp are never of importance, when compared to tidal disruptions, but this statement is seriously challenged by other studies and, in particular, more recent numerical simulations (\\citealt{YSW77,Young77b,DS83,DDC87a,DDC87b}; \\citealt*{MCD91}; \\citealt{Rauch99}).  Unfortunately, the discussion of the contribution of various dynamical processes to the evolution of galactic nuclei has been blurred by uncertainties about the precise outcome of these individual processes.  For instance the amount of gas that is accreted by the SMBH following a tidal disruption is still debated \\cite[and references therein]{ALP00}. As for stellar collisions, most previous works relied on quite unrealistic prescriptions, like complete destruction \\citep{YSW77,Young77b,McMLC81,DS83} or on a semi-analytical recipe proposed by \\citet[hereafter SS66]{SS66} that completely neglects the real hydrodynamical nature of the process \\citep{Sanders70b,DDC87a,DDC87b,MCD91}. The work of \\citet{Rauch99} is a noticeable exception; he used the results of a set of hydrodynamical simulations of stellar collisions by Davies to derive fitting formulae for the quantitative outcome of these events.  The present work originated in our wish to get rid of these annoying uncertainties about the role of collisions in dynamical simulations of galactic nuclei \\citep{FB01d,FB02b}.  Many of the papers we have just cited were not only concerned with the past evolution of galactic nuclei but also (or mainly) with scenarios to feed SMBH and provide quasars' luminosities. Gas-dynamical processes are now favoured candidates for the fuelling of active galactic nuclei (AGN) and the dense cluster hypothesis seems somewhat out-of-fashion \\citep[and references therein]{SBF90,Combes01}. On the other hand, AGN models have been proposed in which large luminosities in electro-magnetic radiation and/or relativistic particles are emitted by the hot gas clouds created by very energetic stellar collisions.  First propositions along that line \\citep{Woltjer64,Sanders70a} postulated that the stars' velocities were due to the cluster's self gravity.  More recent models \\citep{Keenan78,DKT93,CPW96,TCFCP00} invoke a SMBH to provide velocity dispersions ranging from a few $10^3$\\,km\\,s$^{-1}$ to a few $10^4$\\,km\\,s$^{-1}$. These non-standard AGN models may be successful in reproducing observed luminosity-variability relations that are otherwise difficult to explain, but they should be re-examined in the light of a more refined treatment of stellar collisions and stellar dynamics. A third possibility for collisions to contribute directly to the luminosity of AGN is to boost the rate of supernovae through creation of massive stars by mergers \\citep{Colgate67,SW78}.  Finally, even though they are not the dominating luminosity source in AGNs, stellar collisions may be responsible for the formation of massive black holes in dense galactic nuclei, either by run-away merging (\\citealt{Sanders70b,QS90,PZMcM02}; \\citealt*{RFG03,GFR04,FGR04b,FGR04c,FGR05}) or by creating a massive gas cloud that subsequently evolves to a black hole \\citep{SS66,BR78,LFSY90}.   \\subsection{Previous simulations of collisions between main sequence stars.} \\label{subsec:prev_works}  { \\begin{table*} \\caption{Hydrodynamical simulations of high-velocity collisions between MS stars in the literature.} \\setlength\\tabcolsep{3pt} \\begin{center} \\begin{tabular}{lclllll} \\hline\\hline Reference&Abbrev.&Stellar models&$q=M_1/M_2$ & $V_{\\mathrm{rel}}^{\\infty}/V_{\\ast}$$^{\\mathrm{(a)}}$ &$M$--$R$ relation&Method\\\\ \\hline \\protect\\citet{SC72} && polytropes $n=3$ & 1 & 0, 1.6, 3.2 & & Head-on, 2D finite diff. \\\\ \\protect\\citet{BH87}  &BH87& polytropes $n=1.5$ & 1 & 0--2.33 & & SPH 1000 part. \\\\ \\protect\\citet{BH92}  &BH92& polytropes $n=1.5$ & 0.2 & 0--1.5$^{\\mathrm{(b)}}$ & $R_{\\ast}\\propto M_{\\ast}^{0.85}$ & SPH 7000 part. \\\\ \\protect\\citet{LRS93}$^{\\mathrm{(c)}}$ &LRS93& polytropes $n=3$$^{\\mathrm{(d)}}$ & 0.1--1 & 0.2--3.8 & $R_{\\ast}\\propto M_{\\ast}^{0.8}$ & SPH 8000 part. \\\\ Davies \\protect\\citep{Rauch99}$^{\\mathrm{(e)}}$ &R99&  polytropes $n=3$ & 0.25, 0.5, 1 & 1--25 & $R_{\\ast}\\propto M_{\\ast}$ & SPH $\\sim 40\\,000$ part.\\\\ This work && realistic & 0.0013--1 & 0.1--30 & realistic & SPH 4000--36\\,000 part. \\\\ \\hline \\multicolumn{6}{l}{ $^{\\mathrm{(a)}}$See symbols definition in Sec.~\\protect\\ref{subsec:def}. }\\\\ \\multicolumn{6}{l}{ $^{\\mathrm{(b)}}$Up to 5 for head-on collisions. }\\\\ \\multicolumn{6}{l}{ $^{\\mathrm{(c)}}$Fitting formulae are given. }\\\\ \\multicolumn{6}{l}{ $^{\\mathrm{(d)}}$Eddington models. }\\\\ \\multicolumn{6}{l}{ $^{\\mathrm{(e)}}$Results only given as fitting formulae. }\\\\ \\end{tabular} \\end{center} \\label{tbl:coll_simul} \\end{table*} }  Table~\\ref{tbl:coll_simul} lists the previous computations of high-velocity collisions between MS stars. We only mention ``realistic'', multi-dimensional hydrodynamical simulations. This excludes early calculations that were based either on semi-analytical methods \\citepalias{SS66} or on one-dimensional numerical schemes \\citep{Mathis67,DeYoung68}. Such approaches were clearly over-simplifications in which the real 3D hydrodynamical nature of the problem was not properly accounted for. The importance of these ``pre-hydrodynamics'' works should not be underestimated, however. For instance, even though it was always deemed too simplistic to yield better than order-of-magnitude estimates, the \\citetalias{SS66} method had been adapted and used in a few key simulation works. We postpone a presentation of this ``historical'' method to Sec.~\\ref{subsec:comp_lit} where we compare our results to predictions of this approach.  With the historical exception of \\citet{SC72}, all cited works were realised using SPH (Smoothed Particle Hydrodynamics). When featured with a tree to compute gravitation, SPH is a grid-less method which can cope with any asymmetrical three dimensional geometry.  It ignores void spaces completely, imposes no physical limits beyond which matter cannot be tracked, does not come into trouble with large dynamic range as long as variable smoothing lengths are implemented.  SPH is better suited to highly dynamical problems than to near-equilibrium configurations \\citep{SM93}. For all these reasons, SPH is particularly well suited to the simulation of stellar collisions.  From Table~\\ref{tbl:coll_simul}, it is clear that the study of the outcome of high-velocity collisions has not attracted much interest in the last few years, in contrast to parabolic encounters \\citep[][among others]{LRS96,SL97,SFLRW00,SADB02}. As a consequence, the resolutions used seem very modest, by present-day standards; for instance, \\citet{SADB02} present a parabolic collision simulated with $10^6$ particles. Obviously, the simulations presented in this work do not correspond to a break-through in terms of resolution. This reflects the fact that most computations were realised a few years ago, when computing power was more limited and, most importantly, that we had to cover a huge parameter-space, requiring more than 10\\,000 simulations (see Sec.~\\ref{subsec:buildingtable}). This sheer quantity, combined with the use of realistic stellar models instead of polytropes, represent the main improvements over previous works.  \\defcitealias{BH87}{BH87} \\defcitealias{BH92}{BH92} \\defcitealias{LRS93}{LRS93} \\defcitealias{Rauch99}{R99}  For simplicity, in the remaining of this article, we refer to the work of \\citet{BH87} as \\citetalias{BH87}, to \\citet{BH92} as BH92, to \\citet*{LRS93} as \\citetalias{LRS93} and to \\citet{Rauch99} as \\citetalias{Rauch99}. For a more comprehensive list of references on simulations of all kind of stellar collisions, see the web site maintained by MF in the framework of the ``MODEST'' collaboration\\footnote{``MODEST'' stands for Modelling DEnse STellar systems, see \\url{http://www.manybody.org/modest/}. For the collision ``working group'', go to \\url{http://www.manybody.org/modest/WG/wg4.html}.}.   \\subsection{Collisions with non-main-sequence stars}  \\begin{figure*} \\resizebox{\\hsize}{!}{% \\includegraphics{Pcoll_vs_Age.eps} } \\caption{Cumulative collision probability (normalised to 1) integrated over the lifetime of three stellar models. The second star is assumed to be sun-like. In each case, three velocity regimes are considered. At low relative velocities ($V_{\\mathrm{rel}}^\\infty\\ll V_{\\ast}$), the collisional cross-section scales like $R_{\\ast}+\\Rsun$ (solid lines); at very high velocities ($V_{\\mathrm{rel}}^\\infty\\gg V_{\\ast}$), we recover the geometrical cross section, $\\propto (R_{\\ast}+\\Rsun)^2$ (dashed lines). We also plot the case for a relative velocity of 200\\,{\\kms}, an intermediate value typical for galactic nuclei. The end of the MS phase corresponds to the point where the slope of the curves increases suddenly. Stellar evolution models are from the compilation of \\citet{LS01}, available on-line at \\url{http://vizier.cfa.harvard.edu/viz-bin/VizieR?-source=VI/102}.} \\label{fig:pcoll_vs_age} \\end{figure*}  In this work, we only treat collisions between two main-sequence (MS) stars. The motivations for this choice was to keep the number and variety of collisions to consider at a manageable level and that the present version of our Monte Carlo code only includes simplified stellar evolution which skips over the giant phase and turns MS stars directly into remnants. However, in a real stellar system, MS-MS encounters may not dominate the global collision rate. Indeed, a given star of mass $> 1\\,{\\Msun}$ may have a smaller probability for colliding with another star during its MS life than during its red giant (RG) phase despite the latter being about 10 times shorter \\cite[for instance]{BFBC93}. This is made very clear by integrating the collisional cross-section over the lifetime of the star, as we did in Fig.~\\ref{fig:pcoll_vs_age}. In many cases, the collision probability during the red giant phase exceeds its MS counterpart for high relative velocities. RG-RG collisions are less likely than RG-MS events. Indeed, the ratio of probabilities can be estimated as follows \\[ \\frac{P_{\\mathrm{RG-MS}}}{P_{\\mathrm{RG-RG}}} \\sim \\frac{n_{\\mathrm{MS}}}{n_{\\mathrm{RG}}} \\cdot \\frac{R_{\\mathrm{RG}}^2}{(2R_{\\mathrm{RG}})^2} \\simeq 0.25 \\frac{T_{\\mathrm{MS}}}{T_{\\mathrm{RG}}} \\simeq 3-10. \\] \\Added{Although probably more common than MS-MS encounters, RG-MS collisions may not be more important. RG envelopes have very low densities so only little mass is lost in most cases and the RG recovers its appearance. At relative velocities found in galactic nuclei, the MS star cannot be captured unless it is aimed nearly directly at the RG centre \\citep{BD99}. Furthermore, as giants will loose their envelope anyway through winds and a planetary nebula or SN phase, collision with giants will probably make little difference as far as the feeding of a central SMBH is concerned.  Due to mass segregation in clusters and nuclei, collisions between compact remnants (CRs) and MS (or RG) stars are probably much more common than the small CR fractional number would suggest. For instance, the innermost 0.1\\,pc of the {\\SgrA} cluster is likely dominated by invisible stellar BHs \\citep{MEG00} which may collisionally destroy MS and RG stars. CR-MS and CR-RG collisions may also be of great interest as a way to produce exotic objects, such as cataclysmic variables. Unfortunately, due to the high dynamical range involved, the hydrodynamical simulation of these events is challenging and comprehensive predictions for their outcome are still lacking.}  Now that the astrophysical stage is set, we can proceed with a description of our simulation work. In Sec.~\\ref{sec:approach}, we describe the choice and setting of initial conditions and present the numerical methods we use to compute and analyse stellar collisions. In Sec~\\ref{sec:results}, results are reported and we explain how to exploit them in stellar dynamical simulations.  Finally, in Sec~\\ref{sec:concl}, some general conclusions and a discussion of further work to be done are presented. ",
        "conclusions": "\\label{sec:concl}  In this article, we presented a large set of simulations of collisions between two main sequence stars. More than 14\\,000 SPH simulations have been computed over about four years to complete that database. Initial conditions span $M_\\ast=0.1-75\\,\\Msun$ for the stellar masses, impact parameters corresponding approximately to $d_{\\mathrm{min}}/(R_1+R_2) =0-0.9$ and relative velocities at infinity ranging, more or less, from 0.03 to 30 times the stellar escape velocity. This represents a effort of unprecedented breadth in this field.  Our motivation in this work was to incorporate the effects of stellar collisions into models of dense stellar systems like galactic nuclei with as much realism as possible. To reach this goal we developed a module that interpolates between our results to predict the outcome of any collision with initial conditions inside the (large) domain of parameter space we explored. Results of our dynamical simulations of dense clusters including collisions are presented elsewhere (\\citealt{FB01d,FB02b,RFG03,FGR04b,FGR04c,FRB05,FGR05}).  The quest for a handy mathematical description of these results has been unsuccessful so far. This was a source of disappointment but we hope that further study of our simulation data will eventually reveal some way of casting our results in a compact and physically enlightening formulation. Exploring when and why the excellent match with the SS66 prescription is found is a possible way to this goal.  Beyond the scientific exploitation of this important database, --either to develop a better understanding of the physics of collisions or as an ingredient for collisional stellar dynamics studies-- we are aware that other kinds of stellar collisions, not treated here, are also of great astrophysical interest. First, in galactic nuclei, collisions with red giants are certainly more frequent than MS-MS encounters. So we should assess the importance of this process (for stellar evolution and stellar dynamics) and, if needed, compute a set of simulations similar to the one presented here. This would complement the work of \\citet{BD99}. Collisions with compact remnants should also be taken into account as they may be of great importance in producing peculiar objects.  Our work is not well adapted to the physical conditions that are typical in globular clusters. Indeed, we did not study with particular care the low velocity, quasi parabolic encounters, by far the most common in those environments. However, they have already been quite thoroughly studied by other researchers. The conditions for tidal binary formation are probably better determined using other methods and their long term evolution, whose nature and result is still debated \\citep{Mardling95a,Mardling95b}, extends over much too many hydrodynamical time scales to be tackled by SPH. Still in the domain of globular clusters, the study of hydrodynamical effects (including direct collisions) in interactions with or between binary stars is still in its infancy \\citep{GH91,DBH93,DBH94}.  Finally, going back to galactic nuclei in which we are mostly interested, let's mention that the connected problem of the destructive close encounter between a central massive black hole and a star, even though it has been the focus of many papers \\cite[and many others]{CL83,BG83,EK89,LMZD93,Fulbright96,MLB96,ALP00,BogdanovicEtAl04}, has still not been explored systematically. For instance, all studies so far have assumed simplified stellar models. In this problem, however the main uncertainties probably lurk in the post-disruption evolution of the stellar gas, rather than in the ``collision'' process itself.  \\appendix"
    },
    "0403/astro-ph0403417_arXiv.txt": {
        "abstract": "The annihilation of new dark matter candidates with masses $m_X$ in the MeV range may account for the galactic positrons that are required to explain the 511~keV $\\gamma$-ray flux from the galactic bulge. We study the impact of MeV-mass thermal relic particles on the primordial synthesis of $^2$H, $^4$He, and $^7$Li. If the new particles are in thermal equilibrium with neutrinos during the nucleosynthesis epoch they increase the helium mass fraction for $m_X\\alt 10$~MeV and are thus disfavored.  If they couple primarily to the electromagnetic plasma they can have the opposite effect of lowering both helium and deuterium.  For $m_X=4$--10~MeV they can even improve the overall agreement between the predicted and observed $^2$H and $^4$He abundances. ",
        "introduction": "In a recent series of papers the intriguing possibility was explored that the cosmic dark matter consists of new elementary particles with masses in the MeV range~\\cite{Boehm:2002yz, Boehm:2003hm, Boehm:2003bt, Hooper:2003sh, Boehm:2003ha}.  While weakly interacting massive particles (WIMPs) as dark matter candidates, notably supersymmetric particles, are usually thought to have masses exceeding tens of GeV, it is not difficult to come up with viable MeV-mass candidates such as sterile neutrinos~\\cite{Abazajian:2001nj} or axinos~\\cite{Covi:2002vw}. The remarkable aspect of the MeV candidates studied in Refs.~\\cite{Boehm:2002yz, Boehm:2003hm, Boehm:2003bt, Hooper:2003sh, Boehm:2003ha} is that they are taken to be {\\em thermal\\/} relics and thus require a primordial annihilation rate much larger than given by ordinary weak interactions. Contrary to naive intuition, such particles are not excluded by any obvious laboratory measurement or astrophysical argument~\\cite{Boehm:2002yz, Boehm:2003hm, Boehm:2003bt, Hooper:2003sh, Boehm:2003ha, endnote1}. Quite on the contrary, it was argued that the annihilation of these dark matter particles in the galactic bulge can produce enough positrons to explain the 511~keV $\\gamma$-ray signature that was recently confirmed by the INTEGRAL satellite~\\cite{Jean:2003ci,Knodlseder:2003sv} and that seems difficult to explain with traditional astrophysical sources. Alternatively, decaying low-mass dark matter particles have also been proposed as a positron source~\\cite{Picciotto:2004rp,Hooper:2004qf}.  In the early universe, weak interactions freeze out at a temperature of about 1~MeV, just before the epoch of big-bang nucleosynthesis (BBN). Therefore, MeV-mass particles with larger-than-weak interaction rates are expected to be in thermal equilibrium at BBN and thus would add to the primordial energy density and expansion rate. One naively expects such particles to increase the primordial helium mass fraction $Y_{\\rm p}$, exacerbating the tension between observations and the BBN prediction. This expectation bears out for new particles coupling primarily to neutrinos.  However, the new particles proposed in Refs.~\\cite{Boehm:2002yz, Boehm:2003hm, Boehm:2003bt, Hooper:2003sh, Boehm:2003ha} would primarily interact with the electromagnetic plasma. It turns out that the BBN effects of these particles are non-trivial and certainly disfavor masses below about 2~MeV, but masses in the approximate range 4--10~MeV actually have the opposite effect of {\\em lowering\\/} $Y_{\\rm p}$ without significantly affecting deuterium. Therefore, the concordance between the predicted and observed $^4$He and $^2$H abundances is slightly improved.  This little-known but intriguing effect was found a long time ago in a general study of the impact of new particles on BBN~\\cite{Kolb:1986nf}. The main purpose of our work is to re-examine this effect in the context of the MeV-mass dark matter hypothesis and with the help of a modern BBN calculation and current observational data.  In addition to our BBN study presented in Sec.~II we briefly consider two other possible consequences of MeV-mass dark matter particles. In Sec.~III we note that the same mechanism proposed to explain the $\\gamma$-ray signature from the galactic bulge should also produce a diffuse background of low-energy cosmic-ray positrons in the solar neighborhood.  This argument strengthens the conclusion reached independently in Refs.~\\cite{Boehm:2002yz, Boehm:2003hm, Boehm:2003bt, Hooper:2003sh, Boehm:2003ha} that an s-wave annihilation cross section for the reaction $X\\bar X\\to e^+e^-$ is strongly disfavored.  Further, in Sec.~IV we consider possible energy-transfer effects caused by such particles trapped in the Sun or other stars. These effects turn out to be very small because the trapping efficiency is inhibited by the smallness of $m_X$. In Sec.~V we summarize our conclusions. ",
        "conclusions": "We have analyzed some astrophysical and cosmological consequences of the intriguing possibility that the cosmic dark matter consists of MeV-mass particles~\\cite{Boehm:2002yz, Boehm:2003hm, Boehm:2003bt, Hooper:2003sh, Boehm:2003ha}.  These particles are assumed to be thermal relics and thus have interaction cross sections that are larger than weak.  Such particles do not have any apparent consequences for stellar evolution. In supernovae, their interaction is too strong so that neutrinos continue to play the leading role for energy loss and energy transfer.  In the Sun, the particle mass is so small that evaporation prevents the trapped steady-state population from growing large enough for a significant contribution to energy conduction.  We have derived new constraints coming from the low-energy positron component of cosmic rays. An s-wave annihilation cross section for $X\\bar X\\to e^+e^-$ as large as implied by the early-universe freeze-out calculation causes an excessive positron flux for $m_X=20$--90~MeV where experimental upper limits are available. Therefore, only p-wave annihilation is compatible with these constraints. These conclusions agree with those reached in Refs.~\\cite{Boehm:2002yz, Boehm:2003hm, Boehm:2003bt, Hooper:2003sh, Boehm:2003ha}, but in our case they do not depend on the assumed dark-matter profile of the galactic bulge.  The main subject of our work, however, was the impact of MeV-mass particles on big-bang nucleosynthesis (BBN). Significant modifications arise only for $m_X\\alt20$~MeV.  The effects found are largely independent of the energy dependence of the annihilation cross section and even of its exact value, provided that it is high enough to take the new particles to equilibrium with the neutrinos or the electromagnetic plasma.  For the neutrino-coupled case, the impact of the new particles is comparable to that of additional neutrino species.  Notably, the primordial helium abundance is increased, exacerbating the discrepancy between predictions and observations.  Therefore, such $X$-particles are disfavored by BBN for masses up to about 10~MeV.  For the electromagnetically coupled case, the BBN concordance would be severely disturbed for $m_X\\alt2$~MeV.  However, there is a region $m_X=4$--10~MeV where the primordial helium mass fraction $Y_{\\rm p}$ is actually reduced relative to the standard case while the predicted $^2$H remains compatible with observations.  This non-trivial phenomenon is a consequence of the ``entropy effect'' discussed in the paper.  While this effect was already found in Ref.~\\cite{Kolb:1986nf}, it now assumes greater importance because it slightly improves the discrepancy between the BBN predictions for $Y_{\\rm p}$ and its observed value.  In summary, the MeV-mass dark matter particles proposed in Refs.~\\cite{Boehm:2002yz, Boehm:2003hm, Boehm:2003bt, Hooper:2003sh, Boehm:2003ha} as a source for positrons in the galactic bulge are not incompatible with BBN, provided the particle mass exceeds a few MeV. On the contrary, in the mass range $m_X=4$--10~MeV these particles slightly improve the concordance between BBN calculations and the observed helium abundance. It is quite fascinating that such exotic particles, far from being excluded, seem to have several beneficial consequences in astrophysics and cosmology."
    },
    "0403/hep-th0403020_arXiv.txt": {
        "abstract": "Inflation has been the leading early universe scenario for two decades, and has become an accepted element of the successful `cosmic concordance' model. However, there are many puzzling features of the resulting theory. It requires both high energy and low energy inflation, with energy densities differing by a hundred orders of magnitude. The questions of why the universe started out undergoing high energy inflation, and why it will end up in low energy inflation, are unanswered. Rather than resort to anthropic arguments, we have developed an alternative cosmology, the cyclic universe\\cite{cyclic}, in which the universe exists in a very long-lived attractor state determined by the laws of physics. The model shares inflation's phenomenological successes {\\it without} requiring an epoch of high energy inflation. Instead, the universe is made homogeneous and flat, and scale-invariant adiabatic perturbations are generated during an epoch of low energy acceleration like that seen today, but preceding the last big bang. Unlike inflation, the model {\\it requires} low energy acceleration in order for a periodic attractor state to exist. The key challenge facing the scenario is that of passing through the cosmic singularity at $t=0$. Substantial progress has been made at the level of linearised gravity, which is reviewed here. The challenge of extending this to nonlinear gravity and string theory remains. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403692_arXiv.txt": {
        "abstract": "Supplementing existing survey data with VLA observations, we have extended $\\gamma-$ray counterpart identifications down to decl. = -40$^\\circ$ using our Figure of Merit approach.  We find blazar counterparts for $\\sim$ 70\\% of EGRET sources above decl. = -40$^\\circ$ away from the Galaxy. Spectroscopic confirmation is in progress, and spectra for $\\sim$ two dozen sources are presented here. We find evidence that increased exposure in the bulge region allowed EGRET to detect relatively faint blazars; a clear excess of non-blazar objects in this region however argues for an additional (new) source class. ",
        "introduction": "The \\EGRET telescope on the Compton Gamma Ray Observatory (CGRO) satellite has detected 271 sources in a survey of the $\\gamma$-ray (~100 MeV to 10 GeV) sky.  Roughly one quarter of these have previously been identified as blazars \\citep{har99,mat01}.  Recently, we have developed a new counterpart identification technique, which has pushed the identified fraction to $\\sim$70\\% in the northern hemisphere (Sowards-Emmerd \\et 2003, hereafter SRM03).  Here, we extend the application of this technique south (-40$^\\circ$ $<$ decl. $<$ 0$^\\circ$), including the Galactic bulge region.  In order to rank blazar candidates in this zone using our `Figure of Merit' (FoM) method, we have selected flat spectrum radio counterpart candidates in 3EG source error boxes. We have obtained spectroscopic confirmations and estimated redshifts for many of these candidate blazars with the Hobby-Eberly Telescope (HET) Marcario Low Resolution Spectrograph (LRS) and with the  2.7-m Harlan J. Smith Telescope IGI spectrograph at McDonald Observatory. For our final counterpart list, we used the VLA to obtain compact 8.4 GHz fluxes similar to those previously mined from the CLASS (Meyers \\et 2002) survey.  This allows a robust selection of EGRET blazar-like counterparts and a statistical measure of the identification probability of each source.  The blazar label is somewhat heterogeneous, but in the context of the unified AGN model, these sources are believed to be viewed close to the axis of a powerful relativistic jet. As such they are compact flat spectrum radio sources, with apparent superluminal motion at VLBI scales. The optical counterparts exhibit significant polarization and OVV (optically violently variable) behavior \\citep{up95}. Optical spectroscopy yields a dichotomy of sources:  flat spectrum radio quasars with broad emission lines and continuum-dominated BL Lac-type objects showing weak absorption features and occasionally narrow, weak emission lines.  The broad-band spectral energy distribution (SED) is sometimes used to divide these into two classes, with `red' blazars showing a synchrotron peak in the IR-optical with a Synchrotron Self-Compton (SSC) peak in the $\\gamma$-ray while `blue' blazars have a synchrotron component extending into the X-ray with a SSC peak inferred to extend to the TeV range (Urry 1999, and references therein). ",
        "conclusions": "From our radio follow-up, our FoM method has selected 23 new IDs, with 8.4 GHz fluxes, between -40$^\\circ$ $<$ decl. $<$ 0$^\\circ$.  Of these new identifications, seven sources have high confidence `likely' counterparts and the remainder are plausible.  We expect no more than 3 of these new sources are false positives.  A few additional identifications are based on 4.85 GHz extrapolations and require 8.4~GHz compact fluxes for confirmation. We also have spectroscopically identified 26 new optical counterparts selected for this sample.  Several of these targets fell below our FoM after VLA follow-up - the spectroscopic targets were selected before 8.4~GHz fluxes were available - but are included in Table 2 for completeness. Combined with the Northern sample, for high latitude 3EG error boxes ($\\vert b \\vert>10^\\circ$ and $>$30$^\\circ$ from the galactic center), we find blazar IDs for 88 of 129 sources, or 68\\%. Our analysis has now been applied to 157 of the 186 3EG detections located above $\\vert b \\vert>10^\\circ$.  Extending the survey into the Galactic bulge, we find a surplus of unidentified 3EG sources. These likely comprise a new set of Galactic $\\gamma$-ray emitters.  Since the areal density of identified blazars in the bulge region is also larger than in the North, it seems that the longer EGRET exposure in this region has let us look lower on the blazar luminosity function, which should be helpful in refining the blazar contribution to the unresolved high latitude $\\gamma$-ray counts.  We defer detailed calculation of the blazar luminosity function and number count predictions for GLAST to a later paper."
    },
    "0403/astro-ph0403371_arXiv.txt": {
        "abstract": "A method is presented to derive electron temperatures and densities of planetary nebulae (PNe) simultaneously, using the observed hydrogen recombination spectrum, which includes continuum and line emission. By matching theoretical spectra to observed spectra around the Balmer jump at about 3646~{\\AA}, we determine electron temperatures and densities for 48 Galactic PNe. The electron temperatures based on this method [hereafter $T_{\\rm e}$(Bal)] are found to be systematically lower than those derived from [\\ion{O}{3}] $\\lambda4959/\\lambda4363$ and [\\ion{O}{3}] $(88\\mu{\\rm m}+52\\mu{\\rm m})/\\lambda4959$ ratios [hereafter, $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$) and $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm fn}$)]. And the electron densities based on this method are found to be systematically higher than those derived from [\\ion{O}{2}] $\\lambda3729/\\lambda3726$, [\\ion{S}{2}] $\\lambda6731/\\lambda6716$, [\\ion{Cl}{3}] $\\lambda5537/\\lambda5517$, [\\ion{Ar}{4}] $\\lambda4740/\\lambda4711$ and [\\ion{O}{3}] $88\\mu{\\rm m}/52\\mu{\\rm m}$ ratios. These results suggest that temperature and density fluctuations are generally present within nebulae. The comparison of  $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$) and $T_{\\rm e}$(Bal) suggests that the fractional mean-square temperature variation ($t^2$) has a representative value of 0.031. A majority of temperatures derived from the [\\ion{O}{3}] $(88\\mu{\\rm m}+52\\mu{\\rm m})/\\lambda4959$ ratio are found to be higher than $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$), which is attributed to the existence of dense clumps in nebulae -- those [\\ion{O}{3}] IR fine-structure lines are suppressed by collisional de-excitation in the clumps. By comparing $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm fn}$), $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$) and $T_{\\rm e}$(Bal) and assuming a simple two-density-component model, we find that the filling factor of dense clumps has a representative value of 7$\\times10^{-5}$. The discrepancies between $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$) and $T_{\\rm e}$(Bal) are found to be anti-correlated with electron densities derived from various density indicators; high-density nebulae have the smallest temperature discrepancies. This suggests that temperature discrepancy is related to nebular evolution. In addition, He/H abundances of PNe are found to be positively correlated with the difference between $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$) and $T_{\\rm e}$(Bal), suggesting that He/H abundances might have been overestimated generally because of the possible existence of H-deficient knots. Electron temperatures and densities deduced from spectra around the Paschen jump regions at 8250~{\\AA} are also obtained for four PNe, NGC~7027, NGC~6153, M~1-42 and NGC~7009. Electron densities derived from spectra around the Paschen jump regions are in good agreement with the corresponding values derived from spectra around the Balmer jump, whereas temperatures deduced from the spectra around the Paschen jump are found to be lower than the corresponding values derived from spectra around the Balmer jump for all the four cases. The reason remains unclear. ",
        "introduction": "Accurate determinations of metal abundances in planetary nebulae (PNe) are important for understand the chemical evolution of galaxies and the nucleosynthesis and mixing processes in low- and intermediate-mass stars. A long-standing problem in nebular abundance studies has been that the heavy-element abundances derived from optical recombination lines (ORLs) are systematically higher than those derived from collisionally excited lines (CELs; see Liu 2001, 2003 for recent reviews). In extreme cases, the discrepancies exceed a factor of 10. Given that CEL abundances are sensitive to electron temperature and density, knowledge of temperature and density variations is an essential ingredient in understanding the abundance discrepancy problem. Information on temperature and density variations can be obtained by comparing results derived from plasma diagnostics which behave differently as temperature and density vary. Traditionally, nebular electron temperatures and densities are derived from CEL ratios. However, for PNe including high-density regions ($N$e$\\ga10^6$\\,cm$^{-3}$), such as IC~4997 \\citep{hyung1994}, Mz~3 \\citep{zhang02} and M~2-24 \\citep{zhang2003}, CELs are strongly suppressed by collisional de-excitation, and so depending on their critical densities, they may become unusable in probing such high-density regions. On the other hand, hydrogen recombination lines have very high critical densities ($\\ga10^8$\\,cm$^{-3}$). Consequently, the hydrogen recombination spectrum provides a powerful tool to probe high-density regions. Comparison of electron temperatures and densities derived from the hydrogen recombination spectrum (lines plus continuum) with those derived from CELs should enable us to quantify temperature and density variations in nebulae.  By measuring the Balmer discontinuity, \\citet{liu93} determined electron temperatures in 14 PNe and found that they are systematically lower than those derived from the [\\ion{O}{3}] nebular-to-auroral forbidden line ratio. This can be explained by the presence of large temperature fluctuations since the [\\ion{O}{3}] $\\lambda4363$ auroral line of higher excitation energy tends to be emitted in higher temperature regions than the [\\ion{O}{3}] $\\lambda\\lambda4959,5007$ nebular lines of lower excitation energy, as suggested by \\citet{peim67}. Recent reviews on temperature fluctuations of gaseous nebulae are given by \\citet{peim03} and references therein. However, no photoionization models yield such large temperature fluctuations suggested by observations. Moreover, temperature fluctuations are unable to explain the large disparity between the ORL and CEL abundances observed in some PNe \\citep{liua}. Instead, it is found that the hypothesis that PNe may contain clumps of different temperature, density and chemical composition, embedded in diffuse `normal' material, can explain many of the observed patterns \\citep{liubarlow2000}. Peimbert's concept of temperature fluctuations alone may not be a realistic description of the physical conditions in a nebula, although it has proven to be a convenient and useful way to characterize the complexities of real nebulae.  It is expected that there are also variations of electron density inside a PN.  \\citet{rubin1989} pointed out that electron densities derived from various plasma diagnostics should correlate positively with the critical densities of the diagnostic lines involved if there are large density variations within a nebula. \\citet{liubarlow2001} found that electron densities derived from the [\\ion{O}{3}] $88\\mu{\\rm m}/52\\mu{\\rm m}$ ratios are generally lower than those derived from the optical [\\ion{Ar}{4}] and [\\ion{Cl}{3}] doublet ratios, suggesting the presence of high density regions in PNe, where the 52- and 88-$\\mu{\\rm m}$ lines are suppressed by collisional de-excitation due to their relatively low critical densities. \\citet{viegas} showed that the discrepancies between $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$) and $T_{\\rm e}$(Bal) can be explained by the presence of condensations in nebulae and gave an analytical method of estimating the filling factor of condensations. Our detailed studies of the peculiar PNe Mz~3 and M~2-24 suggest that the presence of high-density regions in some PNe can significantly affect abundance determinations \\citep{zhang02,zhang2003}. However, density inhomogeneities alone cannot explain the large disparity between the ORL and CEL abundances either, unless the condensations are also H-deficient \\citep{liubarlow2000}.  \\citet{liubarlow2000} found that the discrepancy between the ORL and CEL abundances correlates with the difference between the temperatures derived from the [\\ion{O}{3}] collisional lines and from the Balmer jump of hydrogen recombination spectrum, thus lending strong support to the argument that some extremely cold H-deficient knots may co-exist with normal nebular material in many PNe. However, it is difficult to explain the origin and evolution of such H-deficient knots with the current theories of element production in intermediate-mass stars.  In this paper, we present determinations of electron temperatures and densities for a large sample of PNe based on the analysis of nebular hydrogen recombination spectrum. Section 2 describes the sample. In Section 3, we describe a new approach to deriving electron temperature and density simultaneously from the observed nebular hydrogen recombination spectrum (lines plus continuum) and present the results. In Section 4, we compare results derived from a variety of temperature and density indicators and discuss the significance of variations in electron temperature and density in PNe. A discussion of He/H abundance determinations is also presented in this Section. In Section 5, we present electron temperatures and densities derived from the Paschen discontinuity and decrements. A summary then follows in Section 6. ",
        "conclusions": "\\subsection{Comparison of densities}  Electron densities deduced from the Balmer decrements, and from the [\\ion{O}{2}] $\\lambda3729/\\lambda3726$, [\\ion{S}{2}] $\\lambda6716/\\lambda6731$, [\\ion{Cl}{3}] $\\lambda5517/\\lambda5537$ and [\\ion{Ar}{4}] $\\lambda4711/\\lambda4740$ doublet ratios are compared in Fig.~\\ref{necorr}. Fig.~\\ref{necorr} shows that for most PNe, the hydrogen recombination spectrum yields higher densities compared to forbidden lines diagnostics. There is a trend such that the discrepancy becomes larger at higher densities. Two extreme cases are IC~4997 and M~2-24. For IC~3568, [\\ion{Cl}{3}] $\\lambda5517/\\lambda5537$ yields an abnormally low density. The cause of it is unclear but is probably due to measurement error. In Fig.~\\ref{necorrir} we compare electron densities derived from the [\\ion{O}{3}] far-infrared fine structure line ratio and from the hydrogen recombination spectrum. Evidently, densities derived from the [\\ion{O}{3}] $52\\mu{\\rm m}/88\\mu{\\rm m}$ ratio are significantly lower than those derived from the Balmer decrements. In fact, it has been found that the far-IR lines yield systematic lower densities, lower than even those deduced from optical forbidden line density-diagnostics \\citep{liubarlow2001}. A linear fit shows that the density discrepancy increases with the electron density.  O$^0$, S$^0$, Cl$^{+}$, O$^{+}$ and Ar$^{2+}$ have ionization potentials of 13.6, 10.4, 23.8, 35.1 and 40.7~eV, respectively, and therefore embrace a wide range of regions of differing states of ionization. On the other hand, the hydrogen recombination spectrum arises from the entire ionized nebula. Thus the discrepancies between densities derived from \\ion{H}{1} recombination spectrum and those deduced from forbidden line ratios are not likely to be caused by ionization stratifications. We notice that the [\\ion{O}{2}], [\\ion{S}{2}], [\\ion{Cl}{3}], [\\ion{Ar}{4}] and [\\ion{O}{3}] density-diagnostic lines have critical densities of $N_{\\rm crit}\\la10^5$~cm$^{-3}$. Therefore, if the nebulae contain regions of higher densities, these forbidden lines will be suppressed in these regions by collisional de-excitation. In contrast, the hydrogen recombination spectrum is practically unaffected by collisional de-excitation. As a result, in such cases the hydrogen recombination spectrum will yield higher average densities compared to the [\\ion{O}{2}], [\\ion{S}{2}], [\\ion{Cl}{3}], [\\ion{Ar}{4}] and [\\ion{O}{3}] lines ratios. The same argument has been developed by \\citep{liubarlow2001} to explain the systematically lower densities derived from the [\\ion{O}{3}] fine-structure line ratio compare to those deduced from optical forbidden line ratios. The [\\ion{Ar}{4}] lines have higher critical densities than the [\\ion{O}{2}], [\\ion{S}{2}], [\\ion{Cl}{3}] lines. Thus, the effects are weaker for [\\ion{Ar}{4}] doublet ratio than for the other cases. Given that the [\\ion{O}{3}] far-infrared fine-structure lines have the lowest critical densities amongst those density-diagnostic forbidden lines ($N_{\\rm cri}$([\\ion{O}{3}]$_{\\rm ff})\\la4\\times10^3$~cm$^{\\rm -3}$), the effects are largest for the [\\ion{O}{3}] $52\\mu{\\rm m}/88\\mu{\\rm m}$ ratio. This can be seen clearly in Fig.~\\ref{necorr} and Fig.~\\ref{necorrir}. A nebular model including these high-density condensations and their effects on temperature determinations will be discussed in Section 4.3.  \\subsection{Comparison of temperatures} Fig.~\\ref{tecorr} and Fig.~\\ref{tecorrir} show that electron temperatures derived from the [\\ion{O}{3}] $\\lambda4959/\\lambda4363$ and $(52\\mu{\\rm m}+88\\mu{\\rm m})/\\lambda4959$ ratios [hereafter, $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na})$ and $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm fn})$] are systematically higher than those deduced from the hydrogen recombination spectrum [hereafter $T_{\\rm e}$(Bal)], although there are a few exceptions where $T_{\\rm e}$(Bal) is higher than $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$). The most extreme case is He~2-118, for which $T_{\\rm e}$(Bal) is 5000~K higher than $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$). In our sample, the four PNe with the largest temperatures discrepancies, Hf~2-2, M~1-42, M~2-36 and NGC~6153, were previously studied in detail by \\citet{liu209}, \\citet{liu2001} and \\citet{liubarlow2000}.  The discrepancy between $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$) and $T_{\\rm e}$(Bal) was first discovered by \\citet{peim71} and was attributed to temperature fluctuations within the nebula \\citep{peim67}. He characterized the temperature structure of a nebula by the average temperature $T_0$ and a mean square temperature fluctuation $t^2$ as follows: \\begin{equation} T_0(N_{\\rm e},N_i)=\\frac{\\int T_{\\rm e}N_{\\rm e}N_idV}{\\int N_{\\rm e}N_idV} \\end{equation} and \\begin{equation} t^2=\\frac{\\int (T_{\\rm e}-T_0)^2N_{\\rm e}N_idV}{T_0^2\\int N_{\\rm e}N_idV}, \\end{equation} where $N_i$ is the ion density for an observed emission line. Assuming that all the oxygen is twice ionized and that there are no density and composition variations in the nebula, $T_0$ and $t^2$ can be derived from two of the three expressions, \\begin{equation} T_{\\rm e}([{\\rm O}~{\\rm III}]_{\\rm na})=T_0[1+\\frac{1}{2}(\\frac{9.13\\times10^4}{T_0}-3)t^2], \\end{equation} \\begin{equation} T_{\\rm e}([{\\rm O}~{\\rm III}]_{\\rm fn})=T_0[1+\\frac{1}{2}(\\frac{2.92\\times10^4}{T_0}-3)t^2], \\end{equation} and \\begin{equation} T_{\\rm e}({\\rm Bal})=T_0(1-1.67t^2), \\end{equation} \\citep{peim67,dinerstein}.  In Fig.~\\ref{tecorr} and Fig.~\\ref{tecorrir}, lines showing $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$) as a function of $T_{\\rm e}({\\rm Bal})$ for the case of $t^2=0.00$, 0.02, 0.06 and 1.00 are over plotted. Fig.~\\ref{tecorr} shows that most nebulae have $t^2\\la0.06$. Excluding a few extreme nebulae having $t^2>0.1$, we obtain a mean value of $<t^2>=0.031$. The value is in excellent agreement with previous result $<t^2>=0.03$ derived by \\citet{liu93} although their analysis contained larger measurement errors. However, even excluding the extreme nebulae, a similar comparison between $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm fn}$) and $T_{\\rm e}({\\rm Bal})$ yields an average $<t^2>=0.162$. As we shall show in the following subsection, it is likely that $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm fn}$) have been overestimated because of density variations in the nebula, resulting in higher $<t^2>$ when comparing $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm fn}$) and $T_{\\rm e}({\\rm Bal})$.  Another interesting point revealed by Fig.~\\ref{tecorr} is that $T_{\\rm e}({\\rm Bal})$ covers a wider range of values than $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$). Except a few extreme cases, $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$) falls between 8000 and 14000~K, whereas $T_{\\rm e}({\\rm Bal})$ varies from 900 to 19000~K. \\citet{peim67} has argued that temperature fluctuations of gaseous nebulae can cause lower $T_{\\rm e}({\\rm Bal})$ compared to $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$). However, even $<t^2>=0.031$ is beyond any photoionization model prediction. Another explanation of higher $T_{\\rm e}({\\rm Bal})$ compared to $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$) is that there are some condensations with $N_{\\rm e}>10^6$~cm$^{-3}$ in PNe, which then suppress the $T_{\\rm e}$([\\ion{O}{3}]) $\\lambda\\lambda$4959,5007 relative to the $\\lambda$4363, leading to an artificially higher $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$) \\citep{viegas}. But no such high density condensations have been found in PNe having extremely large temperature discrepancies \\citep{liubarlow2000,liu2001}. Recently, \\citet{liubarlow2000} presented an empirical nebular model containing two components, each with its own temperature, density and chemical composition to account for the temperature discrepancy. This seems to be the most plausible explanation at present.  For PNe having $T_{\\rm e}({\\rm Bal})>T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$), as in He~2-118, the electron temperature might  conceivably increase in low-ionization regions owing to heating by shock waves in the outer regions. As a result, the average temperature of the entire nebula [$T_{\\rm e}$(Bal)] appears higher than that of high-ionization regions [$T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$)].  \\subsection{A simple model with condensations}  In \\S~4.1, we found that electron densities derived from the \\ion{H}{1} recombination spectrum are systematically higher than those derived from optical forbidden lines ratios, and confirm the earlier result of \\citet{liubarlow2001} that densities derived from the [\\ion{O}{3}] far-IR fine-structure lines are the lowest of all. These results show that nebular electron densities derived from various diagnostics are correlated with the critical density of the diagnostic lines, and suggest that dense clumps are generally present in PNe.  In such cases, the [\\ion{O}{3}] $(52\\micron+88\\micron)/\\lambda4959$ ratio is no longer a good temperature indicator given the relatively low critical densities of the [\\ion{O}{3}] $52\\micron$- and 88-$\\micron$ lines (3500 and 1500~cm$^{-3}$, respectively). When nebulae contain condensations with densities in the range between 10$^4$ and 10$^6$~cm$^{-3}$, the [\\ion{O}{3}] $52\\micron$- and 88-$\\micron$ far-infrared lines will be heavily suppressed by collisional de-excitation, leading to overestimates of $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm fn}$). On the other hand, condensations of densities in this range will not affect $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$) given that the [\\ion{O}{3}] $\\lambda4959$ and $\\lambda4363$ lines have higher critical densities of $7\\times10^5$~cm$^{-3}$ and $3\\times10^7$~cm$^{-3}$, respectively (cf. Fig~~\\ref{tecorrir}).  Another piece of evidence pointing to the presence of dense clumps comes from Fig.~\\ref{oiiite} where we compare $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$) and $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm fn}$). If nebulae have temperature fluctuations but are homogeneous in density, then $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm fn}$) is expected to be always lower than $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$) \\citep{dinerstein}. However,  Fig.~\\ref{oiiite} shows about two-third of the objects have higher $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm fn}$) compared to $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$), suggesting that density variations play an important role in $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm fn}$) determinations.  \\citet{viegas} showed that electron temperatures derived from the [\\ion{O}{3}] $\\lambda4363/(\\lambda\\lambda4959+5007)$ ratio may be overestimated if nebulae contain ionized condensations of electron densities higher than 10$^6$~cm$^{-3}$. However, except for a few extreme PNe, such as IC~4997 \\citep{hyung1994}, Mz~3 \\citep{zhang02} and M~2-24 \\citep{zhang2003} there is no evidence that nebulae contain a substantial amount of ionized gas in condensations with densities in excess of 10$^6$~cm$^{-3}$, as the presence of such condensations should have been revealed by the \\ion{H}{1} recombination spectrum (Liu et al. 2000, 2001a, 2001b; Fig.~\\ref{necorr}). Therefore, we consider $N_{\\rm e}\\sim10^5$~cm$^{-3}$ as a safe upper limit for the density of these condensations.  In order to estimate the filling factor of the condensations, we use an analytical method originally developed by \\citet{viegas} for explaining the discrepancy of $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$) and $T_{\\rm e}({\\rm Bal})$. Assuming that some high-density regions ($N_{\\rm e}>10^6$~cm$^{-3}$) are present, they deduced nebular-to-auroral [\\ion{O}{3}] line ratios (see their Eq.~2.3). Following the same deduction of their Eq.~2.3 but assuming that the electron density of these condensations is about $10^5$~cm$^{-3}$ instead of $10^6$~cm$^{-3}$, we obtain the infrared-to-auroral line ratio \\begin{equation} \\frac{I_{52\\micron+88\\micron}}{I_{4959}}=\\frac{\\varepsilon_{\\rm fL}}{\\varepsilon_{\\rm nL}}\\frac{1+\\mu^2\\omega(\\varepsilon_{\\rm fH}/\\varepsilon_{\\rm fL})}{1+\\mu^2\\omega(\\varepsilon_{\\rm nH}/\\varepsilon_{\\rm nL})}, \\end{equation} where $\\varepsilon$ is line emission coefficient, $4\\pi j=n_{\\rm e}n_i\\varepsilon$, f and n refer to IR fine-structure line and nebular line, L and H refer to the low- and high-density regions, $\\mu=(N_{\\rm e})_{\\rm H}/(N_{\\rm e})_{\\rm L}$ is the density contrast and $\\omega=V_{\\rm H}/V_{\\rm L}$ is the filling factor. In the following calculation, we assume that: 1) density variations are the only cause of the lower $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm fn}$) compared to $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$) and $T_{\\rm e}(\\rm Bal)$; 2) the nebular gas has a homogeneous electron temperature, i.e., $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$) or $T_{\\rm e}(\\rm Bal)$; 3) the condensations have densities lower than the critical density of the [\\ion{O}{3}] auroral line but higher than the critical densities of the [\\ion{O}{3}] infrared lines, whereas low-density regions have a density that is lower than the critical densities of the [\\ion{O}{3}] auroral and infrared lines.  Setting $F(T_{[{\\rm O}~{\\rm III}]_{\\rm fn}})=I_{52\\micron+88\\micron}/I_{4959}$, we have $\\varepsilon_{\\rm fL}/\\varepsilon_{\\rm nL}=F(T_{\\rm Bal})$. For a simple assumption $(N_{\\rm e})_{\\rm H}=10^5$~cm$^{-3}$ and $(N_{\\rm e})_{\\rm L}=10^3$~cm$^{-3}$, the ratio $\\varepsilon_{\\rm fH}/\\varepsilon_{\\rm fL}$ is about 0.04, whereas $\\varepsilon_{\\rm n}$ is practically independent of the density when the densities are lower than the critical density of the nebular line and thus $\\varepsilon_{\\rm nH}/\\varepsilon_{\\rm nL}\\approx1$.  Therefore, $\\mu^2\\omega$ is given by \\begin{equation} \\mu^2\\omega=\\frac{F(T_{[{\\rm O}~{\\rm III}]_{\\rm fn}})/F(T_{\\rm Bal})-1}{0.04-F(T_{[{\\rm O}~{\\rm III}]_{\\rm fn}})/F(T_{\\rm Bal})}. \\end{equation}  In Fig.~\\ref{tecorrir}, the lines showing the variation of $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm fn}$) as a function of $T_{\\rm e}$(Bal) for $\\mu^2\\omega=0.0$, 0.1, 0.5, 1.0 and 2.0 are plotted. We obtain an average value of $<\\mu^2\\omega>=1.0$. However, it should be borne in mind that the $\\mu^2\\omega$ derived by eq.~(18) is a upper limit since temperature variations are actually present and contribute partly to the measured temperature discrepancies.  $\\mu^2\\omega$ can also be estimated by comparing $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$) with $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm fn}$), as shown in Fig.~\\ref{oiiite}. Assuming that $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$) is constant across the entire nebula, $\\mu^2\\omega$ is likewise given by \\begin{equation} \\mu^2\\omega=\\frac{F(T_{[{\\rm O}~{\\rm III}]_{\\rm fn}})/F(T_{[{\\rm O}~{\\rm III}]_{\\rm na}})-1}{0.04-F(T_{[{\\rm O}~{\\rm III}]_{\\rm fn}})/F(T_{[{\\rm O}~{\\rm III}]_{\\rm na}})}. \\end{equation} In Fig.~\\ref{oiiite}, the lines showing the variation of $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm fn}$) as a function of $T_{\\rm e}$(Bal) for $\\mu^2\\omega=0.0$, 0.1, 0.5, 1.0 and 2.0 are also plotted. An average value of $<\\mu^2\\omega>=0.4$ is obtained. It is noteworthy that the presence of temperature variations will increase $T_{\\rm e}$([\\ion{O}{3}]${\\rm na}$) with respect to $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm fn}$), in contrast to density variations. Therefore, $\\mu^2\\omega$ derived from eq.~(19) is a lower limit.  The upper- and lower-limit of $\\mu^2\\omega$ are listed in Table.~\\ref{tbl-3b}. For a few PNe having a lower limit $(\\mu^2\\omega)_{\\rm L}=0.0$, the effect of temperature variations on the temperature determination is larger than that of density variations, causing higher $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na})$ compared to $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm fn}$). For all the PNe, the mean value of $\\mu^2\\omega$ is in the range from 0.4 to 1.0. Thus, we suggest that $\\mu^2\\omega=0.7$ may be a representative value that can be used for a typical nebula. Adopting the density contrast $\\mu\\sim100$, a representative value of the condensations filling factor $\\omega$ is about $7\\times10^{-5}$. The value is similar to that found by \\citet{viegas}. Adopting $\\mu^2\\omega=0.7$ for all the PNe, we obtain new $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm fn}$), for which the contribution of the clumps has been deduced using equation~(17). Comparing the new $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm fn}$) and $T_{\\rm e}$(Bal), we re-evaluate the average value $<t^2>=0.032$, which is in good agreement with the value yielded by the comparison of $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$) and $T_{\\rm e}$(Bal) (see Section~4.2).  \\subsection{The difference between temperatures derived from the [\\ion{O}{3}] forbidden lines and from the Balmer jump versus electron density}  In Fig.~\\ref{crene}, the differences between temperatures derived from the [\\ion{O}{3}] forbidden lines and from the Balmer jump are plotted against electron densities derived from the hydrogen recombination spectrum, the [\\ion{O}{2}] $\\lambda3729/\\lambda3726$, [\\ion{S}{2}] $\\lambda6716/\\lambda6731$, [\\ion{Cl}{3}] $\\lambda5517/\\lambda5537$ [\\ion{Ar}{4}] $\\lambda4711/\\lambda4740$ and the [\\ion{O}{3}] $52\\mu$m/$88\\mu$m ratios, respectively.  In all cases, the differences between the two temperature are anti-correlated with nebular densities derived from various density indicators. A least-squares fit yields the following relation: \\begin{equation} \\Delta T=(1.30\\pm0.12)\\times10^4-(2.93\\pm0.27)\\times10^3\\log N_{\\rm e}, \\end{equation} and a linear correlation coefficient of 0.51, where $\\Delta T=T_{\\rm e}$([\\ion{O}{3}])$-T_{\\rm e}({\\rm Bal})$ and $N_{\\rm e}$ is that derived from hydrogen recombination spectrum, which represents the average density of the entire nebula. \\citet{garnett} found that the discrepancies of O$^{2+}$/H$^{+}$ abundance ratios derived from ORLs and from CELs are greater for larger, lower-surface brightness PNe than for compact, dense ones. On the other hand, \\citet{liu2001} found that the abundance discrepancies are greater for those with larger temperature discrepancies. Our results are therefore consistent with these earlier findings. \\citet{liu2001} found that, \\begin{equation} \\Delta({\\rm O}^{2+}/{\\rm H}^+)=(0.209\\pm0.085)+(2.01\\pm0.33)\\times10^{-4}\\Delta T, \\end{equation} where $\\Delta({\\rm O}^{2+}/{\\rm H}^+)=\\log({\\rm O}^{2+}/{\\rm H}^+)_{\\rm ORL}-\\log({\\rm O}^{2+}/{\\rm H}^+)_{\\rm CEL}$. From the two relations given in equations (20) and (21), we find that the ratio of abundances derived from ORLs and from CELs has a power-law dependence on density, \\begin{equation} \\frac{{\\rm O}^{2+}_{\\rm ORL}}{{\\rm O}^{2+}_{\\rm CEL}}\\propto N_{\\rm e}^{-\\alpha}, \\end{equation} where $\\alpha=0.59\\pm0.15$.  In general, electron densities of PNe are expected to decline monotonically as nebulae expand. \\citet{sk1987} found an approximate relation between nebular mean electron density and nebular size, $N_{\\rm e}\\propto R^{-3}$, where R is the nebular radius. Consequently, it is suggested that the temperature discrepancy is related to the evolutionary state of PNe; more expended PNe show larger temperature discrepancies. On the other hand, the actual evolution in a given nebula may depend on many factors. Furthermore, the electron density measured for a given nebula may depend on the diagnostic used. For example, electron densities of low-ionization regions, as probed by the [\\ion{O}{2}] $\\lambda3729/\\lambda3726$ and [\\ion{S}{2}] $\\lambda6716/\\lambda6731$ line ratios, can be enhanced by interactions of fast stellar winds with earlier slow winds. Therefore, there are slight difference in deriving the relation of $\\Delta T$ and $N_{\\rm e}$ from different density diagnostics, as shown in Fig.~\\ref{crene}. We also use electron density derived from these forbidden line ratios to estimate the $\\alpha$-value in equation (22). The results show that the value is within the range of 0.3 and 0.9.  To explain the relation shown in equation (20) and (22), we apply a chemically inhomogeneous two-component model presented by \\citet{liubarlow2000}. In their hypothesis, the nebula consists of two components, H-deficient clumps and diffuse material with `normal' abundances ($\\sim$solar). These H-deficient clumps are expected to have a low electron temperature due to enhanced cooling caused by high metallicity and pressure equilibrium between the two components, so that only recombination lines are emitted. As a result, ORLs yield overestimated heavy element abundances. In a given PN, the contribution of the two components to the Balmer temperature and abundance determinations depends on their densities, i.e., the contribution becomes weaker as density declines. If the clumps have neutral cores, as found in Abell~30 \\citep{ercolano} and NGC~7293 \\citep{dyson}, they are expected to be quite stable during the evolution of the PNe. As a result, the relative contribution of the low-density and high-temperature material to the average temperature and abundance determinations decrease as nebulae expand. Consequently, the discrepancies in temperature and abundance determinations become larger as the nebulae expend.  \\subsection{Are He/H ratios generally overestimated?}  A long-standing problem in abundance determinations for PNe has been that the heavy element abundances derived from ORLs are systematically higher than those derived from CELs. The ratio of abundances derived from the two types of emission lines is found to vary from target to target. \\citet{tsamis} analysed a PN sample and found the ORL/CEL discrepancy factors for these heavy element abundances span a range from about 2 to as large as 20. For the most extreme known PNe Hf~2-2, the ORL/CEL abundance ratio reaches a record value of 84 \\citep{liu209}. Another problem in nebular astrophysics is that $T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$) are systematically higher than $T_{\\rm e}({\\rm Bal})$, as we have discussed in \\S~4.2. The two problems are found to be related. A positive correlation between the ratio of O$^{2+}$/H$^+$ derived from \\ion{O}{2} ORLs and from [\\ion{O}{3}] CELs and the temperature discrepancy [$\\Delta T$=$T_{\\rm e}$([\\ion{O}{3}]$_{\\rm na}$)-$T_{\\rm e}({\\rm Bal})$] has been shown by \\citet{liu2001}.  The explanations for the dichotomy in abundance and temperature determinations have focussed on temperature and density variations and chemical inhomogeneities. However, it has been shown that the temperature and density variations alone can be ruled out by current observations \\citep{liua,liu209}. At present, the most plausible explanation is that there exist some cold H-deficient clumps in nebula \\citep{liubarlow2000} where only recombination lines are emitted.  To fit the observed pattern of NGC~6153, \\citet{liubarlow2000} have presented several empirical models incorporating H-deficient clumps. \\citet{pequignot} also constructed photoionization models of NGC~6153 and M~1-42, incorporating H-deficient clumps , which reproduce the observed integrated line ratio satisfactorily. In all the models, the H-deficient clumps contain only about $1\\%$ of the total mass, suggesting that the H-deficient components contain only a tiny amount of nebular gas. Nevertheless, they can seriously enhance the intensities of ORLs, causing overestimated ORL abundances. Therefore, the He/H abundance ratio could be overestimated owing to the existence of the H-deficient clumps can lead to since all the observable strong \\ion{He}{1} and \\ion{He}{2} lines are recombination lines. Some H-deficient components have been found in some PNe, such as Abell~30 and Abell~78 \\citep{jacoby}. \\citet{wesson} analysed optical spectra of the H-deficient knots in Abell~30 and found the He/H abundance ratios of the knots are more than ten.  These PNe with extremely large abundance and temperature discrepancies, such as NGC~6153, M~1-42 and Hf~2-2, have been found to have high He/H abundances. \\citet{peimbert1983} defined the PNe having high He/H abundance ($\\ge0.125$) as Type-I PNe. According to our suggestion, these PNe with high He/H abundances might result from contamination from H-deficient clumps. Quantitative analysis of overestimated He/H abundances in PNe is quite difficult. For this purpose, as mentioned by \\citet{pequignot}, one has to rely on the measurement of \\ion{He}{1} $\\lambda$10830 line which is strongly enhanced by collisional excitation from the 2s$^3$S metastable level.  If the He/H abundance is enhanced by H-deficient clumps, one expects that there is a positive correlation between the He/H abundance and the ORL/CEL abundance ratio of heavy element. For most PNe, however, heavy element ORLs are quite weak and thus yield large uncertainties in the resultant abundance. Instead, we can study the relation between the He/H abundance and $\\Delta T$ since the ORL/CEL abundance ratio has been found to be positively correlated with $\\Delta T$ by \\citet{liu2001}. Fig.~\\ref{hete} plots He/H versus $\\Delta T$. The measurement error of the He/H abundance ratio is negligible. In the figure, filled and open circles represent the Galactic disc and bulge PNe respectively and the solid line is a linear fit represented by the relation: \\begin{equation} {\\rm He}/{\\rm H}=(0.056\\pm0.003)+(4.25\\pm0.11)\\times10^{-5}\\Delta T, \\end{equation} which yields a linear correlation coefficient of 0.50. Although there is considerable scatter, there is a trend of increasing He/H abundance with increasing difference between the temperature from [\\ion{O}{3}] forbidden lines and Balmer discontinuity for both the Galactic disc and bulge PNe. Therefore, it seems plausible that the He/H abundances in PNe have been generally overestimated due to the existence of cold H-deficient clumps.  Fig.~\\ref{hete} also shows that the most of PNe have enhancement of helium with respect to the Sun, which is often ascribed to the second and third dredge-up. However, the normal model of dredge-up cannot explain some extreme cases, such as He~2-111 which has a He/H abundance of 0.219 \\citep{kingsburgh}. For these cases, the more likely reason is that pollution of the H-deficient clumps causes an overestimate of the He/H abundance ratio. On the other hand, we cannot completely reject the possibility that helium in nebulae has been partly enhanced by the products of nucleosynthesis of the central star. That may be the main cause of the scatter in the derived He/H abundances. The presence of hydrogen-deficient clumps may be further proof.  The ORL/CEL abundance discrepancies are also found in \\ion{H}{2} regions \\citep{esteban,tsamisa}. And the existence of some high density clumps in \\ion{H}{2} regions has been indicated by previous studies (e.g. O'Dell et al. 2003). Therefore, the H-deficient material might also be present in \\ion{H}{2} regions, leading to overestimated ORL abundances. If this is the case, the primordial helium abundance, which is based on observations of some metal poor extragalactic \\ion{H}{2} regions, could have been overestimated.    In this work, we presented a method to derive nebular electron temperature and density simultaneously by matching the theoretical with the observed hydrogen recombination spectra. Using the method, we derived electron temperatures and densities of 48 Galactic PNe. The temperatures and densities derived by this method are compared with those derived from other indicators, enabling us to study electron temperature and density variations in PNe. Our main findings are as follows.  \\begin{list}{}{} \\item[---] The densities derived from hydrogen recombination spectra are generally higher than those deduced from forbidden line ratios, suggesting that condensations are generally present in nebulae.  \\item[---] The temperatures derived from the nebular-auroral line ratios of [\\ion{O}{3}] are systematically higher than those derived from hydrogen recombination spectra, suggesting that temperature variations are generally present in nebulae. We obtained the rms temperature fluctuation parameter $t^2=0.031$ as a representative value.  \\item[---]The larger discrepancy in temperatures derived from [\\ion{O}{3}] $(52\\micron+88\\micron)/\\lambda4959$ line ratios and from hydrogen recombination spectra, however, cannot be attributed simply to the temperature variations. The presence of clumps with $N_{\\rm e}\\sim10^5$~cm$^{-3}$ can account well for the discrepancy. A comparison between electron temperatures derived from [\\ion{O}{3}] $(52\\micron+88\\micron)/\\lambda4959$ line ratios and those derived from hydrogen recombination spectra and [\\ion{O}{3}] $\\lambda4959/\\lambda4363$ line ratios suggests that the filling factor of the condensations has a typical value of $7\\times10^{-5}$.  \\item[---]The discrepancies in temperatures derived from [\\ion{O}{3}] $\\lambda4959/\\lambda4363$ line ratios and from hydrogen recombination spectra are related to the evolution of PNe. The ORL/CEL abundance ratio is found to have a power-law dependence on electron density ($\\sim N_{\\rm e}^\\alpha$, $\\alpha=0.59\\pm0.15$). We attribute this to different evolutionary scenario of two components in PNe, i.e. expanding diffuse nebular material and stable cold H-deficient clumps. The contribution of the clumps to the temperature and abundance determinations increases as the nebulae expand.  \\item[---]He/H abundance has a positive correlation with the temperature discrepancy, suggesting that He/H abundances might be overestimated due to the existence of H-deficient clumps. If this is the case, traditional up-dredge theory should be modified. The primordial helium abundance determination will be affected if there are significant condensations in \\ion{H}{2} regions. Further investigations are needed.  \\item[---] The densities derived from the spectrum near the Paschen jump are in good agreement with those derived form that near the Balmer jump. However, the former yield generally lower temperatures than the latter. The reason remain unknown. \\end{list}  We conclude that the existence of H-deficient clumps in PNe is the most plausible explanation for the origin of temperature and density variations and the ORL/CEL abundance discrepancies. However, it is a challenge to clarify the origin and evolution of these clumps. In further work, we will study the statistical relationships between the temperature and abundance discrepancies and other physical conditions of PNe such as radio flux, morphology, chemical abundances and the nature of the central star."
    },
    "0403/astro-ph0403147_arXiv.txt": {
        "abstract": "We discuss the structural change and degree of mass segregation of young dense star clusters within about 100\\,pc of the Galactic center. In our calculations, which are performed with GRAPE-6, the equations of motion of all stars and binaries are calculated accurately but the external potential of the Galaxy is solved (semi)analytically.  The simulations are preformed to model the Arches star cluster.  We find that star clusters with are less strongly perturbed by the tidal field and dynamical friction are much stronger affected by mass segregation; resulting in a significant pile-up of massive stars in the cluster center. At an age of about 3.5\\,Myr more than 90 per cent of the stars more massive than $\\sim 10$\\,\\msun\\, are concentrated within the half-mass radius of the surviving cluster.  Star clusters which are strongly perturbed by the tidal field of the parent Galaxy are much less affected by mass segregation. ",
        "introduction": "In recent years a relatively new class of star clusters has been discovered.  These systems are variously referred to in the literature as young populous clusters, super star clusters, proto-globular clusters, and even young globular clusters (although it remains unclear if they are in any way related to the old globular clusters observed in many galactic halos).  We prefer the term ``young dense cluster'' (hereafter YDC, sometimes pronounced ``YoDeC'') because it highlights the key defining properties of these compact stellar systems.  The fact that these clusters are young means that stars of all masses are still present, offering critical insights into the stellar initial mass function and cluster structural properties at formation.  The term ``dense'' means that dynamical evolution and physical collisional processes can operate fast enough to compete with and even overwhelm stellar evolutionary timescales.  Dense stellar systems are places where wholly new stellar evolution channels can occur, allowing the formation of stellar species completely inaccessible by standard stellar and binary evolutionary pathways.  \\begin{table*}[bp!] \\caption[]{\\footnotesize Observed properties of selected young, dense star clusters, mainly in or near the Milky Way.  Mass ($M$), age, tidal radius ($r_{\\rm tide}$) and half-mass radius ($r_{\\rm hm}$) are taken from the literature. } \\begin{flushleft} \\begin{tabular}{ll|rrrrrrr} \\hline Name   &ref&   $\\log M/\\msun$& $\\log$(age/yr) & location & $r_{\\rm tide}$ (pc)& $r_{\\rm hm}$ (pc) \\\\ \\hline Arches      &a& 4.8~~~~~ & 6.5~~~~~  & GC~   &   1~~  & 0.23~  \\\\ Quintuplet  &b& 4.2~~~~~ & 6.6~~~~~  & GC~   &   1~~  & 0.5~~   \\\\ \\hline NGC\\,3603   &c& 4.3~~~~~ & 6.5~~~~~  & disk~ &  10~~  & 0.78~  \\\\ Westerlund~1&d& 4.5~~~~~ & 6.8~~~~~  & disk~ &  10~~  & 0.2~~  \\\\ R\\,136      &e& 4.7~~~~~ & 6.5~~~~~  & LMC~  &$>20$~~ & 0.5~~  \\\\ \\hline MGG-11      &f& 5.5~~~~~ & 7.0~~~~~  & M82~  &$>20$~~ & 1.2$^\\star$~  \\\\ \\hline \\end{tabular} \\\\ \\medskip\\footnotesize References: a) Figer et al.~(1999;2002;2004); b) Glass et al.~(1987); c) Brandl (1999); d) Vrba et al.~(2000); e) Brandl et al.~(1996); f) McCrady et al. (2003)\\\\ $^\\star$ The projected half light radius $r_{\\rm hl}=1.2$\\,pc, and we take $r_{\\rm hm} = \\frac43 r_{\\rm hl}$ (Spitzer 1987). \\end{flushleft} \\label{Tab:observed} \\end{table*}  Table \\ref{Tab:observed} presents an overview of known YDCs in the neighborhood of the Milky Way Galaxy, for which the relevant parameters are well determined.  The clusters listed in the table are selected on the basis of age ($\\aplt 10$\\,Myr) and half-mass relaxation time ($\\aplt 100$\\,Myr); the latter criterion is not universally used, but it best reflects the interplay between dynamics and stellar evolution just described---clusters satisfying this condition are expected to experience significant dynamical evolution before their massive stars explode as supernovae.  All the clusters in Table \\ref{Tab:observed} thus lie in the regime where dynamical and stellar evolution cannot be considered independently.  In the paper we concentrate on the Arches star cluster, which was observed in detail by Figer et al (1999; 2002) and Stolte (2002).  The $\\sim 3$\\,Myr old Arches cluster is located at a projected distance of about 25\\,pc from the Galactic center, has a half mass radius of about 0.23 parsec and contains between 20\\,000 and $10^5$ stars. Together with the Quintuplet star cluster (Glass 1987) they are the only young dense star clusters which are strongly perturbed by the tidal field of their parent Galaxy. McMillan et al (2004) discusses the internal dynamical evolution and the possibility of the formation of an intermediate mass black hole in the star cluster MGG-11 in the starburst galaxy M82 (see also Portegies Zwart et al 2004), where Baumgardt et al (2004) discusses the further consequences of the presence of such a black hole.  Figure~\\ref{Fig:composite} shows a composite image of various star clusters on the same scale. Among these are, next to some other YDCs the Arches and Quintuplet systems. For comparison we added images of the globular cluster M80, the Pleiades and the Trapezium cluster, the star forming region in Orion.  We note that the latter three systems are not considered YDCs, but are added for comparison.   \\begin{figure}[htbp!] \\psfig{figure=./composite.ps,width=0.8\\linewidth,angle=-90} \\caption[]{Composite image of several star clusters, all on the same scale.  The big image to the left is the globular cluster M80. To the right are a selection of young dense star clusters, identified by their name (except the bottom image which represents the Pleiades star cluster.  } \\label{Fig:composite} \\end{figure}  We make the distinction between two families of YDCs; those that are isolated and those that are strongly perturbed by the external tidal field of their parent Galaxy.  We know only two clusters that are in the latter category, Arches and Quintuplet, the other YDCs are relatively unperturbed.  The evolution of tidally perturbed or unperturbed clusters are profoundly different. In the limit of a strong tidal field, stars near the tidal radius of the cluster feel the proximity of the background Galaxy. ",
        "conclusions": "We perform simulations of young dense star clusters of 65536 stars of which 10486 have a close binary companion. The clusters are initially on three different orbits around the Galactic center. At an age of 3.5\\,Myr we compare the mass function of the surviving cluster with the observed present day mass function of the Arches cluster. From this comparison we conclude that the present-day mass function of the Arches cluster is highly dynamically evolved.  In the case this effect is due to the dynamical evolution of the cluster, we conclude that the cluster is not on an elliptic orbit which brings it within about 15\\,pc from the Galactic center."
    },
    "0403/astro-ph0403001_arXiv.txt": {
        "abstract": "We review evidence that the census of Milky Way satellites similar to those known may be incomplete at low latitude due to obscuration and in the outer halo due to a decreasing sensitivity to dwarf satellites with distance. We evaluate the possible impact that incompleteness has on comparisons with substructure models by estimating corrections to the known number of dwarfs using empirical and theoretical models. Under the assumption that the true distribution of Milky Way satellites is uniform with latitude, we estimate a 33$\\%$ incompleteness in the total number of dwarfs due to obscuration at low latitude.  Similarly, if the radial distribution of Milky Way satellites matches that of M31, or that of the oldest sub-halos or the most massive sub-halos in a simulation, then we estimate a total number of Milky Way dwarfs ranging from 1 -- 3 times the known population.  Although the true level of incompleteness is quite uncertain, the fact that our extrapolations yield average total numbers of MW dwarfs that are realistically 1.5 -- 2 times the known population, shows that incompleteness needs to be taken seriously when comparing to models of dwarf galaxy formation. Interestingly, the radial distribution of the oldest sub-halos in a $\\Lambda$CDM simulation of a Milky Way-like galaxy possess a close match to the observed distribution of M31's satellites, which suggests that reionization may be an important factor controlling the observability of sub-halos.  We also assess the prospects for a new SDSS search for Milky Way satellites to constrain the possible incompleteness in the outer halo. ",
        "introduction": "\\label{sec:intro}  The currently favored $\\Lambda$ + cold dark matter ($\\Lambda$CDM) cosmological model successfully reproduces many of the observed large-scale properties of the Universe, including the properties of the Cosmic Microwave Background recently observed by WMAP \\citep{spergel03}, the number, size and clustering of galaxy clusters (e.g. \\citealt{eke96,zehavi02}), and the evolution of galaxy cluster counts \\citep{rosati02}.  However, several major discrepancies between the predictions of $\\Lambda$CDM and the observed properties of the Universe on small scales have presented challenges to the paradigm. One outstanding challenge is that Cold Dark Matter models predict over an order of magnitude more low mass, dark matter halos around the Milky Way than the number of observed satellite dwarf galaxies.  This discrepancy was first pointed out by \\citet{kauffmann93}, and was later confirmed by high resolution N-body simulations \\citep{klypin99,moore99,font01}.   Within the $\\Lambda$CDM framework, a plausible explanation for the discrepancy between the number of predicted sub-halos and observed satellites lies within baryonic physics.  The fraction of dark matter halos with $v_c \\lesssim$ 50 km sec$^{-1}$ that may host a luminous galaxy can be significantly reduced by reionization, feedback, and/or tidal effects (\\citealt{dekel86,efstathiou92,thoul96,quinn96,bkw00,bkw01,benson02,susa03,dijkstra04}, among others).  These baryonic processes are difficult to model. Therefore, the exact extent to which each effects the present day luminosity of dark matter sub-halos remains uncertain. A comparison of the total number and the radial distribution of predicted visible satellites with that of the observed Milky Way satellites may provide a test of the feasibility of particular models (e.g. \\citealt{taylor03,kravtsov04}).  A number of the above implementations of baryonic physics have been able to reproduce the observed Milky Way dwarf population. Unfortunately, existing comparisons are rendered less meaningful by the uncertain completeness of the Milky Way dwarf satellite population.  Due to incompleteness, the observed satellites may not reflect the properties of the underlying population.  Models that provide a good match to the current observations may, therefore, actually underpredict the underlying population.  Past searches for Milky Way companions, although very successful, suffer from unavoidable observational biases that could lead to an undercounting of Milky Way satellites both at low Galactic latitudes and at large ($>$ 100 kpc) distances (see \\S3 for discussion). Furthermore, the extent of these possible biases is not well understood due to a lack of systematic analyses (however, see \\citealt{kleyna97} for a systematic analysis of their survey's sensitivity).  \\citet{willman02} are currently implementing a new search for resolved Milky Way dwarf satellites in the Sloan Digital Sky Survey data.  In contrast to past surveys, this search is sensitive to dwarfs similar to and much fainter than any among the known population, at any distance out to the Milky Way's virial radius.  The SDSS may thus provide the means to evaluate the possibility of undercounting in the outer halo.  In light of the new Willman et al survey, we reevaluate the evidence for incompleteness to Milky Way satellites similar to those in the known population and highlight the possible impact of such an incompleteness on comparisons with substructure models. In \\S2, we describe the N-body cosmological simulation of a Milky Way like galaxy that we use to compare with observations. In \\S3, we review observational evidence for bias in the census of Milky Way companions and estimate the possible number of undetected galaxies similar to those known, based on primarily observational arguments.  In \\S4, we compare the radial distribution of Milky Way dwarfs with that of both the oldest and highest v$_c$ dark matter sub-halos of the simulated galaxy.  We use this comparison both to demonstrate how well radial distributions may be used to distinguish between models and to underscore the possible impact of observational bias on such a comparison.  Finally, in \\S5 we estimate the number of dwarfs similar to the known population that the Willman et al. (2002) survey could detect and still be consistent with either of the two models that we consider. ",
        "conclusions": "In this paper, we have reviewed evidence for incompleteness in the known Milky Way dwarf satellites with properties similar to those in the known population.  This possible incompleteness is due to past observational bias against detecting Milky Way satellites at low latitude and in the outer Galactic halo.  Although the level of incompleteness is very uncertain, the fact that an empirical extrapolation from the M31 distribution yields an average total number of MW dwarfs that is 1.5 -- 2 times the known population shows that incompleteness needs to be taken seriously when comparing to models of dwarf galaxy formation.  We used the oldest and highest v$_c$ sub-halos of a simulated Milky Way-like galaxy to demonstrate how radial distributions may be used to distinguish between proposed models of dwarf galaxy formation. However, KS tests comparing the radial distributions of the Milky Way, M31, and the oldest sub-halos and the highest v$_c$ sub-halos in simulations show that they are all at least marginally consistent with each other. Interestingly, the M31 distribution is consistent with the oldest sub-halo distribution at $> 95\\%$, suggesting that reionization may have a substantial effect on the observabilitiy of sub-halos (however, see Kravtsov et al. 2004). However, small numbers and cosmic scatter permit at least a marginal consistency between a wide range of observations and models. It is thus difficult at present to use radial distributions alone to clearly distinguish between substructure models, although they certainly provide a complimentary test of model predictions. Though faint galaxy membership in other groups is currently controversial, when the satellite populations of galaxies in nearby groups are known more precisely, their radial distributions will provide a stronger discriminant between models. Likewise, a large ensemble of high resolution simulations will allow a more robust assessment of the effects of cosmic scatter and small numbers on the predicted satellite population.  The crude arguments presented in this paper result in predicted total numbers of dwarfs that range from 1 -- 3 times the known number, with the most realistic estimates producing an incompleteness in the Milky Way dwarf census of up to 50$\\%$. The exact level of incompleteness is strongly dependent on the distribution of Milky Way satellites with latitude.  If the Milky Way census is incomplete at the level of 50$\\%$ or more, then many existing models underpredict the number of luminous Milky Way dwarfs.  In particular, models with suppressed small scale power would not produce enough luminous dwarf galaxies, as has already been suggested by \\citet{chiu01}.  Note that our derived ``total incompleteness'' only accounts for dwarfs with properties similar to those known, not any fainter dwarfs, should they exist.  Currently, the largest uncertainty in the known Milky Way population is the underlying distribution of dwarfs with latitude.  Before the Milky Way satellite population can yield a meaningful comparison with substructure models, this uncertainty needs to be investigated with more detail than presented in this paper. The average incompleteness in the current census, assuming a uniform distribution in latitude and no other incompleteness, is 33$\\%$.  Any survey sensitive to faint dwarf satellites at low latitude would thus place very valuable constraints on the Local Galaxy luminosity function."
    },
    "0403/hep-ph0403158_arXiv.txt": {
        "abstract": " ",
        "introduction": "High-energy colliders are the tool to discover new heavy particles with sizable couplings. New  light particles with small couplings can be searched for in many different ways. Neutral fermions with eV-scale mass, called `sterile neutrinos' in the jargon, are typically stable enough to give effects in cosmology, and can affect neutrino oscillation experiments.\\footnote{Small fermions masses are stable under quantum corrections. Cosmology, astrophysics and neutrino experiments are also sensitive to light scalars or vectors, which give different kinds of signals.} Both fields recently discovered new physics (in both cases around meV energies), but so far no sterile neutrinos.  Today the most powerful cosmological probe of sterile effects is standard Big-Bang Nucleosynthesis (BBN)~\\cite{BBN}, which constrains the number of thermalized neutrinos present at $T\\sim 0.1\\MeV$ to be $N_\\nu = 2.5\\pm 0.7$. Since the uncertainty is controversial it is not clear how much $N_\\nu=4$ is disfavoured.  The established solar and atmospheric neutrino anomalies seem produced by oscillations among the three SM neutrinos, with at most minor contributions from possible extra sterile neutrinos. The present 99\\% C.L. bounds on sterile mixing, computed {\\em assuming} that the initial active neutrino $|\\nu_{\\rm a}\\rangle$ oscillates with a large $\\Delta m^2$ into an energy-independent mixed neutrino $\\cos\\theta_{\\rm s}|\\nu'_{\\rm a}\\rangle + \\sin\\theta_{\\rm s}|\\nu_{\\rm s}\\rangle$, are $\\sin^2\\theta_{\\rm s} <0.25$ in solar oscillations (where $\\nu_{\\rm a}=\\nu_e$)~\\cite{etas} and $\\sin^2\\theta_{\\rm s} <0.21$ in atmospheric oscillations (where $\\nu_{\\rm a}=\\nu_\\mu$)~\\cite{etaa}.  \\bigskip  We relax these simplifying assumptions and study the more general $4$-neutrino context. This demands a remarkable effort, but it is an important task for the present, for two reasons. First, because recent discoveries in cosmology and neutrino physics are stimulating  new experiments that  will study the new phenomena with the redundancy necessary to test the minimal models suggested by present data. Second, because oscillations in extra light particles are the natural extension of the emergent massive neutrinos scenario.     Although we list hints that might be interpreted as  sterile neutrino effects (LSND~\\cite{LSND}, certain pieces of solar or atmospheric or BBN data, ...), we do not focus on any of them in particular. Rather, we compute experimental capabilities and constraints on sterile oscillations and compare them with capabilities  and constraints from various present and future cosmological and astrophysical probes. Some important constraints are based on untested assumptions and plagued by systematic uncertainties. In such cases, rather than performing global fits, we identify and compute the key observables, trying to explain the basic physics in simple terms emphasizing the controversial issues, so that the reader can judge.  \\medskip  The paper is organized as follows. In section~\\ref{osc} we briefly review theoretical motivations: as in the case of the known light particles, some fundamental reason would presumably be behind the lightness of an extra sterile neutrino. We also describe the non-standard  parametrization of active/sterile mixing that we choose (because more convenient and intuitive than standard parametrizations) and describe the qualitatively different kinds of spectra on which we will focus.  In section~\\ref{cosmo} we study sterile effects in cosmology, comparing the relative sensitivities of two BBN probes (the helium-4 and deuterium abundances), of Cosmic Microwave Background (CMB) and of Large Scale Structures (LSS). In section~\\ref{solar} we study sterile oscillations in solar (and KamLAND) neutrinos. In section~\\ref{SN} we study sterile oscillations in SN1987A and future supernov\\ae. We also briefly discuss other less promising probes (relic SN background, high-energy cosmic neutrinos,\\ldots). In section~\\ref{atm} we study sterile oscillations in atmospheric and reactor neutrinos, and in short and long-baseline neutrino beams (including LSND).  In order to avoid a unreadably long paper each section is written in a concise way and contains a `Results' subsection, which can be read skipping the other more technical parts. From a computational point of view, exploring 4$\\nu$ oscillations is $3\\div 4$ orders of magnitude more demanding than usual 2 or 3$\\nu$ fits and therefore requires significant improvements of usual techniques. Subsections entitled `technical details' describe how this was achieved. In particular we describe how we employ the density matrix formalism in neutrino oscillation computations, and explain why even in simple $3\\nu$ situations it gives significant advantages with respect to more direct equivalent approaches.   In section~\\ref{all} we conclude by summarizing and comparing the different probes of sterile neutrinos. ",
        "conclusions": "\\label{all} A few years ago active/sterile oscillations were studied as an alternative to active/active oscillations. For example, % it was shown that neither the solar nor the atmospheric anomaly can be produced by oscillations into sterile neutrinos compatibly with standard BBN. These active vs sterile issues  have now been firmly solved by experiments, and the new relevant questions become: \\begin{quote} How large can be the subdominant sterile component possibly present in solar or atmospheric oscillations? How can we discover new anomalies due to sterile effects? \\end{quote} In order to address these issues we systematically  compared oscillation effects generated by one sterile neutrino (including the effects of the now established solar and atmospheric oscillations) with present experiments and studied capabilities of future probes, extending previous analyses in several ways. Almost all previous analyses that studied these new questions considered only the peculiar sterile oscillation pattern that gives the simplest physics: the sterile neutrino mixes with a specific flavour ($\\nu_{\\mu,\\tau}$ in solar analyses, $\\nu_\\tau$ in atmospheric ones) and has a large mass, $\\Delta m^2\\gg \\Delta m^2_{\\rm atm,sun}$. Dropping each one of these assumptions gives quite different physics. We explored the whole parameter space (three sterile mixing parameters and one sterile mass) producing precise results for six representative slices, that span the spectrum of the various possibilities. We considered one extra sterile neutrino with arbitrary mass $m_4$, and allowed it to mix with $$\\nu_e\\quad\\hbox{or}\\quad \\nu_\\mu\\quad\\hbox{or}\\quad\\nu_\\tau\\quad\\hbox{or}\\quad \\nu_1\\quad\\hbox{or}\\quad\\nu_2\\quad\\hbox{or}\\quad\\nu_3$$ where $\\nu_{1,2,3}$ are the mass eigenstates in absence of sterile mixing. The spectrum of active neutrinos is not yet fully known: rather than studying all possible cases we focussed on what we believe is most plausible case: we assumed normal hierarchy of active neutrinos (i.e. $m_1\\ll m_2\\approx m_3/6$) and that $\\theta_{13}$ is small enough that we can neglect its effects. If experiments will contradict these assumptions, it will be easy to update our results; most of them do not depend on these assumptions.  \\medskip   \\begin{figure}[p] $$\\hspace{-7mm}\\includegraphics[width=18cm]{all}$$ \\caption{\\label{fig:all}\\em {\\bf Summary of sterile neutrino effects}. The shaded region is excluded at $99\\%$ C.L.\\ (2 dof) by solar or atmospheric or reactor or short base-line experiments. We shaded as excluded also regions where sterile neutrinos suppress  the SN1987A $\\bar\\nu_e$ rate by more than $70\\%$. This rate is suppressed by more than $20\\%$ inside the {\\color{blue} dashed blue line}, that can be explored at the next SN explosion if it will be possible to understand the  collapse well enough. Within standard cosmology, the region above the {\\color{rossos} red continuous line} is disfavoured (maybe already excluded) by BBN and LSS. Plots in the various sections show how much each probe can be improved by future experiments.} \\end{figure}  We considered the most promising ways to probe the existence of eV-scale sterile neutrinos. Most probes are based on a careful study of natural sources of neutrinos (the universe, the sun, supernov\\ae, cosmic rays,...) which have their own peculiar capabilities and limitations. We studied how one extra sterile neutrino affects BBN (helium-4, deuterium), CMB, LSS and solar, atmospheric, reactor, beam experiments. The sensitivity of some of these probes is enhanced by MSW resonances~\\cite{MSW}. In cosmology, active $\\nu$ and $\\bar\\nu$ encounter a MSW resonance with sterile neutrinos lighter than active ones. Roughly the same happens to supernova $\\bar\\nu_e$ (that also experience less important MSW resonances in the opposite situation). On the contrary, solar $\\nu_e$ encounter a MSW resonance with sterile neutrinos heavier than active ones.   Fig.s\\fig{all} combines present constraints. Each probe is described in greater detail in its specific series of figures: fig.s\\fig{BBN} for cosmology, fig.s\\fig{sun} for solar experiments, fig.s\\fig{SN} for supernov\\ae, fig.s\\fig{SBL} for short base-line experiments, fig.s\\fig{atm} for atmospheric experiments. These figures also show the capabilities of some future experiments, that we now try to summarize in words. \\begin{itemize} \\item Compatibility with standard BBN constrains sterile oscillations occurred at temperatures $T\\circa{>} 0.1\\MeV$. It is very important to improve measurements of the {\\bf helium-4} primordial abundancy (that we parameterize in terms of an effective number of neutrinos $N_\\nu^{^4{\\rm He}}$, see eq.~(\\ref{sys:HeD})) until $N_\\nu^{^4{\\rm He}}=4$ will be safely tested. This requires overcoming `systematic' uncertainties. The helium-4 abundancy is sensitive to two different sterile effects: increase of the total neutrino density, and depletion of electron-neutrinos. The second effect makes $N_\\nu^{^4{\\rm He}}$ sensitive to sterile  oscillations down to $\\Delta m^2\\sim 10^{-8}\\eV^2$, while the first effect becomes negligible at $\\Delta m^2\\circa{<} 10^{-5}\\eV^2$.  However, if BBN were non-standard, a modified density of electron neutrinos could compensate the sterile corrections to $N_\\nu^{^4{\\rm He}}$: for example  helium-4 constraints on sterile oscillations can be evaded by allowing a  large neutrino asymmetry.    \\item For all these reasons it is important to measure a second BBN effect. The {\\bf deuterium} primordial abundancy is affected by milder systematic problems: in the future it might be possible to improve its measurement obtaining an uncertainty on the effective parameter $N_\\nu^{\\rm D}$ (precisely defined in eq.~(\\ref{sys:HeD}))) significantly below 1, possibly making deuterium the most significant BBN probe. We have computed the ranges of active/sterile oscillation parameters that significantly affect $N_\\nu^{\\rm D}$: it is less sensitive than helium-4 to $\\nu_e$ depletion and therefore to  values of $\\Delta m^2$ below $10^{-5}\\eV^2$  (fig.\\fig{BBN}).    \\item Future studies of {\\bf Cosmic Microwave Background} acoustic oscillations should allow to precisely measure the  total neutrino density $N_\\nu^{\\rm CMB}$ at recombination ($T\\sim \\eV$) with $\\pm0.2$ ({\\sc Planck}) or maybe $\\pm0.05$ ({\\sc CMBpol}) error~\\cite{NuCDM}. Neutrinos affect CMB in various ways; neutrino free-streaming offers a clean signature that allows to count neutrinos. Sterile neutrinos affect  $N_\\nu^{\\rm CMB}$ only if  $\\Delta m^2\\circa{>}10^{-5}\\eV^2$; in such a case $N_\\nu^{\\rm CMB}\\approx N_\\nu^{^4{\\rm He}}\\approx N_\\nu^{\\rm D}$. Therefore CMB will cover only a part of the region that BBN could probe with fully reliable measurements of the helium-4 or deuterium abundances.      \\item CMB also allows to probe $\\eV$-scale $\\nu$ masses. Smaller $\\nu$ masses can be probed by measuring how much galaxies are clustered, because neutrinos become non relativistic when the observable universe had a size comparable to present cluster of galaxies. Relativistic neutrinos freely move and tend to reduce the amount of clustering. Recently,  {\\bf Large Scale Structure} data (together with precise CMB measurements) gave a bound on the present energy density in neutrinos $\\Omega_\\nu < 0.0076$ at 95 \\% C.L.~\\cite{WMAP,boundMnu}, dominated by neutrino masses (rather than neutrino energy). With only active neutrinos this implies a significant bound on their mass, $m_\\nu < 0.23\\eV$~\\cite{WMAP,boundMnu}.  Sterile neutrinos contribute to $\\Omega_\\nu$ (a small non-thermal population of relatively heavy sterile neutrinos also modifies the way $\\Omega_\\nu$ manifests). For small active/sterile mixing LSS constrains sterile oscillations more strongly than BBN (fig.\\fig{BBN}).    \\item {\\bf Solar} $\\nu_e$ experiments have explored sterile oscillations not testable by BBN, thanks to two different effects. (1) MSW resonances make solar $\\nu_e$ sensitive to small active/sterile mixing and $\\Delta m^2\\circa{>}10^{-8}\\eV^2$. (2) With large mixing, solar $\\nu_e$ are sensitive down to $\\Delta m^2\\sim 10^{-12}\\eV^2$. Future experiments will explore new aspects of the solar neutrino anomaly, allowing to measure in a redundant way the active oscillation parameters or to discover a new anomaly. We emphasize one qualitative point. Due to  LMA oscillations, neutrinos exit from the sun as almost pure $\\nu_2$ at energies $E_\\nu  \\circa{>}  \\hbox{few MeV}$. Neutrinos with these energies have been precisely studied by SK and SNO, but are almost unaffected by sterile oscillations if they involve mostly $\\nu_1$. This could happen either when $\\nu_{\\rm s}$ mixes with $\\nu_1$, or when $\\nu_{\\rm s}$ experiences a level-crossing with $\\nu_1$. Therefore {\\em there is a whole class of sterile effects which manifest only at $E_\\nu \\circa{<}  \\hbox{\\rm few MeV}$} --- an energy range explored so far only by  Gallium experiments. Future precise measurement of solar $\\nu_e$ at sub-MeV energies will allow to significantly extend searches for active/sterile effects. Part of these extended region can be soon tested by Borexino, where a sterile neutrino can manifest as day/night variations, or as seasonal variations, or even by reducing the total rate.     \\item {\\bf Supernova} neutrinos will be good probes of sterile oscillations because have a different pattern of MSW resonances and a longer base-line than solar $\\nu_e$. Consequently supernova $\\bar\\nu_e$ are more sensitive than solar $\\nu_e$ in two main cases: (a) small $\\Delta m^2\\circa{>}10^{-18}\\eV^2$ with large $\\theta_{\\rm s}$; (b) $\\nu_{\\rm s}$ lighter than $\\nu_1$ with small mixing. Oscillations into one sterile neutrino can reduce the $\\bar\\nu_e$ rate by up to $80\\%$ (see fig.s\\fig{SN}) and, in a more restricted range of oscillation parameters, vary the average $\\bar\\nu_e$ energy by $30\\%$. SN1987A data agreed with expectations. Future SN experiments can perform quantitative test, but it is not clear how to deal with theoretical uncertainties. We also discussed other less promising astrophysical probes.   \\item {\\bf Atmospheric} experiments (SK, MACRO, K2K) indirectly exclude active/sterile oscillations with $\\Delta m^2\\circa{>}10^{-3\\div 4}\\eV^2$ and $\\tan^2\\theta_{\\rm s}\\circa{>}0.1\\div0.2$. Up to minor differences, this applies to all flavours (fig.\\fig{atm}). {\\bf Terrestrial experiments} that mainly probed disappearance of $\\bar\\nu_e$ and $\\nu_\\mu$ ({\\sc Chooz}, CDHS,\\ldots) exclude active/sterile  mixings with these flavours with $\\tan^2\\theta_{\\rm s}\\circa{>}0.03$ and $\\Delta m^2\\circa{>}10^{-3}\\eV^2$  (fig.\\fig{SBL}). Therefore future short-baseline experiments can search for sterile effects with smaller $\\theta_{\\rm s}$. Possible signals are $\\bar\\nu_e$ disappearance in reactor experiments, a deficit of NC events or $\\nu_\\tau$ appearance in beam experiments. Within standard cosmology these effects can be probed by CMB and BBN, which already disfavour them.   \\end{itemize} We listed present anomalies that can be interpreted as due to sterile neutrinos. None looks particularly significant, with the possible exception of the LSND anomaly, that can be due to oscillations of `$3+1$' neutrinos. Our precise study confirms that the extra sterile neutrino suggested by LSND thermalizes almost completely before BBN (in agreement with estimations of most previous analyses). In fig.\\fig{LSND} we plot, as function of the effective LSND oscillation parameters $\\theta_{\\rm LSND}$ and $\\Delta m^2_{\\rm LSND}$, the constraints on $N_\\nu^{^4{\\rm He}}\\simeq N_\\nu^{\\rm D}\\simeq N_\\nu^{\\rm CMB}$ (mainly from helium-4 data) and $\\Omega_\\nu$ (mainly from LSS data): in the relevant region our precise plot negligibly differs from previous estimates~\\cite{instant,2+2}. The 3+1 interpretation of LSND is not compatible with standard BBN, and gives a $\\Omega_\\nu$ which is only marginally compatible with standard cosmology.  \\bigskip   \\paragraph"
    },
    "0403/astro-ph0403093_arXiv.txt": {
        "abstract": "The excitation of gravity waves by penetrative convective plumes is investigated using 2D direct simulations of compressible convection. The oscillation field is measured by a new technique based on the projection of our simulation data onto the theoretical $g$-modes solutions of the associated linear eigenvalue problem. This allows us to determine both the excited modes and their corresponding amplitudes accurately. ",
        "introduction": "Two-dimensional simulations of compressible convection have shown that it is possible to excite internal gravity waves (IGW) in radiative zones of solar-type stars from the downward penetrating plumes (Hurlburt et al., \\citeyear{Hur86}, hereafter HTM86; Hurlburt et al., \\citeyear{Hur94}; Kiraga and Jahn, \\citeyear{Kir95}).  However, detecting IGW with confidence is challenging given the stochastic nature of the excitation mechanism.  We propose here a new detection method which allows us to measure rigorously both the spectrum and amplitude of excited $g$-modes. This method is first applied to the $g$-mode oscillations of an isothermal atmosphere and then, to IGW generated in a 2D-simulation of a convective zone embedded between two stable ones. ",
        "conclusions": ""
    },
    "0403/astro-ph0403570_arXiv.txt": {
        "abstract": "{ We address the question whether the properties of the observed latitude-time diagram of sunspot occurence (the butterfly diagram) provide evidence for the operation of a flux-transport dynamo, which explains the migration of the sunspot zones and the period of the solar cycle in terms of a deep equatorward meridional flow.  We show that the properties of the butterfly diagram are equally well reproduced by a conventional dynamo model with migrating dynamo waves, but without transport of magnetic flux by a flow. These properties seem to be generic for an oscillatory and migratory field of dipole parity and thus do not permit an observational distinction between different dynamo approaches. ",
        "introduction": "In a recent paper, \\citet[][henceforth referred to as HNWR]{Hathaway:etal:2003} presented an analysis of the latitude-time diagram of sunspot observations (commonly called butterfly diagram) and suggested that their results provide ``strong observational evidence that a deep meridional flow toward the equator is driving the sunspot cycle''. This refers to the so-called flux-transport dynamo models, which attribute the equatorward drift of the sunspot zone in the course of the 11-year solar activity cycle to the transport of toroidal magnetic flux towards low heliographic latitudes by an equatorward meridional flow near the bottom of the convection zone, thought to be the return flow of the observed poleward flow in the upper part of the convection zone and at the solar surface. Such physical transport of magnetic flux is not used in the more `traditional' type of dynamo models, which explain the equatorward drift by a latitudinally propagating dynamo wave and thus do not require a material flow \\citep[for a recent comprehensive review of solar dynamo theory, see][]{Ossendrijver:2003}. In this paper we show that the properties of the butterfly diagram analysed by HNWR (drift velocity of the sunspot zone as a function of latitude and its relation to cycle length and amplitude) are well consistent with a dynamo-wave model without meridional flow. ",
        "conclusions": "Results of a simple dynamo model show that the drift velocity of the sunspot zone as a function of latitude and its relation to cycle length and amplitude can be reproduced by a migrating dynamo wave.  This casts doubt upon the suggestion of \\citet{Hathaway:etal:2003} that these properties provide observational evidence for a flux-transport dynamo based upon an equatorward meridional flow in the deep convection zone and that this flow sets the cycle period. In fact, these properties of the butterfly diagram seem to be generic to a field of dipolar parity with equatorward drifting, opposite-polarity branches of toroidal field.  Our result that the butterfly diagram does not permit an observational distinction between dynamo-wave and flux-transport models does not lessen the appeal of the flux-transport dynamo concept.  Indeed, the poleward meridional flow in the outer parts of the convection zone is an observed fact and clearly mass conservation requires an equatorward return flow somewhere below. However, whether the properties of this flow meet the requirements of flux-transport dynamo models can only be clarified by helioseismic measurements of the meridional flow throughout the whole convection zone and over a time period of the order of the solar cycle \\citep[cf.][]{Haber:etal:2002}."
    },
    "0403/astro-ph0403600_arXiv.txt": {
        "abstract": "{ In this paper we estimate the initial He content in about 30\\% of the Galactic globular clusters (GGCs) from new star counts we have performed on the recently published HST snapshot database of Colour Magnitude Diagrams (Piotto et al. 2002). More specifically, we use the so-called $R$-parameter and estimate the He content from a theoretical calibration based on a recently updated set of stellar evolution models. We performed an accurate statistical analysis in order to assess whether GGCs show a statistically significant spread in their initial He abundances, and whether there is a correlation with the cluster metallicity. As in previous works on the subject, we do not find any significant dependence of the He abundance on the cluster metallicity; this provides an important constraint for models of Galaxy formation and evolution. Apart from GGCs with the bluest Horizontal Branch morphology, the observed spread in the individual helium abundances is statistically compatible with the individual errors. This means that either there is no intrinsic abundance spread among the GGCs, or that this is masked by the errors. In the latter case we have estimated a firm 1$\\sigma$ upper limit of 0.019 to the possible intrinsic spread. In case of the GGCs with the bluest Horizontal Branch morphology we detect a significant spread towards higher abundances inconsistent with the individual errors; this can be fully explained by additional effects not accounted for in our theoretical calibrations, which do not affect the abundances estimated for the clusters with redder Horizontal Branch morphology. In the hypothesis that the intrinsic dispersion on the individual He abundances is zero, taking into account the errors on the individual $R$-parameter estimates, as well as the uncertainties on the cluster metallicity scale and theoretical calibration, we have determined an initial He abundance mass fraction $Y_{\\rm GGC}=0.250\\pm0.006$. This value is in perfect agreement with current estimates based on Cosmic Microwave Background radiation analyses and cosmological nucleosynthesis computations. ",
        "introduction": "The determination of the initial He abundance of Galactic Globular Cluster (GGC) stars plays an important role in astrophysics, because of its wide-ranging implications.  First of all, due to the fact that GGC stars are the oldest objects in the Galaxy, their initial He abundance ($Y_{\\rm GGC}$, where $Y$ denotes the mass fraction of He) should mirror the primordial He abundance ($Y_p$) produced during the big bang nucleosynthesis. Secondly, the presence (or absence) of gradients in $Y_{\\rm GGC}$ as a function of the cluster metallicity provides information about Galactic nucleosynthesis and evolution. Thirdly, the existence of an abundance spread at constant metallicity would strenghten the case for He being the so-called 'second parameter' (beside metallicity) that determines the morphology of the Horizontal Branch (HB) in the GGC Colour-Magnitude-Diagrams (CMDs -- Sandage \\& Wildey~1967).  Empirical estimates of $Y_{\\rm GGC}$, are necessarily indirect, since He lines are not detectable in GGC star spectra, apart from the case of hot HB objects, whose atmospheres are however affected by gravitational settling and radiative levitation, which strongly alter their initial chemical stratification (see, e.g., Michaud et al.~1983). $Y_{\\rm GGC}$ estimates take advantage of results from stellar evolution theory, which show that the evolution of low mass Population~II stars is affected by the initial He content. The so-called $R$-parameter (Iben~1968; Buzzoni et al.~1983; Caputo et al.~1987), defined as the number ratio of HB to Red Giant Branch (RGB) stars brighter than the HB level ($R=N_{HB}/N_{RGB}$), is employed to determine $Y_{\\rm GGC}$. At a given metallicity, a higher initial He-content implies a brighter HB and, in turn, a lower value of $N_{RGB}$ ($N_{HB}$ is only slightly affected), with the consequent increase of $R$. Other parameters derived from stellar evolution can also be employed (see, e.g., the discussions in Sandquist~2000; Zoccali et al.~2000), but they are better suited to determine relative He abundances than absolute ones.  In Cassisi, Salaris \\& Irwin~(2003, hereafter Paper~I) we have shown how $R$-parameter determinations in two samples of GGCs (Sandquist~2000, 43 objects; Zoccali et al.~2000, 26 objects), coupled with a theoretical calibration obtained from updated stellar evolution models, provide a constant value of $Y_{\\rm GGC}$, with no trend with respect to [Fe/H]. The weighted average of the individual cluster abundances provides $Y_{\\rm GGC}$=0.244$\\pm$0.006 or $Y_{\\rm GGC}$=0.243$\\pm$0.006, the negligible difference being due to the choice of the observational sample. This value of $Y_{\\rm GGC}$ is in agreement -- within the errors -- with recent determinations of the cosmological baryonic matter density ($\\Omega_b$) from the Cosmic Microwave Background (CMB) power spectrum obtained by the BOOMERANG, DASI, MAXIMA and WMAP experiments (e.g. Pryke et al.~2002; \\\"Odman et al.~2003; Sievers et al.~2003; Spergel et al.~2003) which, coupled to BBN calculations (Burles, Nollett \\& Turner~2001), do consistently provide $Y_p$=0.248$\\pm$0.001. It is  important to recall that earlier studies (Sandquist~2000; Zoccali et al.~2000) that employed older generations of stellar models, found $Y_{\\rm GGC}$$\\sim$0.20, in complete disagreement with the CMB constraint and spectroscopic determinations from HII regions ($Y_p$=0.234$\\pm$0.002 according to Olive, Steigman \\& Skillman~1997; $Y_p$=0.244$\\pm$0.002 according to Izotov \\& Thuan~1998).  As for the abundance spread at a given [Fe/H], the results of Paper~I were contradictory, in the sense that $R$-parameter values from the Zoccali et al. (2000) sample show a spread that is completely consistent with the measurement errors, whereas the sample by Sandquist~(2000) shows an intrinsic additional 1$\\sigma$ spread in the He abundance of $\\sim$0.03.  In this paper we measured new $R$-parameter values for a sample of 57 GGCs from the database of HST GGC observations by Piotto et al.~(2002). We use the HST flight photometric system (without reddening corrections) instead of the Johnson photometric bands. This approach avoids potential subtle errors in the estimate of the visual magnitudes of the HB. In fact, the transformation to the Johnson system requires the knowledge of the cluster reddening, therefore the accuracy of the observed Johnson magnitudes depends on the accuracy of the cluster reddening estimate.  Our determinations of $R$-parameter values constitute the largest existing homogeneous database, encompassing about 30\\% of the total GGC population. It is therefore extraordinarily well suited to address in more detail the question of the intrinsic spread in the individual GGC He abundances, and the ratio of He enrichment with metallicity enrichment ($\\Delta Y/\\Delta Z$). With regard to the latter quantity, its evaluation is extremely important because of the connections with stellar yields and mechanism of formation of the GGCs, and the determination of their absolute ages. In addition, it is relevant to make a comparison with the value $\\Delta Y/\\Delta Z$ between $\\approx$ 6 and 12 (depending on the sample adopted, and with the assumption that the Oxygen abundance is a good tracer of total metallicity) estimated by Olive et al.~(1997) for extragalactic HII regions, and the value $\\sim$ 3 estimated by Izotov \\& Thuan~(2003). We will also address the hypothesis put forward by D'Antona et al.~(2002), which suggests the existence of an He enriched stellar component within individual clusters, due to chemical pollution by the ejecta of massive asymptotic giant branch stars; in particular, we will assess if this phenomenon might seriously affect our determination of $Y_{\\rm GGC}$.  In Sect.~2 we briefly describe the cluster sample and the theoretical models employed. Section~3 presents our results for the $R$-parameter measurements and the determination of $Y_{\\rm GGC}$. A summary and conclusions follow in Sect.~4. ",
        "conclusions": "In this paper we have presented a detailed study of the initial He abundance of the GGC system. We have used the most extensive and homogeneous database of $R$-parameter values to date, and recent updated stellar evolution models. As in Paper~I, we do not find any statistically significant correlation between the individual cluster He abundances and [Fe/H]. This suggests a very homogeneous value of $Y$ for the GGCs, with practically no He abundance evolution over the entire [Fe/H] spanned by the GGC system, at variance with results on the evolution of He in extragalactic HII regions, which measure a gradient of He with Oxygen (hence with [Fe/H]). This result on the evolution of He in GGCs strengthens similar conclusions reached in Paper~I, and it has to be taken into account in studies of Galaxy formation mechanisms. Note that Sandquist~(2000) also found a constant $Y$ for the GGC system using two independent indicators of relative He abundances, namely the $\\Delta$ parameter (difference in magnitude between HB and Main Sequence) and the RR Lyrae mass-luminosity exponent $A$.  As for the existence of an intrinsic spread of the cluster initial $Y$, we found that objects with $HB_{type}<$0.8 -- irrespective of their [Fe/H] -- do show a remarkably homogeneous and symmetric abundance distribution, whose spread is fully compatible with the measurement errors. One can therefore conclude that either the He abundance is the same for all clusters, without any intrinsic dispersion, or that an eventual dispersion is masked by the individual errors. In this case we obtain that our data allow a 0.019 upper limit to a possible 1$\\sigma$ intrinsic dispersion. On the other hand, clusters with blue HB colours ($HB_{type}\\geq$0.8) show a spread in their initial $Y$ incompatible with the individual error bars, skewed towards values higher than those for clusters with redder HB. We interpret this result as due to the increase of evolutionary timescales along the HB phase for progressively bluer HB stars, which causes an overestimate of $Y$ when using our calibration. This effect may be coupled to the possible presence of a subpopulation of He enhanced stars in at least some GGCs (see D'Antona et al.~2002), which also leads to an overestimate of $Y$ (at most by 0.01 -- 0.02) when determined from our calibration.  As already noticed, for clusters with $HB_{type}<$0.8 we obtain a distribution of $Y$ values very close to Gaussian and a 1$\\sigma$ dispersion consistent with the measurement errors. This occurrence goes against the possible existence of a sizable spread towards higher $Y$ values due to subpopulations of He enriched stars, as in blue HB clusters; in fact, in case redder HB clusters, the effect of these subpopulations is at most at the level of $\\sim$0.005, and probably smaller. It is also reasonable to assume that these subpopulations might not be present in all clusters, and the net effect on our estimates of $Y_{\\rm GGC}$ should therefore be negligible. An extreme possibility is that all clusters with $HB_{type}<$0.8 are affected in the same way by this enhancement, whose maximum effect would be to bias our $Y_{\\rm GGC}$ estimate by $\\sim$0.005 towards too high values.  We conclude by providing a best estimate for $Y_{\\rm GGC}$, including the effect of the still uncertain [Fe/H] scale and the sources of systematic errors discussed in Paper~I (i.e., the error on the $^{12}$C$(\\alpha,\\gamma)^{16}$O reaction rate and the method adopted to suppress the breathing pulses during the final stages of central He burning). We have considered as reference value the $Y_{\\rm GGC}$ abundance determined with the ZW84 [Fe/H] scale. Starting with this reference $Y_{\\rm GGC}$ we have generated a set of 100000 synthetic He-abundance values, by applying (through a Monte Carlo simulation) to the reference value a set of random and systematic errors, according to a given probability distribution. In particular, random errors have been modeled according to a Gaussian distribution with mean value equal to the reference one, and 1$\\sigma$ dispersion equal to the corresponding random error on $Y_{\\rm GGC}$. The systematic uncertainties due to the choice of the [Fe/H] scale (which causes a decrease of $Y_{\\rm GGC}$ by 0.005 with respect to the reference value), $^{12}$C$(\\alpha,\\gamma)^{16}$O reaction rate (variation by $\\pm$0.008), and breathing pulses suppression technique (increase by 0.003) have been modeled using a uniform distribution spanning the appropriate range.  The mean value for the final synthetic distribution of He abundances is $Y_{\\rm GGC}$=0.250$\\pm$0.006. An intrinsic dispersion with a firm 1$\\sigma$ upper limit of 0.019 around this value of $Y_{\\rm GGC}$ is a priori possible given the observational errors. This estimate of $Y_{\\rm GGC}$ is in good agreement with the primordial He abundance inferred from the CMB in conjunction with Big-Bang nucleosynthesis computations. Within the respective 1$\\sigma$ errors this value is also in agreement with the results from Paper~I."
    },
    "0403/astro-ph0403485_arXiv.txt": {
        "abstract": "Recent analyses of solar photospheric abundances suggest that the oxygen abundance in the solar atmosphere needs to be revised downwards. In this study we investigate the consequence of this revision on helioseismic analyses of the depth of the solar convection zone and the helium abundance in the solar envelope and find no significant effect. We also find that the revised abundances along with the current OPAL opacity tables are not consistent with seismic data. A significant upward revision of the opacity tables is required to make solar models with lower oxygen abundance consistent with seismic observations. ",
        "introduction": "Recent analyses of spectroscopic data using modern atmospheric models have suggested that the solar abundance of oxygen and other abundant elements needs to be revised downwards (Allende Prieto, Lambert \\& Asplund 2001, 2002; Asplund et al.~2004). Asplund et al.~(2004) claim that the oxygen abundance should be reduced by a factor of about 1.48 from the earlier estimates of Grevesse \\& Sauval (1998). The abundances of C, N, Ne and Ar are also reduced. The measured ratio of oxygen to hydrogen abundance, $[O/H]=8.66\\pm0.05$, is different from earlier estimates at approximately $3\\sigma$ level. As a result, the ratio (by mass) of heavy element to hydrogen abundance, $Z/X$, reduces from 0.023 to 0.0171, which causes the heavy element abundance in the solar envelope to reduce from $Z=0.017$ to 0.0126. This will cause the opacity of solar material to decrease, in turn reducing the depth of the convection zone (henceforth CZ) in solar models. Bahcall \\& Pinsonneault (2004) have constructed a solar model using these revised abundances to find that the depth of the CZ is indeed  reduced significantly as the CZ base is found to be at $r_b=0.726R_\\odot$, which is inconsistent with the seismic estimate of $(0.713\\pm0.001)R_\\odot$ (Basu \\& Antia 1997, henceforth BA97; Basu 1998). Asplund et al.~(2004) have argued that their estimate refers to the abundances near the surface, while because of diffusion the abundance could be much higher in the radiative interior. While this may be true, it should be noted that since the CZ is mixed on a rather small time-scale, the abundances at the surface and above  the base of the CZ should be the same.  It is these abundances, coupled with the corresponding opacity tables, that determine the depth of the CZ in a solar model. It may be possible to increase the hydrogen abundance in the convection zone by increasing diffusion, thereby increasing the convection zone depth, but in that case the helium content of convection zone may reduce below the seismically estimated value.  Since the position of the CZ base, $r_b$, has been determined very accurately through seismic analysis (\\jcd, Gough \\& Thompson 1991; BA97) it provides a constraint on the heavy element abundance or the opacities. In fact, BA97 have shown that the then accepted solar abundances along with the OPAL opacities (Iglesias \\& Rogers 1996) are consistent with the helioseismic data. Thus it is of interest to test how the reduction in abundance affects the conclusions. Before we do this, we examine whether a decrease in the abundances affects the seismic estimate of the CZ depth or the estimate of the helium abundance, $Y$. ",
        "conclusions": "We have investigated the effects that the  revision of the  abundance of oxygen and related elements in the solar atmosphere may have on seismic estimates of the solar helium abundance and the depth of the \\hbox{CZ}. We find neither of these estimates to be sensitive to variations in $Z$. We find that solar envelope models that have reduced abundances of oxygen and related elements do not have the correct density profile in the CZ despite having the seismically determined CZ depth and surface $X$. The density difference is about 10\\%, which is more than 6 times the estimated uncertainties in density. In order to get a seismically consistent solar model it is necessary to increase either the abundances of heavy elements or the computed opacities for a given abundance ratio. Even for a much smaller reduction in the oxygen abundance ($[O/H]=8.80$, the MIX2 abundances) it turns out that $Z/X$ needs to be increased from 0.0218 to 0.0228 to match the sound speed and density in the solar \\hbox{CZ}. The region in the $Z/X$--opacity plane for both mixtures  that is consistent with seismic constraints is shown in Fig.~3. The allowed region is not too sensitive to variations in the oxygen abundance, but  to match the seismic constraints either the opacity or the heavy element abundances, or both need to be increased. From seismic constraints it is not possible to decide between these possibilities.  We find that the recent estimate of the value of the oxygen abundance by Asplund et al.~(2004) along with the OPAL opacities are not consistent with seismic data. In fact, if current opacity tables are accepted, no significant reduction in the oxygen abundance from the values of Grevesse \\& Sauval (1998) is  favored by helioseismology. Models constructed with the new abundances either have incorrect CZ depth, or incorrect $X$, or incorrect density profile in the \\hbox{CZ}. Increasing the diffusion of heavy elements below the CZ base to compensate for a reduction in the atmospheric abundance does not work, the models fail to satisfy the $X$ constraint in the solar envelope, even though they may have the correct CZ depth. If the new values of abundances are confirmed, then opacity sources in the Sun will need to be re-investigated.  The intrinsic opacities will need to be increased to counteract the decrease in opacity due to the reduction in heavy element abundances. It could be argued that the abundances in the atmosphere where the spectral lines are formed is different from those in the CZ, due to some fractionation. This could resolve the problem we face now. However we will be faced with the more fundamental problem of having no way of determining the CZ abundances of the Sun."
    },
    "0403/astro-ph0403350_arXiv.txt": {
        "abstract": "We present the analysis of four 3.6~cm radio continuum archival observations of Orion obtained using the Very Large Array in its A-configuration, with $0\\rlap.{''}3$ angular resolution. The observations were made during the period 1994-1997. In a region of $4' \\times 4'$, we detect a total of 77 compact radio sources. Of the total of detected sources, 54 are detected in one or more of the individual observations and 36 of these show time variability (by more than 30\\%) between the observed epochs. A deep image made from averaging all data shows an additional 23 faint sources, in the range of 0.1 to 0.3 mJy. Of the total of 77 sources, 39 are new centimeter detections. However, only 9 of the 77 sources do not have a previously reported counterpart at near-infrared, optical, or X-ray wavelengths. In particular, we detect three faint sources in the OMC-1S region that may be related to the sources that power the multiple outflows that emanate from this part of the Orion nebula. ",
        "introduction": "The Orion Nebula (NGC 1976, M42) is the closest region of massive star formation, and it is associated with a variety of compact objects that include stars, nebulae, and masers, as well as infrared, optical and X-rays sources (Gaume et al. 1998,  Ali \\& Depoy 1995, O'Dell \\& Wen 1994, and Garmire et al. 2000). It is the best studied region of its class. A recent summary of the properties of Orion is given by O'Dell (2001).  The early identification by Laques \\& Vidal (1979) of six unresolved objects near $\\theta^{1}$ Ori C, determined to be high ionization emission line sources, and subsequent studies in the radio by Garay et al. (1987), Churchwell et al. (1987), and Felli et al. (1993a; 1993b), among others, suggested that possibly these objects were compact neutral clouds photoionized externally by $\\theta^{1}$ Ori C. The observations of these and similar sources in Orion with the HST by O'Dell et al. (1993), established their nature as externally ionized proto-planetary disks (proplyds). We refer the reader to Henney \\& Arthur (1998) for a detailed description of the propylds.  In this paper we present an analysis of radio continuum VLA archival observations of Orion, made at 3.6 cm in the highest angular resolution configuration A. One of the purposes of our analysis was to produce a deeper image that could reveal additional faint radio sources to those previously reported in the literature. Indeed, we detected 77 compact radio sources, 39 of which are new centimeter detections.  A region in Orion that has received considerable attention in recent years is OMC-1S (Keene, Hildebrand, \\& Whitcomb 1982; Bally, O'Dell \\& McCaughrean 2000), because it shows remarkable outflow activity. However, the identification of the exciting sources of these outflows remains unclear. We have used our deep, averaged image to search for continuum sources in this region. ",
        "conclusions": "We have presented the analysis of sensitive, high angular resolution ($0\\rlap.{''}3$) VLA observations at 3.6 cm toward the Orion region of recent star formation. Our main two conclusions are summarized below:  1) In a region of $4' \\times 4'$ around the Orion Trapezium we detect a total of 77 compact 3.6 cm sources, of which 39 are new centimeter detections. We briefly discuss the characteristics of these sources and of their counterparts at other bands.  2) In the OMC-1S region we detect three faint sources that may be related to the sources that power the multiple outflows that emanate from this region.  \\vspace{0.5cm}"
    },
    "0403/astro-ph0403399_arXiv.txt": {
        "abstract": "A critical comparison of estimates for the rates of hypernovae (HNe) and gamma-ray bursts (GRBs) is presented. Within the substantial uncertainties, the estimates are shown to be quite comparable and give a Galactic rate of $10^{-6}$\\,--\\,$10^{-5}\\,$yr$^{-1}$ for both events. These rates are several orders of magnitude lower than the rate of core-collapse supernovae, suggesting that the evolution leading to a HN/GRB requires special circumstances, very likely due to binary interactions. Various possible binary channels are discussed, and it is shown that these are generally compatible with the inferred rates. ",
        "introduction": "While it has now been established for more than 5 years that gamma-ray bursts (GRBs) are caused by some of the most energetic explosions in the Universe (van Paradijs, Kouveliotou, \\& Wijers 2000), no promising channel for their progenitors has been identified, a situation very much resembling that of the progenitors of Type Ia supernovae some 20 years ago. The firm and, unlike the previous case, unambiguous association of a GRB (GRB030329) with a hypernova, SN~2003dh (Hjorth et al.\\ 2003a; Stanek et al.\\ 2003), a highly energetic Type Ic supernova (Mazzali et al.\\ 2003)\\footnote{The term Hypernovae has been used for SNe with energies significantly larger (by about a factor of 10 or more) than the canonical explosion energy of $1\\,{\\rm foe} \\equiv 10^{51}\\,$ergs (Nomoto et al.\\ 2003).}, has confirmed that at least some long-duration GRBs are {\\em observationally} connected with the explosion of massive stars\\footnote{GRBs fall into two classes: short- and long-duration bursts. Presently, very little is known about the progenitors of short bursts from an observational point of view. It is quite possible that they are caused by a completely different physical mechanism, e.g., the merger of two compact objects (van Paradijs et al.\\ 2000). All inferences made in this paper exclusively apply to long bursts.}.  All hypernovae known to date belong to the class of Type Ic supernovae (SNe~Ic), of which they form a subset.  These are SNe that show neither hydrogen nor significant amounts of helium in their spectra. Their progenitors are believed to be either very massive single stars that lost their hydrogen and helium envelopes in a stellar wind or massive stars that lost their envelopes through the interaction with a companion (Wheeler \\& Levreault 1985; Uomoto 1986; Podsiadlowski, Joss, \\& Hsu 1992; Nomoto et al.\\ 1994).  Two out of the three nearby GRBs known to date are associated with hypernovae (GRB980425/SN~1998bw at $z = 0.008$ and GRB030329/SN~2003dh at $z = 0.17$). The third case, GRB031203, is heavily extinguished by dust, so a SN association cannot be firmly ruled out (Hjorth et al.\\ 2003b). This interesting coincidence raises the important question of the general connection between hypernovae and long-duration GRBs, and whether most, or perhaps even all GRBs are associated with/caused by hypernovae.  This paper addresses this question by providing a critical comparison of the rates of GRBs and hypernovae. As is shown in \\S~2, the rates of GRBs and hypernovae are comparable within the uncertainties, and appear to be a small fraction of the global SN rate. This has important implications for the nature of their progenitors, which is discussed in detail in \\S~3 and \\S~4. ",
        "conclusions": "The rates of hypernovae and GRB are quite comparable, suggesting that a large fraction (most?) of hypernovae also produce GRBs, at least in some direction.  Moreover, the rates are significantly smaller than the rates of core-collapse SNe (or even the fraction of SNe that produce black holes).  Furthermore, at least at the present cosmological epoch, special circumstances are required to produce HNe/GRBs.  However, numerous fundamental questions remain unanswered, and no fully self-consistent evolutionary model for the progenitors exists at this time.  As long as this is the case, it is not even clear whether or not the hypernova and the GRB occur concurrently. Does the hypernova occur first and trigger the GRB through the fallback of hypernova ejecta, as may be required in some models (Vietri \\& Stella 1998; Podsiadlowski et al.\\ 2002), or does the GRB occur and then trigger a SN-like event through the interaction of the relativistic jet with the envelope as in the collapsar model (Khokhlov et al.\\ 1999; MacFadyen \\& Woosley 1999; Jason, Woosley, \\& Hoffman 2003)?  Although HNe/GRBs appear to be relatively rare events at the present epoch, this need not be the case for the first generation of stars. Lower metallicity may lead to lower angular-momentum loss from massive stars, and the star-formation environment may be very different. It is even conceivable that, at an early epoch of galaxy formation, hypernovae could provide the missing energy to eject half the baryons from galaxies (Silk 2003).  Finally, another important question concerns the relationship between hypernovae and the class of SNe Ib/c, of which they are a subgroup. Presumably, many normal SNe Ib/c are caused by the collapse of the core of a massive star that lost its H-rich envelope through binary interaction (Wheeler \\& Levreault 1985; Podsiadlowski et al.\\ 1992; Nomoto et al.\\ 1994) forming a neutron star. So perhaps one important distinction between a hypernova and an ordinary supernova is whether a black hole or a neutron star is formed in the aftermath. However, not all black hole formation events can lead to a hypernova: if the minimum mass of a single star that leads to the formation of a black hole is as low as 20\\,--\\,25\\Ms\\ (Maeder 1992; Fryer \\& Kalogera 2001), this would overproduce hypernovae by a large factor (see Table 1).  A natural explanation for this dichotomy may lie in the fact that black holes can either form promptly on a dynamical timescale or on a much longer timescale by continued accretion through a disk phase or fallback. In particular, the disk accretion phase, which is the essential ingredient in collapsar models, requires a rapidly rotating core. In the case of prompt collapse, one would not necessarily expect a bright SN.  This would imply the existence of a class of (very?) dim SNe Ib/c for which there is, however, no observational evidence at present.  \\bigskip This work was in part supported by a European Research \\& Training Network (HPRN-CT-20002-00303) and a Royal Society visitor's grant to K.N. The work of P.P., P.M., and K.N. (the \"Tokyo Think Tank\") at the University of Tokyo has been supported by the 21st Century COE Program of MEXT, Japan."
    },
    "0403/astro-ph0403216_arXiv.txt": {
        "abstract": "{The XMM-Newton survey of the Coma cluster of galaxies covers an area of 1.86 square degrees with a mosaic of 16 pointings and has a total useful integration time of 400 ksec. Detected X-ray sources with extent less than $10^{\\prime\\prime}$ were correlated with cataloged galaxies in the Coma cluster region. The redshift information, which is abundant in this region of the sky, allowed us to separate cluster members from background and foreground galaxies. For the background sources, we recover a typical $LogN-LogS$ in the flux range $10^{-15}-10^{-13}$ ergs s$^{-1}$ cm$^{-2}$ in the 0.5--2.0 keV band. The X-ray emission from the cluster galaxies exhibits X-ray colors typical of thermal emission. The luminosities of Coma galaxies lie in the $10^{39}-10^{41}$ ergs/s interval in the 0.5-2.0 keV band. The luminosity function of Coma galaxies reveals that their X-ray activity is suppressed with respect to the field by a factor of 5.6, indicating a lower level of X-ray emission for a given stellar mass. ",
        "introduction": "Studies of X-ray emission from normal galaxies can be used to reveal signs of recent star-formation activity (Grimm et al. 2003), the presence of the hot gas, filling the potential well of the galaxy and its immediate surroundings (Forman, Jones, Tucker 1985), and the population of discrete X-ray galactic sources (e.g. Pietsch et al. 1994, Irwin et al. 2003). With the advent of high-spatial resolution observations, such studies become feasible also for nearby clusters of galaxies, with the advantage that many galaxies are observed simultaneously, all at essentially the same distance from the observer. Such observations probe the effect of the cluster environment on properties of galaxies, in particular ram-pressure stripping and compression.  In this paper we present the properties of point sources in the XMM-Newton mosaic of the Coma cluster of galaxies, which covers most of the cluster virial radius. The advantages of studying the Coma cluster are the availability of deep optical catalogs, plus the results of many studies at various wavelengths, providing a unique opportunity for the identification and in-depth study of the X-ray characteristics of cluster galaxies. We use the catalog of Godwin, Metcalfe, Peach (1983, hereafter GMP) to associate X-ray sources with galaxies and employ the GOLDmine\\footnote{http://goldmine.mib.infn.it/} (Gavazzi et al, 2003) database to provide a wealth of additional information about the galaxies.  The paper is organized as follows, in \\S\\ref{s:data} we describe our analysis of XMM-Newton data. The X-ray emission of Coma galaxies is analyzed in \\S\\ref{s:glx}. In \\S\\ref{s:lf} we derive the X-ray luminosity function of the Coma cluster galaxies. We adopt $H_{\\rm o}=70$~km~s$^{-1}$~Mpc$^{-1}$ and $\\Omega=1$, with a resulting $D_{L}=100$ Mpc. One degree corresponds to 1.67 Mpc.   \\begin{figure*} \\includegraphics[width=17.cm]{0765fig1.eps}  \\figcaption{DSS-2 B band image of the core of the Coma cluster (grey scale), overlaid with EPIC-pn contours of the small-scale X-ray emission in the 0.5-2.0 keV band.  Galaxies are labeled by their GMP catalog number. The image is centered on 194.913, 27.958 (J2000) and is $36^\\prime \\times 26^\\prime$. North is up east is to the left. \\label{f:imh} } \\end{figure*}  \\begin{figure*} \\includegraphics[width=17.cm]{0765fig2.eps}  \\figcaption{Same as Fig.\\ref{f:imh} for the north-east region of the Coma cluster. The image is centered at 195.240, 28.095 (J2000) and is $37^\\prime \\times 39^\\prime$. North is up east is to the left \\label{f:imhne} } \\end{figure*} \\begin{figure*} \\includegraphics[width=17.cm]{0765fig3.eps}  \\figcaption{Same as Fig.\\ref{f:imh} for the north-west region of the Coma cluster. The image is centered at 194.499, 28.134 (J2000) and is $39^\\prime \\times 44^\\prime$. \\label{f:imhnw} } \\end{figure*} \\begin{figure*} \\includegraphics[width=17.cm]{0765fig4.eps}  \\figcaption{Same as Fig.\\ref{f:imh} for the south-east region of the Coma cluster. The image is centered at 195.186, 27.409 (J2000) and is $43^\\prime \\times 40^\\prime$. \\label{f:imhse} } \\end{figure*} \\begin{figure*} \\includegraphics[width=17.cm]{0765fig5.eps}  \\figcaption{Same as Fig.\\ref{f:imh} for the south-west region of the Coma cluster. There is one arcminute gap between this image and Fig.\\ref{f:imhnw}. The image is centered at 194.401, 27.440 (J2000) and is $35^\\prime \\times 34^\\prime$. \\label{f:imhsw} } \\end{figure*}  \\begin{figure*} \\includegraphics[height=18.cm,angle=-90]{0765fig6.eps}  \\figcaption{DSS-2 B band image of the core of the Coma cluster center, overlaid with the fluctuations in the entropy map of Coma cluster gas, as observed by pn.  Galaxies are labeled by their GMP catalog number. The image is centered at 194.958, 27.939 (J2000) and is $40^\\prime \\times 27^\\prime$. North is up east is to the left \\label{f:imhent} } \\end{figure*}    \\begin{figure*} \\includegraphics[width=8.2cm]{0765fig7a.eps}\\hfill\\includegraphics[width=8.2cm]{0765fig7b.eps}  \\figcaption{{\\it Left panel.}  Sensitivity map for source detection on the $8^{\\prime\\prime}$(radius) wavelet scale. Contours enclose two areas of similar sensitivity, $10^{-14}$ ergs s$^{-1}$ cm$^{-2}$, which includes the entire image, and $2\\times10^{-15}$ ergs s$^{-1}$ cm$^{-2}$. {\\it Right panel.} Surveyed area vs. source flux in the $0.5-2.0$ keV band. Dashed lines show at which flux the area reaches 10\\% and 90\\% of the total area surveyed of 1.86 square degrees. \\label{f:area} } \\end{figure*} ",
        "conclusions": ""
    },
    "0403/astro-ph0403202_arXiv.txt": {
        "abstract": "Several authors have recently explored the idea that physical constants such as $c$ and $G$ might vary over time and have formulated theories describing this variation that can address a range of cosmological problems. Such work typically invokes a generic parameterization which assumes a power-law variation with the expansion scale factor, $R$. This work offers alternative, physically motivated definitions for the parameters $c$, $G$, and $h$ based on the Machian premise that these dimensional quantities reflect global dynamics of the expansion geometry. Together with a postulated conservation law and equations of motion, the implications of this theory for Friedmann models are examined, and found to yield several interesting conclusions including: (1) natural solutions to the horizon, flatness, and lambda problems, (2) the prediction of a flat, $\\Omega_0 = 1$ universe, (3) different forms for some cosmological scaling laws, (4) an apparent fit to observations of Type Ia supernovae without invoking a cosmological constant, (5) equivalence between our Universe and a black hole and apparent consistency of the model with the Holographic Principle, and (6) potentially testable predictions for the time variation of physical parameters, including values for $\\dot{c_0}$ and $\\dot{h_0}$ that are small but non-zero today and a value for $\\dot{G}$ that was negative and nonzero during radiation domination and decayed to effectively zero upon the epoch of matter domination. While this work does not attempt to provide the complete theoretical foundation that must ultimately underlie any theory that could naturally marry traditional physics with the notion of time-varying physical parameters, it is written in the hope that it might stimulate further progress towards this end. ",
        "introduction": "For over three centuries, theories of physics have been characterized by the existence of physical constants.  Despite their ubiquity, however, the significance of these dimensional parameters remains relatively unexplored. While some, such as the Rydberg, have historically been revealed to be comprised of other more fundamental quantities, others such as the gravitational constant $G$, the speed of light $c$, and Planck's constant $h$, are of seemingly primary significance. If precedence suggests that constants might be viewed as placeholders for as yet undiscovered physics in our theories, then one of the challenges of physics is surely to reveal the physical meaning of these parameters. Indeed, it might be generally supposed that as physics progresses, fewer constants will be required as more fundamental theories are put forth--that a true ``theory of everything'' might contain no such quantities, explaining nature from first principles alone. In practice, of course, the Universe sometimes reveals complexities faster than theorists can explain them, and new constants are born. One such example is the cosmological constant $\\Lambda$, introduced (and later bemoaned) by Einstein in an attempt to arrive at a steady-state characterization of the Universe\\footnote{A $\\Lambda$-dominated model was later shown to be, in fact, unstable.}, and which in light of recent observations is now nearly universally accepted (Reiss \\etal 1998; Perlmutter \\etal 1999; Knop \\etal 2003; Bennett \\etal 2003; Tegmark \\etal 2003).  Several papers in recent years have explored the significance of physical ``constants'' (henceforth termed physical parameters) by assuming an ad hoc form for their possible variation over time and investigating the corresponding implications. For example, Albrecht \\& Magueijo (1999) and Barrow (1999) show that simple varying speed of light (VSL) theories offer an alternative to inflation for solving several of the problems associated with the Standard Big Bang (SBB) cosmology, such as the horizon, flatness, and lambda problems. Barrow (2003) demonstrates, however, that these simple models may face challenges, for example, in solving the isotropy problem and generating constant-curvature fluctuations, clearly highlighting the need for additional work to make such models viable.  The present work seeks to explore the significance of physical parameters from the Machian standpoint that these dimensional quantities reflect global characteristics of the Universe. From this premise, we will formulate two simple postulates relating $c$ and $G$ to the dynamics of the expansion geometry, offering a physical motivation for the meaning of these parameters.\\footnote{If physical parameters reflect global properties of the Universe, it is only natural that they {\\em should} vary in time in accordance with cosmological dynamics, i.e., the expansion. As there are obviously no physical constants in mathematics, one may suppose that it dimensionless, rather than dimensional, quantities that are in fact invariant. Indeed, since these assumptions will be shown to necessitate the existence of a preferred cosmological frame, it may be philosophically more reasonable to suppose that the distance traveled by a photon per unit time in such a frame frame might vary, than that the number 3 should evolve into 1 or 59$\\sqrt{\\pi}$. As a dimensionless quantity, the fine-structure parameter $\\alpha$ is assumed to be constant in this model.} A definition for $h$, which will also be linked to the expansion geometry, will follow from these postulates. Together with two additional postulates (a conservation law and equations of motion), we begin construction of a theory which appears to have several interesting features. A limited subset of these is presented here, including: (1) natural solutions to the horizon, flatness, lambda, and potentially other SBB problems, (2) the prediction of a flat, $\\Omega_0 = 1$ universe, (3) different forms for some cosmological scaling laws, (4) an apparent fit to observations of Type Ia supernovae without invoking a cosmological constant, (5) equivalence between our Universe and a black hole and apparent consistency of the model with the Holographic Principle, and (6) potentially testable predictions for the time variation of physical parameters, including values for $\\dot{c_0}$ and $\\dot{h_0}$ that are small but non-zero today and a value for $\\dot{G}$ that was negative and nonzero during radiation domination and decayed to effectively zero upon the epoch of matter domination. Note that in the context of this theory, any predicted time variation (or lack thereof) in the physical parameters arises naturally from their physical definitions, rather than from an assumed parameterization.  Section \\ref{section1} of this paper puts forth the postulates for the physical definitions of $c(t)$ and $G(t)$ and two additional postulates: a conservation equation and equations of motion. Section \\ref{section2} explores some of the resulting cosmological implications in Friedmann models, including specification of $G(t)$, solutions to SBB problems, specification of the expansion geometry, modified scaling laws, calculation of the luminosity distance, and application to Type Ia SN observations. Section \\ref{section3} investigates the quantitative behavior of $c(t)$ and $G(t)$ and draws the mathematical connections between our Universe and a black hole. Section \\ref{section4} puts forth a corollary for the physical definition of $h(t)$ and further develops the notion of ``black hole equivalence'' via thermodynamic analogy and apparent consistency with the Holographic Principle. Additional thoughts and a discussion on the resolution of this model with Special and General Relativity are presented in the Conclusion.  It is clear that any theory which motivates the existence of physical parameters and/or introduces their time variation must contend not only with the classical SBB problems and with the strong observational limits that have been measured (Rich 2003, Magueijo 2003), but ultimately must produce a compelling model of physics that expands our understanding of natural phenomena, rather than simply interesting heuristic exercises. Like most advances, however, the development of such a theory is likely to occur in small steps. The present exercise is merely an attempt to take one of those steps in the hopes that it might stimulate further discussion. It will focus on deriving a set of conclusions based on a series of postulates, but does not attempt to provide a complete theory nor to explore the entire range of implications. ",
        "conclusions": "This work has put forth several radical postulates on the definitions and variations of physical parameters, and drawn a number of conclusions which are different from those of standard physics. As mentioned in the Introduction, it has not presented a complete theoretical framework in which these ideas can be fully embedded. However, it is hoped that a viable path for further inquiry has been laid out for both theoretical and experimental investigation. In this concluding section we highlight a few starting points for such an inquiry.  From a theoretical perspective, there is nothing which prevents the fundamental ``constants'' from evolving. The VPP model is based on the Machian premise that physical parameters reflect global characteristics of the Universe, and are therefore {\\em expected} to evolve in accordance with the expansion. In practice, there are of course observational constraints on the variation of these quantities. To the degree that ``standard'' physics has arisen from laboratory experimentation conducted over a cosmologically infinitesimal timeframe, it can be viewed simply as the specialized case where $G(t) \\longrightarrow G_0$, $c(t) \\longrightarrow c_0$, $h(t) \\longrightarrow h_0$, etc., provided that $G(t)$, $c(t)$, and $h(t)$ do not violate experimental constraints. The present model is only one in a long history of models that have been proposed under this rationale.  At first glance, the VPP model might appear to be at odds with special or general relativity, but in fact this may not be the case. Consider special relativity (SR) which is based on the concept of inertial reference frames and follows from the postulates that: (1) the laws of physics are the same in all inertial reference frames, and (2) the speed of light as measured by any observer in an inertial frame is constant at $c$. The VPP model, however, posits the existence of a preferred frame in which we and everything else exists, so how can these be reconciled? Consider a classical particle with arbitrarily small mass moving at a speed $v$ that is arbitrarily close to $c_0$. The VPP model states that since this particle is moving at $c_0$ in the preferred frame, an observer in an ``inertial'' SR frame, will also measure a speed of $c_0$. Apparently, what SR calls an inertial frame is in the VPP model any frame at time $t_*$ in which the speed of a massless particle is measured to be $c(t_*)$. Under this equivalence, the first postulate of SR still holds and the second postulate is then merely a direct consequence of the first. SR, however, also requires local and global energy conservation, whereas the VPP model violates energy conservation stating that $E(t) \\propto c(t)^2 \\propto 1/R(t)$. Quantifying this violation, we find that for a system of {\\em any} equivalent mass \\begin{equation} \\frac{\\dot{E}}{E} = 2\\frac{\\dot{c}}{c} = -2q\\frac{c}{R} \\end{equation} which implies a present-day percent variation in $E$ of order $10^{-18}$ s$^{-1}$, or nearly one part in a billion per decade, using the values of $c_0$ and $R_0$. Thus, the predicted degree of non-conservation may not be at odds with the predictions of SR in terrestrial laboratory experiments.  The VPP model is of course assumed to be consistent with General Relativity (GR) by virtue of Postulate 4. It is hoped that a suitable generalization of GR can prove this. GR of course treats all reference frames on equal footing, but only defines reference frames in a local sense, and thus does not have a global energy conservation law. The results of the SBB model are derived by assuming a global energy conservation law for the Universe, in combination with the Einstein equations. The VPP model instead assumes, via Postulate 3, a global mass conservation law, which has been shown in Section \\ref{section2} to be consistent with the assumed forms of the physical parameters. It is noted that in principle it might be possible to formulate an alternative, but equivalent, VPP model in which energy is conserved and $M=M(t)$. If this is possible, however, such a model would seem to lack the same physical intuition and would not appear to imbue the physical parameters with the same geometric significance as in the current model.  As a final comment, it is noted that the statement $c = \\dot{R}$ suggests the possibility that a generalized theory consistent with the results of the VPP model might somehow couple massless particles directly to the cosmic expansion, so that they are in effect swept along geodesics by the expansion. By contrast, massive particles would be only weakly coupled to the expansion--governed on small scales by local forces that dominate over the expansion, and only at large scales by the expansion dynamics.  From a theoretical standpoint, the VPP model appears attractive in that it predicts $\\Omega=1$ for all time, explains a range of cosmological problems without invoking inflation, appears to fit Type Ia supernovae observations without invoking a cosmological constant, and gives rise to a series of other results that suggest the possibility of a deeper theoretical significance. Experimentally, the VPP model appears to offer several potentially testable predictions, some of which have been described in this paper. Whether this model proves to be correct, incomplete, or wrong, it is hoped this work will stimulate further discussion that may lead towards a greater understanding.  \\bigskip The author would like to thank E. Blackman for helpful discussions and comments."
    },
    "0403/astro-ph0403034_arXiv.txt": {
        "abstract": "The detection and characterization of accretion processes in the disks surrounding young stars may be directly relevant to studies of planet formation.  Especially the study of systems with very low accretion rates ($<<$ 10$^{-10}$ M$_\\odot$~yr$^{-1}$) is important, since at those rates radial mixing becomes inefficient and disk material will have to be dissipated into larger bodies at its present location.  In these proceedings, we compare the different methods of tracing accretion onto Herbig Ae/Be stars and conclude that high-resolution infrared spectroscopy is currently the only reliable method that offers the required sensitivity to shed light on this problem. ",
        "introduction": "The mass accretion rate is thought to be a key parameter in the evolution of young stellar objects (YSOs).  As a function of time, the accretion rate traces the build-up of material onto the young star and severely affects the evolution of the disk itself.  Local disk structure is affected by the rate of mass flow, which in turn is determined by the rate of gravitational energy release. Several studies (e.g. Bertout et al. 1988; Hartigan et al. 1995; Hartmann et al, 1998) have shown that accretion rates around typical T Tauri stars stars are low: of the order 10$^{-8}$--10$^{-10}$ M$_\\odot$~yr$^{-1}$.  Compared to typical disk-masses of $<$ 0.1~M$_\\odot$, and disk lifetimes of $<$ 10 Myr, these low accretion rates imply that at the evolutionary stage of a typical T Tauri star, disk accretion is no longer important for building up the mass of the central star.  Furthermore, at these accretion rates, the heating of the disk is dominated by reprocessing of light from the central star: the disk is passive.  However, there are several reasons why it is especially important to study accretion in YSOs with the lowest accretion rates: (1) if the accretion rate is observed to be lower than $\\sim$ 10$^{-7}$ M$_\\odot$~yr$^{-1}$ (the disk mass (typically $<$ 0.1~M$_\\odot$)/the disk dissipation timescale ($<$ 10 Myr)), the disk cannot disappear due to accretion onto the central star(s).  The conclusion that a significant mass fraction of the disk may be coagulated into larger bodies, such as comets or full-fledged planets, seems justified. (2) recent studies of infrared emission from dust in the disks surrounding Herbig Ae/Be stars (e.g. Bouwman et al. 2003) have shown that crystalline silicates are found at temperatures of a few hundred K; temperatures that are much lower than their sublimation temperature of $\\sim$ 1500~K.  Therefore they cannot have formed at their present location.  It is currently believed that this provides strong evidence for the importance of radial mixing in Herbig stars.  However, radial mixing can only occur with the required efficiency in the presence of mass accretion rates larger than a few times 10$^{-9}$ M$_\\odot$~yr$^{-1}$.  If accretion rates are proven to be lower than this limit, radial mixing becomes too inefficient to affect disk structure as a whole: not only must the mass in these disks be dissipated into larger bodies, but they also have to form from the radial distribution of material as it is observed  In these {\\it proceedings} we will discuss the different methods of tracing accretion onto Young Stellar Objects and conclude that high-resolution infrared spectroscopy may be the most sensitive method in our current arsenal of determining accretion rates. ",
        "conclusions": ""
    },
    "0403/astro-ph0403344_arXiv.txt": {
        "abstract": "The linear anisotropies in the temperature of the cosmic microwave background (CMB) radiation and its polarization provide a clean picture of fluctuations in the universe some $370\\,\\mathrm{kyr}$ after the big bang. Simple physics connects these fluctuations with those present in the ultra-high-energy universe, and this makes the CMB anisotropies a powerful tool for constraining the fundamental physics that was responsible for the generation of structure. Late-time effects also leave their mark, making the CMB temperature and polarization useful probes of dark energy and the astrophysics of reionization. In this review we discuss the simple physics that processes primordial perturbations into the linear temperature and polarization anisotropies. We also describe the role of the CMB in constraining cosmological parameters, and review some of the highlights of the science extracted from recent observations and the implications of this for fundamental physics. ",
        "introduction": "\\label{adc:sec:intro}  The cosmic microwave background (CMB) radiation has played an essential role in shaping our current understanding of the large-scale properties of the universe. The discovery of this radiation in 1965 by Penzias and Wilson~\\cite{adc:pw65}, and its subsequent interpretation as the relic radiation from a hot, dense phase of the universe~\\cite{adc:dicke65} put the hot big bang model on a firm observational footing. The prediction of angular variations in the temperature of the radiation, due to the propagation of photons through an inhomogeneous universe, followed shortly after~\\cite{adc:sw67}, but it was not until 1992 that these were finally detected by the Differential Microwave Radiometers (DMR) experiment on the Cosmic Background Explorer (COBE) satellite~\\cite{adc:smoot92}. The fractional temperature anisotropies are at the level of $10^{-5}$, consistent with structure formation in cold dark matter (CDM) models~\\cite{adc:peebles82,adc:bond84}, but much smaller than earlier predictions for baryon-dominated universes~\\cite{adc:sw67,adc:peebles70}. Another experiment on COBE, the Far InfraRed Absolute Spectrophotometer (FIRAS), spectacularly confirmed the black-body spectrum of the CMB and determined the (isotropic) temperature to be 2.725\\,K~\\cite{adc:mather94,adc:mather99}.  In the period since COBE, many experiments have mapped the CMB anisotropies on a range of angular scales from degrees to arcminutes (see \\cite{adc:bond03} for a recent review), culminating in the first-year release of all-sky data from the Wilkinson Microwave Anisotropy Probe (WMAP) satellite in February 2003~\\cite{adc:bennett03}. The observed modulation in the amplitude of the anisotropies with angular scale is fully consistent with predictions based on coherent, acoustic oscillations~\\cite{adc:peebles70}, derived from gravitational instability of initially adiabatic density perturbations in a universe with nearly-flat spatial sections. The amplitude and scale of these acoustic features has allowed many of the key cosmological parameters to be determined with unprecedented precision~\\cite{adc:spergel03}, and a strong concordance with other cosmological probes has emerged.  In this review we describe the essential physics of the temperature anisotropies of the CMB, and its recently-detected polarization~\\cite{adc:kovac02}, and discuss how these are used to constrain cosmological models. For reviews that are similar in spirit, but from the pre-WMAP era see e.g.~\\cite{adc:hu02,adc:hu03}. We begin in Sect.~\\ref{adc:sec:fundamentals} with the fundamentals of CMB physics, presenting the kinetic theory of the CMB in an inhomogeneous universe, and the various physical mechanisms that process initial fluctuations in the distribution of matter and spacetime geometry into temperature anisotropies. Section~\\ref{adc:sec:params} discusses the effect of cosmological parameters on the power spectrum of the temperature anisotropies, and the limits to parameter determination from the CMB alone. The physics of CMB polarization is reviewed in Sect.~\\ref{adc:sec:pol}, and the additional information that polarization brings over temperature anisotropies alone is considered. Finally, in Sect.~\\ref{adc:sec:highlights} we describe some of the scientific highlights that have emerged from recent CMB observations, including the detection of CMB polarization, implications for inflation, and the direct signature of dark-energy through correlations between the large-scale anisotropies and tracers of the mass distribution in the local universe. Throughout, we illustrate our discussion with computations based on $\\Lambda$CDM cosmologies, with baryon density $\\Omega_\\mathrm{b}h^2 = 0.023$ and cold dark matter density $\\Omega_\\mathrm{c} h^2 = 0.111$. For flat models we take the dark-energy density parameter to be $\\Omega_\\Lambda = 0.75$ giving a Hubble parameter $H_0 = 73\\, \\mathrm{km\\,s}^{-1}\\,\\mathrm{Mpc}^{-1}$. We adopt units with $c=1$ throughout, and use a spacetime metric signature $+---$. ",
        "conclusions": "\\label{adc:sec:summary}  The linear anisotropies of the cosmic microwave background have been studied theoretically for over three decades. The physics, which is now well understood, employs linearised radiative transfer, general relativity, and hydrodynamics to describe the propagation of CMB photons and the evolution of the fluid constituents in a perturbed Friedmann-Robertson-Walker universe. A number of bold predictions have emerged from this theoretical activity, most notably the existence of acoustic peaks in the anisotropy power spectrum due to oscillations in the photon--baryon plasma prior to recombination. Observers have risen to the challenge of verifying these predictions, and their detection is proceeding at a staggering rate. The large-scale Sachs--Wolfe effect, acoustic peak structure, damping tail, late-time integrated Sachs--Wolfe effect, polarization and reionization signature have all been detected, and the first three have been measured in considerable detail. Already, the size and scale of these effects is allowing cosmological models to be constrained with unprecedented precision. The results are beautifully consistent with almost-scale-invariant adiabatic initial conditions evolving passively in a spatially flat, $\\Lambda$CDM universe.  Much work still remains to be done to exploit fully the information contained in the CMB anisotropies. The Planck satellite, due for launch in 2007, should provide definitive mapping of the linear CMB anisotropies, and a cosmic-variance limited measurement of the power spectrum up to multipoles $l\\sim 2000$. This dataset will be invaluable in assessing many of the issues hinted at in the first-year release of WMAP data, such as the apparent lack of power on large scales and possible violations of rotational (statistical) invariance. Prior to Planck, a number of ground-based programmes should shed further light on the issue of whether departures from a power-law primordial spectrum are required on cosmological scales, and the implications of this for slow-roll inflation. In addition, these small-scale observations will start to explore the rich science of secondary anisotropies, due to e.g.\\ scattering in hot clusters~\\cite{adc:sz72} or bulk flows modulated by variations in the electron density in the reionized universe~\\cite{adc:sz80,adc:ostriker86}, and the weak lensing effect of large-scale structure~\\cite{adc:blanchard87}.  Detections of CMB polarization are in their infancy, but we can expect rapid progress on this front too. Accurate measurements of the power spectra of $E$-mode polarization, and its correlation with the temperature anisotropies, can be expected from a number of ground and balloon-borne experiments, as well as from Planck. The ultimate goal for CMB polarimetry is to detect the $B$-mode signal predicted from gravitational waves. This would give a direct measure of the energy scale of inflation, and, when combined with measurements of the spectrum density perturbations, place tight constraints on the dynamics of inflation. Plans are already being made for a new generation of polarimeters with the large numbers of detectors and exquisite control of instrument systematics needed to detect the gravity-wave signal if the energy scale of inflation is around $10^{16}\\,\\mathrm{GeV}$. Ultimately, confusion due to imperfect subtraction of astrophysical foregrounds and the effects of weak lensing on the polarization limit will limit the energy scales that we can probe with CMB polarization; see~\\cite{adc:hirata03} and references therein."
    },
    "0403/astro-ph0403172_arXiv.txt": {
        "abstract": "We study secular stability against a quasi-radial oscillation for rigidly rotating stars with soft equations of state in general relativity. The polytropic equations of state with polytropic index $n$ between 3 and 3.05 are adopted for modeling the rotating stars. The stability is determined in terms of the turning-point method. It is found that (i) for $n \\agt 3.04$, all the rigidly rotating stars are unstable against the quasi-radial oscillation and (ii) for $n \\agt 3.01$, the nondimensional angular momentum parameter $q \\equiv cJ/GM^2$ (where $J$, $M$, $G$, and $c$ denote the angular momentum, the gravitational mass, the gravitational constant, and the speed of light, respectively) for all marginally stable rotating stars is larger than unity. A semi-analytic calculation is also performed, and good agreement with the numerical results is confirmed. The final outcome after axisymmetric gravitational collapse of rigidly rotating and marginally stable massive stars with $q > 1$ is predicted, assuming that the rest-mass distribution as a function of the specific angular momentum is preserved and that the pressure never halt the collapse. It is found that even for $1 < q \\alt 2.5$, a black hole may be formed as a result of the collapse, but for $q \\agt 2.5$, the significant angular momentum will prevent the direct formation of a black hole. ",
        "introduction": "At the final stage of their evolution, massive stars with initial mass $\\agt 8M_{\\odot}$ form an iron core. Because of high temperature and density, the photodissociation of the iron to lighter elements proceeds rapidly (Shapiro \\& Teukolsky 1983). In the photodissociation, the thermal pressure is consumed, decreasing the adiabatic index $\\Gamma$ below $4/3$. Consequently, the core is destabilized, resulting in collapse to a neutron star or a black hole.  Recent numerical simulations of the collapse of primordial molecular clouds have suggested that the first generation of stars contains many massive members of mass $\\sim 100$--$1000M_{\\odot}$ (Bromm et al. 1999; Abel et al. 2000; Nakamura \\& Umemura 2001). Such massive stars form large helium cores that reach carbon ignition. It is known that after helium burning, cores of such stars encounter the electron-positron pair creation instability, bringing $\\Gamma$ below $\\sim 4/3$ and triggering the collapse (Barkat et al. 1967; Bond et al. 1984; Kippenhahn \\& Weighrt 1990; Fryer et al. 2001).  For sufficiently massive stars, of mass $M \\agt 260 M_{\\odot}$, such unstable cores form a black hole, while for less massive cases, the outcome is a pair-unstable supernova (Fryer et al. 2001).  These examples illustrate that the gravitational collapses set in when the stability of progenitors changes, and hence, a star marginally stable against quasi-radial oscillations may be regarded as an initial condition of the collapse. Massive stars in nature are rapidly rotating in general (Bond et al. 1984). Thus, determining the quasi-radial stability for rotating stars is an important subject in astrophysics for elucidating a plausible initial condition for the stellar collapse.  In many cases, one simply assumes that progenitors of neutron stars or black holes start collapsing when their adiabatic index $\\Gamma$ decreases below a well-known critical value, 4/3, against a quasi-radial instability. The number $4/3$ is the correct value only when the progenitor is not rotating. If the progenitor is rotating, it is not the valid criterion, since angular momentum has the effect of stabilizing the star against gravitational collapse. In particular, the number 4/3 would not be a good indicator of the instability for pair-unstable collapse, since the progenitor stars are likely to be rapidly rotating (Bond et al. 1984, Fryer et al. 2001).  According to a Newtonian theory, slowly rotating stars that satisfy the following relation are unstable against a quasi-radial oscillation (Tassoul 1978; see also \\S 2): \\beq 3\\Gamma - 4 - 2 (3\\Gamma-5)\\beta < 0. \\label{eq1} \\eeq Here, \\beq \\beta \\equiv {T \\over W}, \\eeq and $T$ and $W$ are rotational kinetic and gravitational potential energies, respectively. In the following, we restrict our attention to the rigidly rotating case for simplicity. Based on our numerical analysis in Newtonian gravity, the maximum value of $\\beta$ for rigidly rotating stars with polytropic equations of state, which is achieved when the velocity at the equatorial surface is equal to the Kepler velocity (i.e., at the mass-shedding limits), is approximately written as \\beq \\beta_{\\rm max} \\approx 0.00902 + 0.124 \\biggl( \\Gamma-{4\\over 3}\\biggr). \\label{eq2} \\eeq Combining equations (\\ref{eq1}) and (\\ref{eq2}), we find that for $\\Gamma \\alt 1.328$, all the rigidly rotating stars are unstable against a quasi-radial oscillation in Newtonian gravity.  However, even for $1.328 \\alt \\Gamma < 4/3$, the stars can be stabilized by the effect of rotation.  General relativistic gravity destabilizes the rapidly rotating stars even in the case of $\\Gamma > 1.328$. One of the purposes of this paper is to numerically determine the criteria for the onset of the instability. In particular, we focus on the secular stability of rigidly rotating stars with soft equations of state in equilibrium against a quasi-radial oscillation. The stability is determined using the turning-point method (Friedman et al. 1988). To model the rotating stars, we adopt the polytropic equations of state with $n \\agt 3~(\\Gamma \\alt 4/3)$.  {}From a general relativistic point of view, it is interesting to ask how large the nondimensional angular momentum parameter $q \\equiv J/M^2$ can be for the marginally stable stars that may be plausible progenitors of gravitational collapse. Axisymmetric hydrodynamic simulations in general relativity (Nakamura 1981; Stark \\& Piran 1985; Piran \\& Stark 1986; Nakamura et al. 1987; Shibata 2000) have provided many numerical results that show that for $q > 1$, a black hole is not formed after the collapse. That is, the value of $q$ may be a good indicator in predicting the final fate of the stellar collapse of massive objects. As Baumgarte and Shapiro (1999) indicate, the value of $q$ for a marginally stable star with $\\Gamma=4/3$ is very close to unity. Cook et al. (1994) show that the critical value of $q$ increases with decreasing of $\\Gamma$ for $4/3 < \\Gamma \\alt 5/3$. This implies that for $\\Gamma < 4/3$, the critical value is likely to be larger than unity, even for the case of $4/3-\\Gamma \\ll 1$. The second purpose of this paper is to confirm this fact.   The third purpose is to predict the outcomes of gravitational collapse of the marginally stable stars. As mentioned above and confirmed below, the value of $q$ for the marginally stable stars with $\\Gamma < 4/3$ can be larger than unity. This suggests that a rotating star with $\\Gamma < 4/3$ of mass large enough collapse to a black hole in the nonrotating case may not form a black hole because of its significant effect of the angular momentum. In the analysis, we assume that (i) the collapse proceeds in an axisymmetric manner, (ii) the rest-mass distribution as a function of the specific angular momentum is preserved, and (iii) the pressure never halt the collapse. We indicate that the criterion for no black hole formation, $q>1$, which has been suggested so far, is not very good for rigidly rotating stars with the soft equations of state. We predict that even for $1 < q \\alt 2.5$, a black hole may be formed as a result of gravitational collapse of the rigidly rotating and marginally stable stars. However, for $q \\agt 2.5$, the significant angular momentum will prevent the collapsing stars from directly forming a black hole.   The paper is organized as follows. In \\S 2, we present a semi-analytic calculation for determining the stability of rotating stars against a quasi-radial oscillation. In \\S 3, the secular stability of rotating stars in general relativity is numerically determined. In \\S 4, we predict the final outcomes of the stellar collapse for very massive stars, assuming that the initial condition is a marginally stable star, as determined in \\S 3. Section 5 is devoted to a summary.  Throughout this paper, the pressure is given by the polytropic equation of state as \\beq P= K \\rho^{\\Gamma},~~\\Gamma=1+{1 \\over n}, \\eeq where $P$ is the pressure, $\\rho$ the baryon rest-mass density, $K$ the polytropic constant, and $n$ the polytropic index, which is chosen in the range between 3 and 3.05. We adopt the geometrical units $G=c=1$, where $G$ and $c$ denote the gravitational constant and the speed of light, respectively. ",
        "conclusions": "We have investigated the secular stability of rigidly rotating stars against a quasi-radial oscillation in general relativity. It is found that all the rigidly rotating stars with $\\Gamma \\leq 1.328$ are unstable against a quasi-radial oscillation. Stars with $1.329 \\alt \\Gamma < 4/3$ that are unstable for the spherical case can be stabilized by the effect of rotation. Therefore, for rapidly rotating stars for which the adiabatic index gradually decreases from $\\sim 4/3$ to smaller values, the instability against gravitational collapse will set in when the value of $\\Gamma$ decreases below $\\approx 1.329$ (not 4/3).  The marginally stable stars for $\\Gamma=1.329$--4/3 have been determined numerically. It is found that for $1.329 \\leq \\Gamma \\leq 1.332$, the nondimensional angular momentum parameter $q$ for {\\it all} rigidly rotating and marginally stable stars is larger than unity. The value of $q$ for the marginally stable stars increases with decrease of the compactness. Therefore, for a sufficiently small compactness, the value of $q$ exceeds unity even for $1.332 < \\Gamma < 4/3$. These results are in good agreement with those derived by a simple semi-analytic calculation presented in \\S 2.  The final outcomes of axisymmetric collapse of the rigidly rotating and marginally stable stars are predicted assuming that the pressure never halts the collapse and that the mass distribution as a function of the specific angular momentum is preserved. It is found that even for the case $q > 1$, a black hole will form as a result of the gravitational collapse if the value of $q$ is smaller than $\\sim 2.5$. This is due to the fact that the angular momentum in the central region of the rigidly rotating stars is not large enough and that the effective value of the nondimensional angular momentum parameter $q$ for the inner region is smaller than unity. The outcome of the collapse in such cases will be a rapidly rotating black hole surrounded by a massive disk. For $q \\agt 2.5$, even the central region has angular momentum large enough to prevent collapsing to a black hole. In this case, the outcome will be a disk or a torus. It will be subsequently unstable against nonaxisymmetric deformation. Once the nonaxisymmetric structure is developed, the angular momentum will be transported outward by the nonaxisymmetric effects, decreasing the value of $q$ in the inner region below unity and eventually forming a seed black hole.  Although the conjecture mentioned above is reasonable, the final fate of the collapse of marginally stable stars with $q \\agt 1$ can be determined only by fully general relativistic simulations. In particular for $q \\agt 2.5$, three-dimensional simulations are necessary since the collapsing star is likely to become unstable against nonaxisymmetric deformation. Since the length scale will change by a factor of $10^4$ during the collapse, huge computational resources will be necessary for the simulation. The simulation for this problem will be one of the computational challenges in the field of numerical relativity."
    },
    "0403/astro-ph0403491_arXiv.txt": {
        "abstract": "The central regions of many interacting and early-type spiral galaxies are actively forming stars. This process affects the physical and chemical properties of the local interstellar medium as well as the evolution of the galaxies. We observed near-infrared \\hh\\ emission lines: $v=$ \\vone, \\vthree, \\vonez, and \\vtwo\\ from the central $\\sim 1$~kpc regions of the archetypical starburst galaxies, M82 and NGC 253, and the less dramatic but still vigorously star-forming galaxies, NGC 6946 and IC 342. Like the far-infrared continuum luminosity, the near-infrared \\hh\\ emission luminosity can directly trace the amount of star formation activity because the \\hh\\ emission lines arise from the interaction between hot and young stars and nearby neutral clouds. The observed \\hh\\ line ratios show that both thermal and non-thermal excittion are responsible for the emission lines, but that the great majority of the near-infrared \\hh\\ line emission in these galaxies arises from energy states excited by ultraviolet fluorescence. The derived physical conditions, e.g., far-ultraviolet radiation field and gas density, from \\cii\\ and \\oi\\ lines and far-infrared continuum observations when used as inputs to photodissociation models, also explain the luminosity of the observed \\hh\\ \\vone\\ line. The ratio of the \\hh\\ \\vone\\ line to far-IR continuum luminosity is remarkably constant over a broad range of galaxy luminosities; $L_{H2}/L_{FIR}$ $\\simeq 10^{-5}$, in normal late-type galaxies (including the Galactic center), in nearby starburst galaxies, and in luminous IR galaxies (LIRGs: $L_{FIR}$ $>$ $10^{11}$~\\lsun). Examining this constant ratio in the context of photodissociation region models, we conclude that it implies that the strength of the incident UV field on typical molecular clouds follows the gas density at the cloud surface. ",
        "introduction": "The {\\it Infrared Astronomical Satellite (IRAS)} detected a significant population of galaxies that emit large fractions of their bolometric energy in the mid- and far-infrared (IR) bands. The majority of luminous IR galaxies (LIRGs: $L_{FIR}$ $>$ $10^{11}$~\\lsun ) are powered by massive stars that have recently formed in their central kiloparsec \\citep{Genzel98}. The far-IR luminosities of these galaxies therefore measure the amount of star formation activity \\citep{Kennicutt98}. Local starburst galaxies with far-IR luminosities of $10^{10} - 10^{11}$ \\lsun \\ and normal spiral galaxies with $L_{FIR}$ $\\simeq 10^9$ \\lsun \\ also show evidence for active star formation in their central ($< 1$~kpc) regions, even though their far-IR luminosities are much lower than those of LIRGs \\citep{Thronson87,Stacey91}.  The formation of O and B stars affects the physical and chemical state of the interstellar medium through stellar UV radiation and through the effects of stellar winds and supernovae \\citep{Telesco88}. Strong near-IR \\hh\\ emission lines have been observed from the central regions of galaxies, especially from starburst galaxies and LIRGs. Since \\hh\\ vibration-rotation emission can trace the interaction of the neutral ISM with the UV radiation from hot young stars and its interaction with sources of dynamical energy such as supernovae, the \\hh\\ luminosity may also be a good tracer of star formation activity. The main motivation of this paper is to explore how well we can use this additional, readily observable tracer of the interaction of massive stars with the dense interstellar medium to study massive star formation in the inner parts of galactic disks.  The observational part of this paper emphasizes new, large-beam observations of emission in the vibration-rotation transitions of \\hh\\ in the near-IR from nearby late-type galaxies and low luminosity starbursts. For normal galaxies, observations of near-IR \\hh\\ lines have been restricted to local bright regions because the average surface brightness of the \\hh\\ emission is low. In a previous study, we used an instrument optimized for observations of low surface brightness line emission in the near-IR to detect enormously extended ($\\sim 3\\degr$ or $\\sim 400$~pc) \\hh\\ emission along the plane of the Galaxy near the Galactic center \\citep{Pak96a,Pak96b}. The integrated and dereddened \\hh\\ $v=$ \\vone\\ luminosity from the extended region ($8 \\times 10^3$ \\lsun) dominates over that from the previously known 5~pc circumnuclear gas ring ($4.5 \\times 10^2$ \\lsun, \\citealt{Gatley84}).  Our results for the inner 400 pc of the Milky Way motivated us to begin a study of \\hh\\ emission from the central kiloparsec of other nearby galaxies where it may be possible to detect extended, low surface brightness \\hh\\ emission using the same large-aperture Fabry-Perot spectrometer used in the Galactic work. The large-beam observations we present here make it possible to compare for the first time the \\hh \\ emission to the far-IR continuum and far-IR fine-structure line emission at similar angular resolution and over similar angular scales. The observations also give us measurements of \\hh \\ line strengths from nearby objects with linear scales comparable to those obtained with slit spectrometers for the more distant luminous IR galaxies. With this broad range of spatially compatible observations, we investigate whether it is possible to formulate a generalized star formation law which can be applied to galaxies over a wide range of levels in star formation activity.  In order to analyze the \\hh\\ line data and to compare them with other star formation tracers, we need to identify the  excitation mechanism responsible for the \\hh\\ emission. In the central regions of galaxies, however,  the  excitation mechanism for the \\hh \\ emission has long been the subject of intense controversy (see the reviews of \\citealt{Mouri94}; \\citealt{Goldader97a}). \\citet{Puxley88} observed nearby spiral galaxies using a circular variable filter ($\\lambda/\\Delta \\lambda = 120$) and a beam full-width at half maximum (FWHM) of 19\\farcs 6. They detected \\vone, \\vtwo, and \\vonez\\ lines and concluded that the \\hh\\ line populations are non-thermal. However, \\citet{Moorwood90} presented high spectral resolution data ($\\lambda/\\Delta \\lambda = 1500$) and pointed out that the \\vtwo\\ flux measurement of \\citet{Puxley88} could be affected by unresolved stellar absorption lines. They argued that the \\hh\\ emission is mostly from shocks. \\citet{Mouri94} suggested that high spectral resolution observations are essential to discriminate between the excitation mechanisms.  The ultimate energy source responsible for the excitation of the \\hh\\ vibration-rotation emission can lie either in the absorption of UV photons from the ambient interstellar radiation field or from local OB stars, or in shock waves excited by supernovae, protostellar flows, or other energetic events \\citep{Black87,Burton92}. Hydrogen molecules that absorb far-UV radiation make a transition to an excited electronic state.  At low densities, this transition is followed by fluorescence to the ground electronic state and a cascade through that state's vibration-rotation levels \\citep{Sternberg89a,Pak03}. This interaction between \\hh\\ molecules and UV radiation takes place at the edges of molecular clouds where hydrogen makes the transition from predominantly atomic to predominantly molecular gas, in the photodissociation region (PDR: \\citealt{Tielens85}; \\citealt{Hollenbach99}). If $n_H$ $<$ $5 \\times 10^4$~cm$^{-3}$, the relative intensities of vibration-rotation transitions in the near-IR arising in UV-excited \\hh\\ are insensitive to density or to UV field strength \\citep{Black87} and depend only on the branching ratios in the downward cascade. In dense ($n_H$ $>$ $5 \\times 10^4$~cm$^{-3}$) gas heated by intense far-UV radiation, \\citet{Sternberg89b} showed that inelastic H$-$\\hh\\ and \\hh$-$\\hh\\ collisions can heat the gas at the cloud edge to temperatures approaching 1000~K. A combination of the collisional de-excitation and collisional excitation in the high temperature layer alters the \\hh\\ level populations in this UV-excited gas. The expected changes in relative line strength with density have been verified by observations of molecular cloud edges \\citep{Luhman97}. In shocked regions, the \\hh\\ vibration-rotation energy level populations in the ground electronic state are thermalized by collisions and the level populations are usually in LTE \\citep{Shull82,Black87}. Because the level populations are thermalized and reflect a similar range of temperatures, shock-heated gas and hot, dense gas heated by UV radiation are difficult to distinguish.  Most previous observations of the central regions of galaxies showed that the \\hh\\ line ratios of \\vtwo\\ to \\vone\\ lie between the ratios one would expect for fluorescently excited \\hh\\ and what one would see in observations of UV-heated but collisionally excited gas or of gas collisionally heated in shocks \\citep{Doyon94,Kawara90,Koornneef96,Lumsden94,Moorwood90,Puxley88,Puxley90,Davies03}. These intermediate ratios suggest that both pure UV-fluorescence and thermalized UV-excited (or shock excited) gas are responsible for the \\hh\\ emission \\citep{Davies03}. Detection of \\hh\\ lines arising from upper vibrational states with $v > 2$ is a strong indication of the presence of UV-fluorescent excitation since the high energies of the upper states make it very hard to excite these transitions thermally. There are marginal detections of such transitions from a few galaxies, e.g., a giant \\hii\\ region NGC 5451 in M101 \\citep{Puxley00}, Arp 299 \\citep{Sugai99}, and some ultraluminous IR galaxies \\citep{Davies03}. Another method for discriminating between the \\hh\\ excitation mechanisms is to measure the \\hh\\ nuclear spin degeneracy ratio (ortho- to para-ratio), which may deviate from its formation value if the gas is thermally processed, but remains constant in the case of non-thermal excitation. Observations of \\vone\\ and \\vonez\\ showed that the \\hh\\ emission in the galactic centers arises mostly from UV-fluorescence \\citep{Harrison98,Sugai99,Puxley00}. More recently, \\citet{Davies03} have observed vibration-rotation transition lines of \\hh\\ from the nuclei of seven ULIRGs and concluded that the line emission arises from dense PDRs irradiated by intense far-UV radiation fields. In galactic sources, we have used observations of \\hh\\ lines from high vibrational states \\citep{Luhman95}, as well as comparisons of extended \\hh\\ emission and tracers of UV excited PDRs such as far-IR continuum and \\cii\\ emission to show that much of the \\hh\\ emission arises from UV-excited gas \\citep{Luhman96,Luhman97,Klumpe99}. This comparison with far-IR lines and continuum has been more difficult in other galaxies because of the significant differences between the \\cii\\ and far-IR\\ beam sizes and the typical slit sizes used for near-IR spectroscopy of \\hh. In this work, we will use both \\hh\\ line ratios and a study of various tracers of star forming clouds in the context of PDR models to learn more about the relative significance of various types of \\hh\\ excitation.  In Section~\\ref{sec:Observations}, we present a description of our spectroscopic observations of the starburst galaxies M82 and NGC 253, and the normal spiral galaxies NGC 6946 and IC 342. In Section~\\ref{sec:Results} we first show the observed \\hh\\ spectra and their extinction-corrected line fluxes. We then show the spatial distribution of the \\hh\\ \\vone\\ line intensity and compare it to the distribution of other PDR tracers, e.g., CO line emission and sub-mm continuum. In Section~\\ref{sec:Discussion} we analyze the nature of the \\hh\\ excitation mechanism by comparing the measured \\hh\\ line ratios with thermal and non-thermal models. We also calculate the luminosity ratios of \\hh\\ to far-IR and \\cii\\ to far-IR, and discuss the implications of the observed ratios by comparison to predictions by PDR models. Section~\\ref{sec:Conclusions} summarizes our conclusions. ",
        "conclusions": "\\label{sec:Conclusions}   We have analyzed near-IR \\hh\\ emission from the central $\\sim 1$~kpc regions in the nearby starburst galaxies, M82 and NGC 253, and the vigorously star-forming galaxies, NGC 6946 and IC 342. The line ratios of various \\hh\\ vibration-rotation lines, e.g., \\vtwo, and \\vonez\\ to \\vone, show that, while both thermal and non-thermal excitation are responsible for the emission, non-thermal excitation predominates. In addition to the \\hh\\ line ratios, the constant ratio between \\hh\\ \\vone, \\cii, and far-IR luminosity suggests that the emission in the two lines and the far-IR continuum originate together in photodissociation regions.  We also compared scan maps of \\hh\\ \\vone\\ emission along the major axes (or molecular bars) of the galaxies with maps of CO and sub-mm emission. The distribution of \\hh\\ line emission correlates with the sub-mm continuum emission, and also agrees well with the distribution of CO emission. The \\hh\\ emission traces active star formation regions where the far-UV fields are intense and the gas is dense and warm. The correlation of \\hh\\ and sub-mm emission suggests that clouds have similar structure or that the star formation and cloud distributions are similar in the central regions of these galaxies.  We compare the luminosities of \\hh\\ \\vone, \\cii\\ 158 \\micron, and far-IR continuum. The $L_{H2}/L_{FIR}$ ratios of normal spiral and starburst galaxies are the same. Comparing the observed ratio to PDR models, we conclude that the column of warm \\hh\\ line emitting gas is limited by the extinction of far-UV photons by dust, rather than by \\hh\\ self-shielding in the PDRs. In this case, the \\hh\\ intensity is only sensitive to the gas density. The constant $L_{H2}/L_{FIR}$ ratio over the large far-IR luminosity ranges of the galaxies, then requires that the gas density in the PDRs increase as the far-UV (or far-IR) intensity increases.      We thank Mark R. Swain for help with observations at the Palomar Observatory, Ewine van Dishoeck for providing PDR model results, and Michael Luhman and Thomas Benedict for their work on the UTFPS. This work was supported by NSF Grant 95-30695. S.P. appreciates the warm hospitality of the staff of the Institute of Space and Astronautical Science in Japan.      \\appendix"
    },
    "0403/hep-th0403223_arXiv.txt": {
        "abstract": "When $p$-dimensional branes annihilate with antibranes in the early universe, as in brane-antibrane inflation, stable $(p-2)$-dimensional branes can appear in the final state. We reexamine the possibility that one of these $(p-2)$-branes could be our universe.   In  the low energy effective theory, the final state branes are cosmic string defects of the  complex tachyon field which describes the instability of the initial state.   We quantify the dynamics of formation of these vortices.  This information is then used to estimate the production of massless gauge bosons on the final branes, due to their coupling to the time-dependent tachyon background, which would provide a mechanism for reheating after inflation.  We improve upon previous estimates indicating that this can be an efficient reheating mechanism for observers on the brane. ",
        "introduction": "In the last few years significant progress has been made in constructing string theoretic  cosmological models where inflation is driven by the naturally occuring potentials between D-branes and their antibranes \\cite{DbraneInflation}-\\cite{newinfrefs}. The formation of lower-dimensional branes at the end of inflation can lead to interesting signals: cosmic-string-like or higher dimensional defects could be observable remnants  \\cite{KibbleMechanism}, possibly providing a rare clue to the stringy origin of inflation.  A more radical idea was explored in \\cite{Cline:Reheating}: perhaps our own observable universe is such a defect in the higher-dimensional spacetime predicted by string theory.  Since the stable branes in Type IIB string theory have spatial dimensionalities which are odd, a 3-brane would have descended from annihilation of 5-branes in this picture, and thus they would be codimension-two defects in of  the effective 6D theory.  Codimension-two braneworlds have attracted interest lately because of their novel features, which might have some bearing on the cosmological constant problem \\cite{codimension2,CDGV}.  Our interest in this scenario is motivated by questions about the efficiency of reheating in brane-antibrane inflation \\cite{ProblemsWithReheating}.  It is possible that the energy  density liberated from the brane collisions will be converted mostly into closed string states,  ultimately gravitons, and not necessarily into visible radiation \\cite{ClosedStringRadiation}.  A generic mechanism which could avoid this problem was proposed in \\cite{Cline:Reheating}, wherein the reheating in D-brane driven inflation is due to the coupling of massless gauge fields to a time-dependent tachyon condensate, which describes the annihilation process.  However, ref.\\ \\cite{Cline:Reheating} considered only the formation of a tachyon kink instead of the more realistic case of a vortex, and it used a somewhat crude ansatz for the background tachyon field.   The problem of finding the actual tachyon background predicted by string theory was studied numerically in  \\cite{Cline:DbraneCondensation} but no attempt was made to improve on the reheating computation. In this paper we aim to analytically determine the dynamics of formation of lower dimensional branes described as tachyon defects---both kinks and vortices---and to improve on the reheating calculation of \\cite{Cline:Reheating}.  Let us begin by describing the scenario we have in mind.  In the simplest version of D-brane inflation  a parallel brane and antibrane begin with some separation between them in one of the extra dimensions. Although parallel branes are supersymmetric and have no force between them, the brane-antibrane system breaks supersymmetry so that there is an attractive force and hence a nonvanishing potential energy.  It is the latter which drives inflation.  Once the branes have reached a critical separation one of the stretched string modes between the branes, $T$, becomes tachyonic and the branes become unstable to annihilation.  The tachyon field starts from the unstable maximum $T=0$ and rolls towards the vacuum $T \\rightarrow \\pm \\infty$.  However, topological defects may form through the Kibble mechanism \\cite{KibbleMechanism} so that $T=0$ stays fixed at the core of the defect.  These defects are known to be consistent descriptions of branes whose dimension is lower than that of the original branes \\cite{Sen:TachyonKinkAndVortex,DescentRelations}.  For example, the brane-antibrane system has a complex tachyon field, leading to vortices which represent codimension-two branes.  On the other hand, an unstable brane has a real tachyon which leads to kinks representing codimension-one branes.  The formation of tachyon defects at the endpoint of D-brane inflation is a dynamical process where the tachyon couples to gauge fields which will be localised on the descendant brane.  It is thus expected that some radiation will be produced by the rolling of the tachyon and the problem of reheating becomes quantitative: can this effect be efficient enough to strongly deplete the energy density of the tachyon fluid so the the universe starts out being dominated by radiation rather than cold dark matter?  It is important to stress that though the situation is somewhat analogous to that of hybrid inflation  (where the tachyon plays the role of the unstable direction in field space which allows for inflation to end  quickly) the mechanism for reheating is qualitatively different.  The difference is that in the low  energy effective field theory which describes the tachyon $T$, the potential is minimized at $T=\\pm\\infty$ and there are no oscillations about the minimum of the potential. In a normal hybrid inflation model, $T$ would have a minimum at some finite value and the oscillations of $T$ around its minimum would give rise to reheating in the usual way.  In the present case, the time dependence of the background is monotonic, not oscillatory.  Reheating thus might seem to resemble gravitational particle production \\cite{quintinf} rather than the standard picture in which the inflaton decays.  However, in this work we highlight an important difference between reheating through tachyon condensation and gravitational particle production, which can make the former much more efficient: there is a divergence in the stress-energy tensor of the tachyon field within a finite time, which corresponds to the formation of the lower-dimension D-brane.  In this paper we study analytically the dynamical formation of the tachyon vortex and improve  the reheating calculation, using a slightly simplified model of particle production by tachyon condensation, which captures the essential physics revealed by the analysis of vortex formation. In section II we review the  formation of tachyon kinks which describes the condensation of a brane to a brane of codimension-one. In  sections III-V we study analytically the formation of a tachyon vortex on the brane-antibrane pair. Section III introduces Sen's action for the complex tachyon field describing this situation. Section IV presents new analytic results for the time-dependent, complex tachyon field representing vortex formation, both near to and far from the vortex core.  In section V we show that the stress-energy tensor for the system splits into a localized, singular piece describing the descendant branes, plus  a bulk contribution that describes the rolling tachyon condensate.  Section VI introduces the effective action for U(1) gauge bosons which become localized on the final-state 3-brane, in the rolling tachyon background.  This provides a model for the visible radiation produced during reheating. In section VII we calculate the energy density of this produced radiation on the 3-brane, using some reasonable simplifying assumptions. Section VIII gives our conclusions, including speculation about how the final brane-antibrane system could be stabilized. ",
        "conclusions": "We have argued for the possibility that our visible universe might be a codimension-two brane left over from annihilation of a D5-brane/antibrane pair at the end of inflation.  In this picture, reheating is due to production of standard model particles ({\\it e.g.} photons) on the final branes, driven by their couplings to the tachyon field which encodes the instability of the initial state as well as the vortex which represents the final brane. We find that reheating can be efficient, in the sense that a sizable fraction of the energy available from the unstable vacuum can be converted into visible radiation, and not just gravitons.   The efficiency of reheating is greatest  if the radius of compactification of the extra dimensions is larger than $2$-$3$ times the string length $l_s$.  The efficiency also depends on phenomenological parameters we had to introduce by hand in  order to cut off ultraviolet divergences in the calculated particle production rate: namely $\\epsilon$, a nonvanishing radius for the final brane, and $\\Lambda$, an explicit cutoff on the momentum of the photons produced.  The latter must be introduced to correct for discontinous time derivatives in our simplified model of the background tachyon condensate; the actual behavior of the condensate corresponds to a cutoff of order $\\Lambda\\sim 1/l_s$.  It is less obvious why the effective field theory treatment should give divergent results as the thickness of the final brane ($2\\epsilon$) goes to zero, but it seems clear that a fully string-theoretic computation would give no such divergence, and therefore it is reasonable to cut off the field-theory divergence at $\\epsilon\\sim l_s$.  Given only these mild assumptions, our estimates (\\ref{mainresult},\\ref{rhoceq}) predict that the fraction of the available energy which is converted into visible radiation is $\\rho/\\rho_c \\cong \\pi\\alpha(M_s) \\cong 0.25$.  This simple estimate counts only photons; in a more realistic calculation, it would be enhanced by the number of light degrees of freedom which couple to the tachyon, which could be much greater 1.  Moreover it could also be enhanced by the production of massive KK modes, which correspond to string excitations when the vortex has formed \\cite{HN,MZ}.  Our analysis makes significant improvements to the previous work of \\cite{Cline:Reheating}.  We considered the formation of a vortex in the tachyon field rather than a kink, in line with the descent relations for stable Dp-branes.  We found analytic solutions for the tachyon field which give the time dependence in the vicinity of the defect while it is forming, for both the vortex and the kink solutions.  In the latter case we verified that this solution reproduces the  known dynamics of kink formation which was determined numerically in \\cite{Cline:DbraneCondensation}, giving us more confidence in the vortex solutions, which are quite analogous. Our explicit solution for the tachyon background is nevertheless too complicated for computing the production of particles on the defect.  We therefore approximated it by a simpler ansatz with the same qualitative behavior, which allows for analytic solutions of the gauge fields in  the background.  This ansatz resembles a gauge field in a 6D spacetime with two compact spatial  dimensions which are contracting with time, and leads to simple analytic results for the energy density of photons produced during the contraction, in the regime where the extra dimensions are large compared to the string scale.    There are still some outstanding questions to be addressed concerning this scenario.  First, we have made reference to the Kibble mechanism for the creation of the final state defects.  If we assume the causal bound of one defect per Hubble volume then this would imply that the size of the extra dimensions must exceed the inverse Hubble rate; otherwise there would be enough time for the fields to straighten themselves out and the putative vortex-antivortex pair would immediately annihilate.  For example if we take the string scale $M_s$ to be $10^{16}$ GeV then  $H\\sim M_s^2/M_p$ and we would need the compactification scale to be of order $R_0 \\sim (M_p/M_s) l_s$. Our results indicate that efficient reheating is compatible with a large compactification scale.  However, taking the remaining four extra dimensions (which the initial 5-brane/antibrane pair do not wrap) to be string scale is not consistent with getting inflation since the initial state 5-brane and antibrane cannot be sufficiently separated to satisfy the slow roll conditions.  One the other hand, our results indicate that the reheating \\emph{can} be efficient for $R_0$ only a few times the string length, though in this scenario a naive application of the Kibble mechanism does not favor having the final state defects span the three large dimensions.  These requirements may not be prohibitive since the question of how these defects form dynamically at the end of inflation is a quantitative one which merits further investigation. In principle the correlation length for the initial fluctuations of the tachyon field could be as small as the string length.  We point out also that it is possible that the dynamics of the formation of tachyon defects is qualitatively different from defect formation in a conventional scalar field theory. For example, the numerical investigation of \\cite{Cline:DbraneCondensation}, it was found that small kinks in the initial  configuration which are in causal contact with each other {\\it do not} dynamically straighten themselves out as they would in a conventional, nontachyonic field theory.  Instead, every place where the field crosses zero in the initial state develops a full-blown kink, so long as there was enough energy in the bulk to produce the required number of kinks.  Another indication that the dynamics of the tachyon field may be qualitatively different from an ordinary scalar field theory comes from the \\cite{Causal} in which the causal structure of the tachyon Dirac-Born-Infeld action was studied.  The authors of \\cite{Causal} found that small fluctuations of the tachyon field propagate according to an effective metric which depends on the tachyon background.  In the case of a homogeneous rolling background it was found that as the condensation proceeds the effective metric contracts to the Carroll limit of the Lorentz group so that the tachyon light cone collapses into a timelike half line and the tachyon fields at different spatial points are decoupled.  We feel that quantitatively determining the dynamics of the formation of tachyon defects at the endpoint of D-brane inflation is a question which deserves further investigation.    In the present analysis we have not included the gravitational or Ramond-Ramond forces between the vortex and antivortex which would attract them toward each other and lead to their eventual annihilation.  How do we insure that the braneworld on which we are supposed to live is safe from annihilation with an antibrane in the bulk?   One possibility is to have warping caused by a stack of branes which wraps only the equator of the extra dimensions, as illustrated in fig. \\ref{lastfig}.  Such a braneworld scenario using the AdS soliton solution for the bulk has been considered in ref.\\ \\cite{LMW,BCCF,CDGV}.  The advantage for our scenario is that the warping can provide a barrier to the annihilation of the brane-antibrane pair, since it is energetically favorable for them to remain within their respective throats.  \\begin{figure} \\centerline{\\epsfxsize=3.5in \\epsfbox{adssol.eps}}\\caption{Warped compactification with branes localized in throats on opposite sides of a stack of branes. \\label{lastfig}} \\end{figure}  In solving for the tachyon background we have also ignored the possibility of caustic formation in the bulk \\cite{Felder:Caustics} by taking initial profiles without too much  curvature. It is possible that caustic formation may be an artifact of the derivative truncation which leads to the Born-Infeld type of Lagranian for the tachyon.  See \\cite{p-adic} for a discussion of the problems of dynamical equations with infinitely many derivatives.  Brane-antibrane inflation and braneworld cosmology are two of the most important applications of string theoretic ideas to cosmology.  We find it intriguing that these two ideas might be combined in the way we have described.  An outstanding challenge is to find some observable signatures that would be able to test our scenario, for example through the gravitational wave component which is expected to be a major component of the radiation produced during reheating.  We thank Horace Stoica for helpful discussions concerning warped compactifications as well as  Robert Brandenberger and Henry Tye for helpful comments regarding the Kibble mechanism. Our research is supported by the Natural Sciences and Engineering Research Council (NSERC) of Canada and Fonds qu\\'eb\\'ecois de la recherche sur la nature et les technologies (NATEQ)."
    },
    "0403/astro-ph0403614_arXiv.txt": {
        "abstract": "Scalar-tensor theories of gravity provide a consistent framework to accommodate an ultra-light quintessence scalar field. While the equivalence principle is respected by construction, deviations from General Relativity and standard cosmology may show up at nucleosynthesis, CMB, and solar system tests of gravity. After imposing all the bounds coming from these observations, we consider the expansion rate of the universe at WIMP decoupling, showing that it can lead to an enhancement of the dark matter relic density up to few orders of magnitude with respect to the standard case. This effect can have an impact on supersymmetric candidates for dark matter. ",
        "introduction": "According to our current understanding \\cite{rev-cosmo}, Dark Matter (DM) and Dark Energy (DE) represent the two major components of the present Universe.  Surprisingly, it is found that the DM and DE energy densities, $\\rho_{DM}$ and $\\rho_{DE}$, are today roughly the same (differing only by a factor of two), while their ratio has been varying by several orders of magnitude in the past history of the Universe.  It seems quite natural, then, to explore the possibility of a DM--DE interaction which could account for this coincidence. This approach, however, is not free from problems if the DE component is interpreted in terms of a dynamical quintessence \\cite{quint-review} scalar field. Indeed, such a scalar is constrained to be extremely light in order to fit the data, giving rise to unwanted long-range forces which may represent a severe threat to the equivalence principle. In addition, couplings of the quintessence scalar with the gauge field strengths are potential sources of dangerous time variations of the fundamental constants\\footnote{See, for example, Refs. \\cite{carroll1,mpr,dpv,dam3a,dam3b} for different approaches on these problems.}.  A possible way of coping with a very light scalar while avoiding these shortcomings, is to choose to work in the framework of a scalar-tensor gravity (ST) theory \\cite{bd}, where by construction matter has a purely metric coupling with gravity.  It has been shown \\cite{max} that in this case, the scalar field can benefit from an attraction mechanism which, during the matter dominated era, makes ST overlapping with standard General Relativity (GR). At the same time, ST may possess a second ``attraction mechanism'' \\cite{max} which will ensure the correct evolution of $\\rho_{DE}$ along a so-called `tracking' \\cite{tracker} solution.  While ST can very closely reproduce standard GR at the present time, it may however lead to major differences in the past evolution of the Universe, differences which may result in observable consequences for us today.  For example, it has been shown \\cite{joyce,damour-pichon,santiago,carroll2} that ST theories may have a profund impact on nucleosynthesis.  At the same time, a curious fact has recently come to attention: a non--conventional dynamics of the quintessence scalar in the past history of the universe may remain `hidden' to the available cosmological observations, but manifest itself through the DM relic abundance \\cite{salatirosati,ullio,comelli}.  It is then worth studying if ST may provide a viable quintessence candidate and at the same time have an impact on the DM relic abundance. We have considered this possibility and computed explicitly the differences from the standard case.  When considering ST theories, the departures from standard cosmology are mainly due to the different expansion rate ($\\tilde{H}$) which they determine. Such deviation from the usual expansion rate of GR bears potentially relevant consequences in all those phenomena which are closely dependent on the timing in which they occur.  The aim of the present work is to study the possible modifications of the expansion rate of the Universe in ST at the time of Cold Dark Matter (CDM) freeze-out, focussing on Weekly Interacting Massive Particles (WIMPs) as the most natural candidates for CDM.  As it is well known, their present relic density depends on the precise moment they decouple and, in turn, on the precise moment the WIMP annihilation rate equals the expansion rate of the Universe. We expect then that a variation of $\\tilde{H}$ in the past may lead to measurable consequences on the WIMP relic abundance.  In order to assess the allowed departure of $\\tilde{H}$ from its standard value at WIMPs freeze-out, we have to take into account the bounds imposed on ST by phenomena at later epochs \\cite{Riaz}.  We already mentioned the crucial test of nucleosynthesis: passing the BBN exam will imply rather severe restrictions on the coupling of the scalar field of ST to matter and also on its initial conditions at temperatures much higher than the WIMP freeze-out. Coming to later epochs, we have to consider the photon decoupling and the consequent restrictions imposed by CMB data \\cite{CMB}. Interestingly enough, we will show that the BBN ``filter'' on ST is so efficient that it drastically limits any visible effect on the CMB spectrum (in particular, shifts in the peaks positions are forced to be smaller than the experimental error). More restrictive than the CMB probe turn out to be the GR tests, in particular after the tight bound on the post-newtonian parameter $\\gamma$, recently provided by the Cassini spacecraft \\cite{cassini}. This constraint becomes quite relevant when combined with that on the scalar equation of state of $w_\\phi$ coming from SNe Ia data \\cite{SNe}. We will find that significant deviations from standard cosmology are possible only if the DE equation of state differs appreciably from $-1$ today. In other words, if DE is a cosmological constant, it is unlikely that future cosmological observations will be able to discriminate between ST and GR.  The question we intend to explicitly tackle is the following: taking all the abovementioned restrictions (BBN, CMB, GR tests) into account, how much can the Hubble parameter $\\tilde{H}$ at the time of WIMP freeze-out differ from its canonical value if ST replaces GR?  In other words, how much is the WIMP relic density allowed to vary, if we consider ST instead of GR?  We find that in ST theories the expansion rate of the Universe at few GeVs can profoudly differ from the usual value obtained in GR (with variations up to five orders of magnitude) and, yet, allow the correct light elements production at BBN.  This situation is perfectly analogous to the `kination' effect studied in \\cite{salatirosati}, where a modification of $\\tilde{H}$ at WIMP freeze-out was induced by a short period of dominance of the scalar kinetic energy, although with some deal of fine-tuning.  In the case considered here, the effect of ST on $\\tilde{H}$ depends on the strength of the scalar-matter coupling, however no particular fine-tuning is needed to pass the severe nucleosynthesis test even when large modifications of $\\tilde{H}$ at freeze-out occur. This means that the attraction of ST towards GR proceeds very rapidly during the cooling of the Universe from the few GeVs of WIMPs freeze-out down to the MeV range of nucleosynthesis. The overlap of ST with GR can subsequently be very efficient leading to ST scenarios which can hardly be disentangled from ordinary GR in present tests at the post-newtonian level. The fact that ST strongly affects the number of CDM particles may turn out to be the major signature of these theories.  >From the point of view of particle physics model building, these large variations in the WIMPs number density today is of utmost relevance. Particles which were not considered suitable to play a significant role in CDM scenarios can be rescued because of their enhanced number density.  On the other hand, particles (or regions of the parameter space for certain WIMPs candidates), which in usual GR scenarios constitute promising CDM candidates, would be excluded because their boosted number would overclose the Universe. These considerations become of particular interest if we focus on the case where the WIMPs correspond to the lightest supersymmetric particle. A complete analysis of the cosmologically excluded and cosmologically interesting regions of the SUSY parameter spaces in different SUSY contexts, when ST is considered, is presently in progress \\cite{inprogress}. ",
        "conclusions": "The idea of exploiting primordial (ultralight) scalars in order to shed some light on a dynamical interpretation of DE is by now a widespread research topic in the literature. In this paper we follow the promising proposal of considering the quintessence scalar as embedded in a scalar--tensor theory of gravity.  This approach is at variance with the usual interpretation of quintessence as a new light scalar whose interactions with matter are subject to the tight phenomenological constraints on the equivalence principle violation and time-variation of the fundamental coupling constants. Identifying the quintessence field with the scalar component of a ST theory, instead, does not pose any threat on the equivalence principle, since by construction matter has a purely metric coupling with gravity.  We focus on quintessence ST models which possess a double ``attraction mechanism'', one to GR and the other ensuring $\\rho_{DE}$ to follow a tracking solution. These two simultaneous mechanisms act as a ``protection'' for the theory to prevent its fall into immediate troubles (for instance, large departures from GR predictions). Nevertheless, we still obtain important phnomenological signatures which might disentangle this theory from GR or other alternative proposals. The tests of our ST scenario divide  into two classes: deviations from GR and departures from standard cosmology, in particular concerning the expansion rate of the Universe.  The latter effect has a big impact on the most distant epoch of the Universe for which we have ``direct'' information, i.e. nucleosynthesis. However, we pointed out that we can further extend the implications of a non-standard $\\tilde{H}$ in the early Universe to times prior to nucleosynthesis.  Sticking to the standard WIMP picture of DM, one of the most relevant events before BBN is the WIMP decoupling which is expected to have occurred at a temperature of a few GeVs. Our work shows that, despite the severe ``filters'' on ST quintessence models which are provided by BBN, solar system tests of gravity and, to a lesser degree, by CMB, it is still possible to find remarkable enhancements on the expansion rate of the Universe at  WIMP freeze-out, yielding to relic WIMP abundances which can vary up to a few orders of magnitude with respect to the standard case.  In this paper we pointed out  some general features of the new ``WIMP story''  around its decoupling temperature in the presence of ST quintessence. In particular, we noticed that some unexpected effect can take place, such as a short phase of WIMP ``re-annihilation'' when ST approaches GR. Needless to say, such potentially (very) large deviations entail new prospects on the WIMP characterization both for the choice of the CDM candidates and  for their direct and indirect detection probes. A thorough reconsideration of the ``traditional'' WIMP identified with the lightest neutralino in SUSY extensions of the SM as well as the identification of other potentially viable CDM candidates in the ST context  is presently under way \\cite{inprogress}."
    },
    "0403/astro-ph0403422_arXiv.txt": {
        "abstract": "We report a possible detection of TeV gamma-rays from the Galactic Center by the Whipple 10m gamma-ray telescope.  Twenty-six hours of data were taken over an extended period from 1995 through 2003 resulting in a total significance of 3.7 standard deviations.  The measured excess corresponds to an integral flux of \\Flux above an energy of $2.8\\ \\mathrm{TeV}$, roughly 40\\% of the flux from the Crab Nebula at this energy.  The 95\\% confidence region has an angular extent of about 15 arcmin and includes the position of Sgr A*. The detection is consistent with a point source and shows no evidence for variability. ",
        "introduction": "The central region of our galaxy is now thought to contain a super-massive black-hole of 2.6$\\times 10^6 M_\\odot$ \\citep{2002AAS...201.6804G,2002Natur.419..694S} coincident with the unresolved radio source Sgr A* \\citep{balick1974}. Chandra observations reveal X-ray emission from an unresolved point source as well as an extended structure ($\\sim 1.5 \\ \\mathrm{arcsec}$), both of which appear to be physically associated with Sgr A* (e.g., \\citet{2003ApJ...591..891B}).  The recent discovery of hour-scale X-ray \\citep{2001Natur.413...45B} and rapid IR flaring \\citep{2004ApJ...601L.159G} point to an active nucleus, albeit with very low bolometric luminosity compared with the luminosity inferred from the Bondi accretion rate or with that which is typical of more powerful AGNs.  More recently, \\emph{INTEGRAL} (\\emph{the International Gamma-Ray Astrophysics Laboratory}) has detected time-variable 20-100 keV emission from within 0.$^{'}$9 of Sgr A*\\citep{2004ApJ...601L.163B}.  Polarization measurements show the signature of synchrotron radiation in a Keplarian accretion disk \\citep{2002ApJ...566L..77L}. Taken together, these multi-wavelength data are not easily described by a one-component model, and the current theoretical framework combines thermal emission from a radiatively inefficient Keplarian accretion flow with synchrotron inverse-Compton emission produced either by electrons accelerated in the disk or further out in a hypothetical jet-like outflow (e.g. \\citet{2002ApJ...566L..77L,2002A&A...383..854Y,2003astro.ph..4125Y}). From the present measurements, the maximum energy of the non-thermal electron distribution in the jet models is ambiguous and theories alternately explain the high-energy emission as inverse-Compton or the high-energy extension of the synchrotron spectrum; gamma-ray measurements may eventually break this degeneracy.  The EGRET experiment detected a strong unidentified source of GeV gamma-rays marginally consistent with the position of the Galactic Center \\citep{1999ApJS..123...79H}.  Both the Whipple and Cangaroo groups have presented preliminary evidence for TeV emission at the position of Sgr A* as well \\citep{buckley97, cangaroo2003, kosack2003}.  \\citet{dingus2002} re-analyzed the higher energy gamma-ray data from EGRET and found that the most likely position of the EGRET source may be offset from Sgr A*.  However, systematic uncertainties in the gamma-ray background models and limited angular resolution make the analysis of the source in the Galactic Center region difficult. Observations of the Galactic Center are complicated since Sgr A* is surrounded by a dense cluster of stars and stellar remnants (including low-mass X-ray binaries and black hole candidates), molecular clouds, and a large structure that may be the remnant of a powerful supernova remnant, Sgr A East \\citep{2003ApJ...596.1035F}.  Source confusion is particularly difficult for high-energy gamma-ray observations given the limited angular resolution of present experiments.   High energy gamma-ray observations of the Galactic Center are also the subject of particular theoretical interest given the possibility of detecting halo dark matter in our galaxy (e.g. \\citet{1998APh.....9..137B}).  Sgr A*, at the dynamical center of our galaxy, may well be surrounded by a cusp or spike in the dark matter halo distribution (e.g.  \\citet{1991ApJ...378..496D, 1996ApJ...462..563N,1999PhRvL..83.1719G, 2003astro.ph.11594M}). Annihilation of these hypothetical weakly interacting massive particles could also contribute to the luminosity in the vicinity of Sgr A* in the radio through gamma-ray waveband.  Annihilation of dark matter would be enhanced by a factor proportional to the density squared, and might result in an observable gamma-ray line (from direct annihilation to gamma rays) as well as continuum emission (from secondary products of annihilation to quarks and fermions) \\citep{1987ApJ...313L..47S, 1989pafe.conf..255B, 1989PhRvD..39.3549R, 1989PhRvD..40.2549G,1989ApJ...343..169S,1995PhRvD..51.3121J}. The presence of a massive black hole could further steepen the density profile of the dark matter halo, producing very high radio and gamma-ray fluxes that exceed the observational upper bounds \\citep{1999PhRvL..83.1719G}. The details of the halo model on scales $<100\\mathrm{pc}$ and the formation history of the central black hole are critical to predicting the gamma-ray flux, but are, unfortunately, still poorly understood.  Given the limited angular resolution of GeV and TeV instruments, a number of different sources could contribute to a signal near the Galactic Center. The key to distinguishing between all of the possible emission scenarios is to measure the position, angular extent, variability and spectrum of the gamma-ray signal. Here we present first results from an analysis of Whipple telescope data. In \\S\\ref{sec:method} we describe the observational method and data analysis procedure used to observe the Galactic Center at TeV energies. In \\S\\ref{sec:analysis} and \\S\\ref{sec:results}, we discuss a possible weak detection and consider its impact on various gamma-ray production scenarios in \\S\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  The TeV excess observed near the position of the Galactic Center is unlikely to have occurred by chance and constitutes a probable, as yet unconfirmed, detection of a new TeV source.  Possible systematics that could contribute to a false detection include the effects of additional noise from the relatively bright off-source region.  While we have largely corrected for these effects, some systematic uncertainties remain.  We have taken into account trials factors by formulating an explicit \\emph{a priori} hypothesis that we would only look for emission at the exact position of the Galactic Center after a pointing correction was applied. Statistical variations in the analysis method (due to the addition of simulated noise in padding) have been taken into account by repeating the analysis ten times, and taking the average significance, giving a conservative estimate of $3.7\\sigma$ for the detection significance.  The lack of significant variability in our data makes it difficult to uniquely identify the source with a compact point source such as Sgr A*, but inspires some confidence in the stability of our observations at large zenith angle. Note that the analysis procedure was designed to mitigate against changes in the count rate due to variations in the instrument. The same ISU simulation package was used here to analyze Whipple observations of the Crab Nebula giving a spectrum in good agreement which that measured a small zenith angle \\citep{1999ApJ...511..149K}.  In the past, our group reported a positive excess of $2.4\\sigma$ for 1995-1997 observations \\citep{buckley97} and $2.4\\sigma$ for 1999-2003 observations \\citep{kosack2003} at the position of Sgr A*.  The combined significance for our refined results is consistent with these earlier analyses.  The large error circles for both EGRET (7.2 arcmin) and Whipple (15 arcmin) observations make identification with a particular source difficult, but given the dearth of TeV sources, an accidental angular coincidence of a new source along the line of sight is unlikely, and it is probable that the emission comes from a non-thermal source physically near the Galactic Center.  The high level of emission $\\approx 0.4\\ \\mathrm{Crab}$ at a distance of roughly four times that of the Crab Nebula, qualifies this as an unusually luminous galactic source.  Previous TeV observations of relatively nearby galactic sources such as X-ray binaries and shell-type and plerionic supernovae have produced numerous upper-limits, or (at best) unconfirmed detections, making the detection of such an object at 8.5kpc even more unlikely. \\citet{1998A&A...335..161M} came to a similar conclusion about the GeV emission based on the high luminosity of the EGRET unidentified source and lack of significant variability. If the Sgr A East supernova shock were the source of the EGRET gamma rays, it would have been an unusually intense explosion \\citep{1996ApJ...457L..61K} and a density of $1000\\ \\mathrm{cm^{-3}}$ and magnetic field of $B \\sim 0.18 \\mathrm{mG}$ (well above the canonical values) would be required \\citep{2003ApJ...596.1035F}. While a typical galactic source such as an SNR, pulsar, or stellar mass black hole is unlikely, an active nucleus at our Galactic Center is still a viable possibility and the detection of correlated variability in future gamma-ray and X-ray observations is required to make the identification with Sgr A* compelling.  If we associate this emission with either the super-massive black hole Sgr A* or supernova remnant Sgr A East, the observed emission could come from self-Compton scattering by electrons with energies up to at least 2.8 TeV or from pion-decay gamma-rays from primary protons of even higher energy (kinematics require their energy to be at least several times the maximum gamma-ray energy).  The lack of significant variability and the consistency with the Galactic Center position allow more exotic possibilities such as the annihilation of very high mass ($> 2\\ \\mathrm{TeV}$) dark matter particles at the Galactic Center.  While not particularly constraining, these results and a more detailed analysis of the spectrum will be used to derive upper limits for dark matter annihilation in a subsequent paper.  In summary, based on 26 hours of data taken with the Whipple 10m gamma-ray observatory, we report a probable detection of a gamma-ray point source consistent with the position of the Galactic Center. We also describe a modified analysis procedure that we have developed to analyze gamma-ray data taken by Imaging Atmospheric \\Cerenkov telescopes operating at large zenith angles. The continued observation of TeV gamma rays from the Galactic Center has important theoretical implications for understanding a variety of astrophysical phenomena in the nuclear region of the Milky Way.  \\paragraph{Acknowledgements} {\\small We would like to thank Fluvio Melia, Paolo Gondolo, Johnathan Katz, and Ramanath Cowsik for useful discussion. The VERITAS Collaboration is supported by the U.S. Dept. of Energy, N.S.F., the Smithsonian Institution, P.P.A.R.C. (U.K.) and Enterprise-Ireland.}"
    },
    "0403/astro-ph0403087_arXiv.txt": {
        "abstract": "I  examine  the  possibility  that  the Cold  Neutral  Medium  of  the interstellar medium  in galaxies contains  a molecular gas  phase that may represent a  significant and even the dominant  amount of its mass in metal-poor  regions. In spiral  galaxies such regions are  found at large  galactocentric  distances  where  diffuse  H$_2$  gas  will  be untraceable through  its feeble $  ^{12}$CO J=1--0 emission,  the very lack  of it  being also  responsible its  higher  kinetic temperatures ($\\rm T_k\\sim 60-100$ K).  Sensitive sub-mm imaging of spiral galaxies has  demonstrated the  existence  of  dust well  inside  their HI  gas distribution  while recent  observational work  suggests a  high H$_2$ formation rate from HI association onto grains. The latter is indeed a critical  unknown and  a  high  value can  easily  compensate for  the reduction   of  the   available  grain   surface  in   the  metal-poor environments giving  rise to a  CO-poor, diffuse H$_2$ phase  that may then  contribute  significantly  to  the  mass  and  pressure  of  the interstellar medium in such environments. ",
        "introduction": "The use of CO rotational lines to detect the bulk of H$_2$ gas mass in galaxies and deduce its  physical properties is now a well-established technique (e.g.   Dickman, Snell, \\& Schloerb 1986;  Young \\& Scoville 1991) employed  successfully  in the  local  Universe (e.g.   Sanders, Scoville  \\& Soifer  1991; Downes  \\& Solomon  1998) as  well  as high redshifts  (e.g.  Brown  \\& vanden  Bout  1991; Omont  et al.   1996). Nevertheless early  studies have indicated serious  limitations of the method in  tracing metal-poor H$_2$  gas (e.g. Maloney \\&  Black 1988; Israel  1997).    These  become  particularly   severe  in  metal-poor environments  where the  reduced dust-shielding  allows UV  photons to dissociate the  CO molecule  while leaving the  largely self-shielding H$_2$ intact.  This is now corroborated by firm observational evidence for the metal-poor outer parts of typical spirals (Nakai \\& Kuno 1995; Arimoto  et  al.   1996),  in  globally metal-poor  objects  like  the Magellanic Clouds  and Magellanic  Irregulars (Israel 1997;  Madden et al.  1997),  and blue  compact dwarf galaxies  (Barone et  al.  2000). Interestingly  in  many  such  environments pressure  ``probes''  like diffuse CO-bright  molecular clouds (Heyer, Carpenter,  \\& Snell 2001) or HII regions (Rudolph et al.  1996; Elmegreen, \\& Hunter 2000) imply unexpectedly  high pressures  which  then suggest  the  presence of  a pressure agent other than HI or CO-bright H$_2$ gas. ",
        "conclusions": "The Cold  Neutral Medium  in galaxies may  contain an H$_2$  gas phase with  similar  temperatures and  densities  that  can be  particularly prominent  in their  metal-poor regions.   Its spatial  extent depends sensitively on  the value of the  H$_2$ formation rate  on dust grains and the  most promising  tool for revealing  its mass  distribution is through the emission of  its optically thin S(0) $\\rm J_{u}-J_{l}=2-0$ line."
    },
    "0403/astro-ph0403368_arXiv.txt": {
        "abstract": "\\noindent We present submillimetre (SCUBA) photometry of a sample of eight high-redshift ($2.5 \\leq z < 3.5$) radio quasars from two redshift surveys: the {\\sc TexOx}-1000, or TOOT, survey; and the 7C Redshift Survey (7CRS). Unlike the powerful high-redshift radio sources observed previously in the submillimetre, these radio sources are typical of those dominating the radio luminosity density of the population. We detect just two of the TOOT/7CRS targets at 850 $\\mu\\rm{m}$, and one of these detections is probably due to synchrotron emission rather than dust. The population represented by the other six objects is detected in a statistical sense with their average 850 $\\mu\\rm{m}$ flux density implying that they are similar to low-redshift, far-infrared-luminous quasars undergoing at most moderate ($\\ltsimeq 200 ~ \\rm M_{\\odot} ~ yr^{-1}$) starbursts. By considering all the SCUBA data available for radio sources, we conclude that positive correlations between rest-frame far-infrared luminosity $L_{\\rm FIR}$, 151-MHz luminosity $L_{151}$ and redshift $z$, although likely to be present, are hard to interpret because of subtle selection and classification biases, small number statistics and uncertainties concerning synchrotron contamination and $K$-correction. We argue that there is not yet any compelling evidence for significant differences in the submillimetre properties of radio-loud and radio-quiet quasars at high redshift. ",
        "introduction": "The advent of the Submillimetre Common-User Bolometer Array (SCUBA; Holland et al.\\ 1999) camera on the James Clerk Maxwell Telescope (JCMT) opened up the field of submillimetre cosmology. This, and other, instrument/telescope combinations now routinely detect dusty star-forming galaxies at high redshifts, a process greatly aided by the way in which the steep submillimetre spectral index compensates for distance-dependent dimming, so that two objects of similar far-infrared luminosity have similar observed submillimetre flux density if both lie anywhere in the redshift range $1 \\ltsimeq z \\ltsimeq 10$ (Blain \\& Longair 1993).  A major problem in understanding how submillimetre surveys constrain the history of star formation in the Universe is the difficulty of measuring redshifts for the sources of the submillimetre emission. This problem has not yet been solved, with only small ($< 10$ per cent) fractions of objects in submillimetre-flux-density-limited samples having reliable spectroscopic redshifts (e.g. Blain et al.\\ 2002).  An alternative strategy for studying the cosmic star-formation history is to make targeted submillimetre observations of objects with known redshifts (e.g.\\ Hughes, Dunlop \\& Rawlings 1997). Powerful radio sources are important targets, not because the amount of star formation associated with them seems likely to be cosmologically significant, but because of two crucial properties: (i) they appear to be associated with galaxies, and also presumably black holes, with a mass distribution with a large mean and a small dispersion (e.g.\\ Jarvis et al.\\ 2001b); and (ii) studies of the physics of their extended radio structures allows the calculation of the time since their jet-producing AGN were first triggered (e.g.\\ Blundell \\& Rawlings 1999). The first property is potentially very useful for determining how physical properties derived from submillimetre observations (e.g.\\ dust mass and star-formation rate) depend on galaxy, or black hole, mass. The second property can be used to constrain models in which starburst and AGN activity are physically related.  The first large study of the submillimetre emission of high-redshift radio sources was made by Archibald et al.\\ (2001). They used SCUBA on the JCMT to make sensitive (to an 850 $\\mu$m noise level $\\approx 1$ mJy rms) observations of 47 radio galaxies covering the redshift range $1<z<4.5$. The chief result of their study was a rise in the detection rates of radio sources from $\\approx 15$ per cent at $z \\ltsimeq 2.5$ to $\\gtsimeq 75$ per cent at higher redshifts. They interpreted their results as indicating that the submillimetre luminosity of radio galaxies, and hence physical quantities like dust mass, increases systematically with redshift.  One important limitation of the Archibald et al.\\ (2001) study was that it targeted only radio galaxies and excluded radio quasars\\footnotemark. This was largely a practical decision, because it was thought that the bright radio-emitting cores of quasars might produce troublesome synchrotron contamination. The decision could be defended on the basis of the simplest unified scheme (e.g.\\ Antonucci 1993) in which no differences are predicted in the joint distributions of isotropically-emitted quantities like synchrotron low-frequency-radio and dust submillimetre emission. A follow-up study of the submillimetre properties of powerful radio quasars at $z \\sim 1.5$ (Willott et al.\\ 2002a) revealed a significant difference between the submillimetre luminosities of radio-luminosity-matched quasars and galaxies -- the quasars are brighter in the submillimetre by a factor $\\geq 2$, accounting for synchrotron contamination. Willott et al.\\ show that their result can still be accommodated within a unified scheme provided one allows for systematic changes in the properties of the obscuring material with time after the jet-triggering event (e.g.\\ Blundell \\& Rawlings 1999; Hirst, Jackson \\& Rawlings 2003) and/or correlations between quasar luminosity and the probability of viewing a given object as a quasar (e.g.\\ as in the `receding torus' model; Lawrence 1991, Simpson 1998).  A second important limitation of the Archibald et al.\\ (2001) study was that it was confined only to the most radio-luminous sources: such objects contribute marginally to the luminosity density of the radio source population and could plausibly be entirely unrepresentative. The main aim of this paper is to establish the submillimetre properties of typical high-redshift radio sources.  \\footnotetext{ We will use the term radio quasar to refer to lobe-dominated steep-spectrum radio sources associated with quasar nuclei. These are the dominant population of radio-loud quasars in the low-frequency-selected samples considered in this paper, whereas compact flat-spectrum quasars can dominate in high-frequency-selected samples. Unified schemes (e.g.\\ Antonucci 1993) invoke an optically-thick obscuring torus which hides the quasar nucleus when the jet axis makes a large angle with the line-of-sight, and such objects are termed radio galaxies. }  In Sec.~\\ref{sec:sample} we describe how a sample of radio sources was selected for SCUBA follow-up, and how these observations were made. The results of these observations are presented and discussed. In Sec.~\\ref{sec:toot} we use these results to establish the submillimetre properties of typical high-$z$ radio sources, to look for evidence of correlations between far-IR luminosity, radio luminosity and redshift and to compare and contrast these results with those obtained previously by Archibald et al.\\ (2001) and Willott et al.\\ (2002a). Some concluding remarks are made in Sec.~\\ref{sec:conclusions}.  We use J2000.0 positions throughout. The convention for all spectral indices, $\\alpha$, is that flux density $S_{\\nu} \\propto \\nu^{-\\alpha}$, where $\\nu$ is the observing frequency. We assume throughout a low-density, $\\Lambda$-dominated Universe in which $H_{\\circ}=70~ {\\rm km~s^{-1}Mpc^{-1}}$, $\\Omega_{\\rm M}=0.3$ and $\\Omega_ {\\Lambda}=0.7$. ",
        "conclusions": "\\label{sec:conclusions}  We have seen in Sec.~\\ref{sec:correlationz} that evidence for evolution of the dust properties of radio sources with redshift emerges only with the adoption of the synchrotron-corrected values for $S_{850}$, and is confused by the influence of the correlation between $L_{\\rm FIR}$ and $L_{151}$ (Sec.~\\ref{sec:correlationl151}). Since these corrections and this correlation are very uncertain, we conclude that, for radio sources, direct observational evidence that $L_{\\rm FIR}$ rises systematically with redshift is not yet compelling. Although it would be rather surprising if an increase in the dust and gas masses of active galaxies with redshift were not an important part of the overall picture (Archibald et al.\\ 2001), the results of Sec.~\\ref{sec:correlation} are now in reasonable accord with the results obtained on radio-quiet quasars. Priddey et al.\\ (2003) noted that variations in the submillimetre detectability of optically-selected quasars with redshift are consistent with the K-correction of a template spectrum which does not evolve in normalisation with redshift. Although, submillimetre studies of radio-quiet quasars do not suffer from the problems of synchrotron contamination, it is interesting to note that, for radio-quiet quasars, searches for systematic changes in dust emission with redshift are confused by possible correlations between $L_{\\rm FIR}$ and optical luminosity (Willott, Rawlings \\& Grimes 2003; Priddey et al.\\ 2003).  In Sec.~\\ref{sec:correlation} we have seen that hyperluminous infrared objects contribute substantially to the average values of $S_{850}$ used in the correlation studies. Looking at the location in Fig.~\\ref{fig:pd} of the radio-loud objects with associated hyperluminous infrared galaxies hints at complicated physical processes behind any correlations. As has been shown previously by Willott et al.\\ (2002a), and was conjectured by Blundell \\& Rawlings (1999), there seems to be a clear association of hyperluminous infrared galaxies with the radio sources of smallest projected linear size $D$. The underlying physics here (see caption of Fig.~\\ref{fig:pd} for details) may be that a short time $\\tau \\ltsimeq 10^{7} ~ \\rm yr$ since the jet-triggering event is a necessary, if not sufficient, condition for far-infrared hyperluminosity as at later times dust may have been destroyed or dispersed by the expanding radio source (De Young 1998). With two important exceptions, the hyperluminous objects appear to be concentrated amongst the most radio-luminous sources. The underlying physics here may be that extreme radio luminosity is strongly linked to the presence of an extreme-luminosity quasar nucleus which in turn promotes, perhaps by direct heating of dust, infrared hyperluminosity. The two hyperluminous objects at relatively low $L_{151}$ (6C1045+3513 and TOOT08\\_061) both show, unusually for radio-loud objects, evidence for broad absorption lines (Willott et al.\\ 2002a; Rawlings et al., 2004). The underlying physics here may be that young bursts of less-than-extreme radio luminosity, probably associated with the early evolution of less powerful jets (e.g.\\ Kuncic 1999), as well as BAL features may be indicative of processes triggered by galaxy mergers (e.g.\\ Canalizo \\& Stockton 2001).  \\begin{figure*} \\begin{center} \\setlength{\\unitlength}{1mm} \\begin{picture}(150,150) \\put(-50,-20){\\special {psfile=\"fig4.ps\" vscale=80 hscale=80angle=0}} \\end{picture} \\end{center} {\\caption[junk]{\\label{fig:pd} The 151-MHz radio luminosity $L_{151}$ versus projected linear size $D$ plane for various complete redshift surveys of radio sources (surveys and symbols as in Fig.~\\ref{fig:pz}; filled symbols are radio quasars, open symbols are radio galaxies). As in Fig.~\\ref{fig:pz}, large rings represent confirmed hyperluminous infrared galaxies and smaller rings represent objects observed with SCUBA, including objects not in the complete redshift surveys. The lines are calculated from the radio source model of Willott et al.\\ (1999) assuming $\\alpha_{\\rm rad}=1$: the dotted lines are contours of constant $\\log_{10}(\\tau / \\rm yr)$, where $\\tau$ is the time since the jet-triggering event, and the dashed lines are contours of constant $\\log_{10}(Q / \\rm W)$, where $Q$ is the power in both jets. The values adopted for the fixed parameters of the radio source model (see Willott et al.\\ 1999 for definitions) are as follows: $\\theta=60^{\\circ}$; $R_{T}=5$; $f=20$; $k=0$; $\\phi = 90^{\\circ}$; $\\eta = 1$; $\\beta=1.5$ (here $\\beta$ is the slope of the density power law, not the dust emissivity index); $c_{1}=2.3$; $n_{100} = 3000 ~ \\rm e^{-} ~ m^{-3}$; none of these parameters, most crucially those of the environment ($\\beta$ and $n_{100}$), are assumed to vary with cosmic time, and some, most notably $f$, are sufficiently uncertain that values of $Q$ and $\\tau$ should be regarded as order-of-magnitude estimates.  }}  \\end{figure*}"
    },
    "0403/astro-ph0403642_arXiv.txt": {
        "abstract": "A 20-hour VLBI observation with the NRAO VLBA and GBT in the NOAO Bo\\\"otes field reaches an rms noise of 9 microJy per beam at 1.4~GHz. Three sources were detected at 10-milliarcsecond resolution within the GBT primary beam of FWHM 8.6\\arcmin, including the 20-milliJy calibrator and two sub-milliJy sources.  By tapering the visibility data, portions of the VLBA primary beam of FWHM 29\\arcmin\\, were imaged at poorer sensitivity and resolution to yield five further detections.  New developments at JIVE will permit deeper and wider VLBI surveys at full sensitivity and resolution, enabling new types of survey science. ",
        "introduction": "Very Long Baseline Interferometry (VLBI) observations, and VLBI surveys in particular, have contributed significantly to our understanding of active galaxies.  For example, the discovery of superluminal motion in these systems provided the first clue that viewing angle was an important parameter in interpreting and eventually unifying different classes of active galaxies, including radio galaxies and quasars (Urry \\& Padovani 1995).  More recently, VLBI has provided direct evidence for an evolutionary scenario wherein very young compact sources are the precursors of giant extended sources associated with radio galaxies (Taylor et al.\\ 2000). Numerous other advances are summarized in these Proceedings.  At GHz frequencies, VLBI surveys usually target sources distributed across the sky.  The sources are also often pre-selected to be strong ($S_{T} > 200$~mJy) and to exhibit flat spectra (e.g., Beasley et al.\\ 2002).  Such samples are largely dominated by intrinsically luminous active galaxies at moderate redshifts $z\\sim 1-2$.  Self-calibration techniques can be used for such bright sources so calibration overhead is minimal.  But the observation time $t$ per source is reduced by needing to slew the array from source to source.  Also, most observers impose a typical image field of view ($fov$) of $\\sim 100$~milliarcseconds, ensuring that only one source is detected in any given observation time $t$.  Such VLBI surveys are therefore slow and the biases introduced are substantial.  Recent attempts have been made to survey weaker ($S_{T} > 10$~mJy) sources using phase referencing in the switching style.  Such VLBI surveys take advantage of the huge catalog of FIRST sources now available at 1.4~GHz with a resolution of 5\\arcsec\\, (White et al.\\ 1997).  These VLBI surveys are usually localized to a region subtending a few square degrees and including a strong phase calibrator close to weaker targets.  Example surveys are presented in these Proceedings by Gurvits et al.\\ (2004) and by Wrobel et al.\\ (2004).  These surveys are both unbiased as no spectral pre-selection is made and efficient as they minimize array slewing.  But phase referencing in the switching style causes substantial calibration overhead, leaving a typical observation time per source of $t \\sim 10$~minutes.  For the weaker VLBI detections, the dynamic range can be 10 or less.  For the stronger VLBI detections, the dynamic range is often limited by residual phase-referencing errors.  These limitations mean that jets and other secondary features can easily be missed, making morphological classification difficult.  Although a $fov \\sim 1000$~milliarcseconds is required to accommodate FIRST's astrometric accuracy, this is still small enough to ensure only one target per $fov$.  Moreover, the 10-milliJy sources from FIRST still involve essentially the same radio source populations that are targeted in the bright surveys (Ivezi\\'c et al.\\ 2002).  Indeed, at 1.4~GHz most sources slightly stronger than $S_{T} \\sim 1$~mJy are energized by active nuclei (e.g., Condon et al.\\ 2003). What steps can be taken toward studying the more diverse sub-milliJy population at resolutions of, say, 10~milliarcseconds (80~pc at $z \\sim 2$ for $\\Omega_m = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and $H_0 = 70$~km~s$^{-1}$~Mpc$^{-1}$)?  Sources brighter than 0.1~milliJy have an areal density of 0.2 per square arcminute, so about a dozen are expected within the primary beam of a 100-m class antenna.  Therefore, progress can be expected if two technical goals can be reached: achieve wider fields of view and lower detection thresholds, such that many sources can be observed simultaneously with the VLBI array.  For a connected array with short baselines, the $fov$ is often set by the primary-beam response of the antennas in the array.  For VLBI this is rarely the case.  In VLBI, a more demanding limitation is set by the frequency and time sampling used during data correlation.  This sampling must be fine enough to mitigate the effects of both bandwidth and time-average smearing in images.  Also, since preserving $fov$ scales computationally with the square of the baseline length, generating wide-field VLBI images is often limited by the off-line computing resources available.  This latter restriction has introduced a psychological barrier, which is that most VLBI observers over-average their data in both the time and frequency domains to make data sets more manageable.  Data-averaging collapses both the $fov$ and the total information content, leaving us with the postage-stamp VLBI images with which we are all so familiar. ",
        "conclusions": ""
    },
    "0403/astro-ph0403474_arXiv.txt": {
        "abstract": "We review a range of stastistical methods for analyzing the structures of star clusters, and derive a new measure ${\\cal Q}$ which both quantifies, and distinguishes between, a (relatively smooth) large-scale radial density gradient and multi-scale (fractal) sub-clustering.  The distribution of separations $p(s)$ is considered, and the Normalised Correlation Length $\\bar{s}$ (i.e. the mean separation between stars, divided by the overall radius of the cluster), is shown to be a robust indicator of the extent to which a smooth cluster is centrally concentrated. For spherical clusters having volume-density $n \\propto r^{-\\alpha}$ (with $\\alpha$ between 0 and 2) $\\bar{s}$ decreases monotonically with $\\alpha$, from $\\sim 0.8$ to $\\sim 0.6$. Since $\\bar{s}$ reflects all star positions, it implicitly incorporates edge effects. However, for fractal star clusters (with fractal dimension $D$ between 1.5 and 3) $\\bar{s}$ decreases monotonically with $D$ (from $\\sim 0.8$ to $\\sim 0.6$). Hence $\\bar{s}$, on its own, can quantify, but cannot distinguish between, a smooth large-scale radial density gradient and multi-scale (fractal) sub-clustering.  The Minimal Spanning Tree (MST) is then considered, and it is shown that the Normalised Mean Edge Length $\\bar{m}$ (i.e. the mean length of the branches of the tree, divided by $({\\cal N}_{\\rm total}A)^{1/2}/({\\cal N}_{\\rm total}-1)$, where $A$ is the area of the cluster and ${\\cal N}_{\\rm total}$ is the number of stars), can also quantify, but again cannot on its own distinguish between, a smooth large-scale radial density gradient and multi-scale (fractal) sub-clustering.  However, the combination ${\\cal Q} ={ \\bar{m}}/{ \\bar{s}} $ does both quantify, and distinguish between, a smooth large-scale radial density gradient and multi-scale (fractal) sub-clustering. IC348 has ${\\cal Q} = 0.98$ and $\\rho$ Ophiuchus has ${\\cal Q} = 0.85$, implying that both are centrally concentrated clusters with, respectively, $\\alpha \\simeq 2.2 \\pm 0.2$ and $\\alpha \\simeq 1.2 \\pm 0.3\\,$. Chamaeleon and IC2391 have ${\\cal Q} =0.67$ and ${\\cal Q} =0.66$ respectively, implying mild substructure with a notional fractal dimension $D \\simeq 2.25 \\pm 0.25$. Taurus has even more sub-structure, with ${\\cal Q} =0.45$ implying $D' \\simeq 1.55 \\pm 0.25$. If the binaries in Taurus are treated as single systems, ${\\cal Q} $ increases to $0.58$ and $D'$ increases to $1.9 \\pm 0.2$. ",
        "introduction": "Since most stars are formed in clusters, it would be useful to have quantitative and objective statistical measures of their structure, with a view to comparing clusters formed in different environments, and tracking changes in structure as clusters evolve. This is particularly important for young, embedded clusters, where the structure may yield important clues to the formation process but is changing rapidly. It is also important for comparing observed clusters with numerical simulations.  At present, we do not have sufficiently robust statistical measures for this purpose. Features which are easily identified by the human eye, such as sub-clusters, or linear features, can be strangely elusive to objective statistical analysis. For example, it is difficult to distinguish, statistically, between a degree of fractal or random sub-clustering, and the existence of a density gradient (Bate, Clarke \\& McCaughrean 1997). This paper explores some possible measures, and evaluates their usefulness. In particular, we find a robust objective measure which both quantifies, and distinguishes between, a smooth large-scale radial density gradient and multi-scale (fractal) sub-clustering.  In Section 2 we describe our methodology. In Section 3 we look again at the Mean Surface Density of Companions (MSDC), a tool pioneered by Larson (1995) and subsequently used by several others (e.g. Simon 1995; Bate, Clarke and McCaughrean 1997; Nakajima et al. 1998; Brandner \\& K\\\"ohler 1998; Gladwin et al. 1999; Klessen \\& Kroupa 2001). We focus on measures which reflect the clustering regime (wide separations) rather than the binary regime (close separations). In Section 4 we explore the use of the Minimal Spanning Tree (Barrow, Bhavasar \\& Sonoda 1985) and its derivatives. In Section 5 we combine the MSDC and the MST to derive a single measure ${\\cal Q}$ which is able both to quantify, and to distinguish between, a smooth radial density gradient and multi-scale (fractal)  sub-clustering.  All the measures are tested and calibrated on multiple realizations of artificial star clusters, and applied to $\\rho$ Ophiuchus, Chamaeleon, Taurus, IC348 and IC2391. Our results are discussed in Section 6, and the main conclusions are summarized in Section 7.      \\section {Methodology}  Three different types of artificial star cluster have been created, using random numbers ${\\cal R}$ to generate the individual star positions. The first type (2D$\\alpha$) are circular clusters (i.e. two-dimensional disc-like clusters) with surface density $N \\propto r^{-\\alpha}$ and $\\alpha = 0 \\;{\\rm or}\\; 1$. The second type (3D$\\alpha$) are spherical clusters (i.e. three-dimensional clusters) having volume density $n \\propto r^{-\\alpha}$ with $\\alpha = 0,\\;1,\\;2,\\;{\\rm and}\\;2.9$. The third type (F$D$) are fractal star clusters (again three dimensional) with fractal dimension $D = 3.0,\\;2.5,\\;2.0,\\;{\\rm or}\\;1.5$.  The different types are listed in Column 1 of Table~\\ref{Measures}. All of the artificial clusters are created with 100 to 300 stars, as the numbers of stars within the  five real clusters lie within that range. The data for the five real clusters used are illustrated in  Appendix~\\ref{RawData}, and the sources listed in Table 2.     \\begin {table}  \\caption{Clustering measures obtained for artificial and real star clusters. Column 1 lists the cluster type (for artificial clusters) or name (for real clusters). Column 2 gives the Normalized Correlation Length $\\bar{s}$ (i.e. the ratio of the mean separation to the cluster radius, see Section 3). Column 3 gives the Normalised Mean Edge Length $\\bar{m}$ (see Section 4). Column 4 gives the mean value of the standard deviation of the edge length, $\\bar{\\sigma}_m$. Column 5 gives ${\\cal Q} = \\bar{m} / \\bar{s}$. For the artificial star clusters, means and standard deviations are computed from 100 realisations of each type, with $100 \\leq {\\cal N}_{total} \\leq 300$.} \\begin {tabular}  [tbp]{lcccc} Cluster type  & & &\\\\ or name&  $\\bar{s}$ & $\\bar{m}$ & $\\bar{\\sigma}_m$ & ${\\cal Q}$\\\\ \\hline & &  & & \\\\ 2D0($N\\!\\propto\\!r^0$)\t & $.88\\pm.03$ &$.65 \\pm0.02$ &$.31\\pm.02$&$.74\\pm.02$\\\\ 2D1($N\\!\\propto\\!r^{-1}$) & $.70\\pm.03$ &$.61 \\pm.02 $ &$.38\\pm.02$&$.85\\pm.03$\\\\  & &  & & \\\\ 3D2.9($n\\!\\propto\\!r^{-2.9}$) &$.16\\pm.02$&$.24\\pm.05$& $.59\\pm.07$&$1.50\\pm.13$\\\\ 3D2($n\\!\\propto\\!r^{-2}$) & $.60\\pm.03$ &$.55 \\pm.02 $& $.41\\pm.03$&$.93\\pm.03$\\\\ 3D1($n\\!\\propto\\!r^{-1}$) & $.73\\pm.03$ &$.61 \\pm.02 $& $.33\\pm.03$&$.84\\pm.02$\\\\ 3D0($n\\!\\propto\\!r^0$) \t & $.80\\pm.02$ &$.63\\pm.02  $& $.31\\pm.02$&$.79\\pm.02$\\\\ & & & & \\\\ F3.0($D\\!=\\!3.0$)& $.81\\pm.03$ \t& $.64 \\pm.02$ \t&$ .30\\pm.02$\t&$.80\\pm.02$\\\\ F2.5($D\\!=\\!2.5$)& $.74\\pm.09$ \t& $.54 \\pm.05$ \t&$ .28\\pm.03$\t&$.73\\pm.06$\\\\ F2.0($D\\!=\\!2.0$)& $.67\\pm.13$ \t& $.41 \\pm.04$ \t&$ .28\\pm.02$\t&$.61\\pm.08$\\\\ F1.5($D\\!=\\!1.5$)& $.62\\pm.18$ \t& $.27 \\pm.07$ \t&$ .35\\pm.07$\t&$.45\\pm.09$\\\\ & & & &\\\\ IC2391\t\t& 0.74 \t\t& .49 \t\t& .30 \t\t& .66\\\\ Chamaeleon\t& 0.63 \t\t& .42 \t\t& .45 \t\t& .67\\\\ Taurus\t\t& 0.55   \t& .26 \t\t& .56 \t\t& .47\\\\ $\\rho$ Ophiuchus  & 0.53 \t& .45 \t\t& .39 \t\t& .85\\\\ IC348\t\t& 0.49 \t\t& .48 \t\t& .46 \t\t& .98\\\\  \\hline \\end {tabular} \\label{Measures}  \\end{table}     A cluster of type 2D$\\alpha$ is created by positioning the stars according to \\begin{equation} \\left. \\begin{array}{lll} r & = & \\left\\{(2-\\alpha){\\cal R}_r/2 \\right\\}^{1/(2-\\alpha)} \\,, \\\\ \\phi & = & 2 \\pi {\\cal R}_\\phi \\,, \\\\ x & = & r \\, {\\rm cos}(\\phi) \\,, \\\\ y & = & r \\, {\\rm sin}(\\phi) \\,. \\\\ \\end{array} \\right\\} \\end{equation} where ${\\cal R}_r$ and ${\\cal R}_\\phi$ are random numbers in the range 0-1.  A cluster of type 3D$\\alpha$ is created by positioning the stars according to \\begin{equation} \\left. \\begin{array}{lll} r & = & \\left\\{(3-\\alpha){\\cal R}_r/3 \\right\\}^{1/(3-\\alpha)} \\,, \\\\ \\theta & = & {\\rm cos}^{-1} \\left( 2 {\\cal R}_\\theta - 1 \\right) \\,, \\\\ \\phi & = & 2 \\pi {\\cal R}_\\phi \\,, \\\\ x & = & r \\, {\\rm sin}(\\theta) \\, {\\rm cos}(\\phi) \\,, \\\\ y & = & r \\, {\\rm sin}(\\theta) \\, {\\rm sin}(\\phi) \\,, \\\\ z & = & r \\, {\\rm cos}(\\theta) \\,. \\\\ \\end{array} \\right\\} \\end{equation} where ${\\cal R}_r$, ${\\cal R}_\\theta$ and ${\\cal R}_\\phi$ are random numbers in the range 0-1. Clearly this method cannot be used for $\\alpha$ = 3, so to have a cluster type approximating to $\\alpha$ = 3 we use $\\alpha$ = 2.9.  A cluster of type F$D$ is created by defining an ur-cube with side 2, and placing an ur-parent at the centre of the ur-cube. Next, the ur-cube is divided into ${\\cal N}_{\\rm div}^3$ equal sub-cubes, and a child is placed at the centre of each sub-cube (the first generation). Normally we use ${\\cal N}_{\\rm div} = 2$, in which case there are 8 sub-cubes and 8 first-generation children. The probability that a child matures to become a parent in its own right is ${\\cal N}_{\\rm div}^{(D-3)}$, where $D$ is the fractal dimension. For lower $D$, the probability that a child matures to become a parent is lower, and the cluster is more `porous'. Children who do not mature are deleted, along with the ur-parent. A little noise is then added to the positions of the remaining children, to avoid an obviously regular structure, and they then become the parents of the next generation, each one spawning ${\\cal N}_{\\rm div}^3$ children (the second generation) at the centres of ${\\cal N}_{\\rm div}^3$ equal-volume sub-sub-cubes, and with each second-generation child having a probability ${\\cal N}_{\\rm div}^{(D-3)}$ of maturing to become a parent. This process is repeated recursively until there is a sufficiently large generation that, even after pruning to impose a spherically symmetric envelope of radius 1 within the ur-cube, there are still more children than the required number of stars. Children are then culled randomly until the required number is left, and the surviving children are identified with the stars of the cluster. At each generation, the survival of a child is determined by generating a random number ${\\cal R}$ in $(0,1)$; survival then requires that ${\\cal R} < {\\cal N}_{\\rm div}^{(D-3)}$.  Clusters of type 2D$\\alpha$ are investigated for two purposes. First, we wish to clarify the effect of a sharply defined circular edge on an otherwise statistically uniform, two-dimensional distribution of stars. Clusters of type 2D0 enable us to isolate this effect. Second, we wish to explore how readily two-dimensional and three-dimensional distributions can be distinguished. This could be important if stars are being formed in layers, for example at a shock front.  For each type of artificial cluster, 100 realisations are analysed, so that means and standard deviations can be obtained for the parameters extracted. Three-dimensional clusters (types 3D$\\alpha$ and F$D$) are projected onto an arbitrary plane prior to analysis. Two-dimensional clusters are viewed face-on.     \\begin{figure*} \\centerline{\\psfig{figure=fig2a.ps,width=6cm} \\psfig{figure=fig2b.ps,width=6cm} \\psfig{figure=fig2c.ps,width=6cm}}  \\centerline{\\psfig{figure=fig2d.ps,width=6cm} \\psfig{figure=fig2e.ps,width=6cm} \\psfig{figure=fig2f.ps,width=6cm}} \\caption{Distribution function $p(s)$ for separations between randomly chosen stars in artificial (non-fractal) cluster of type (a) 2D0, $N \\propto r^0$; (b) 2D1, $N \\propto r^{-1}$; (c) 3D0, $n \\propto r^0$; (d) 3D1, $n \\propto r^{-1}$; (e) 3D2, $n \\propto r^{-2}$; and (f) 3D2.9, $n \\propto r^{-2.9}$. The solid line is the value of $p(s)$ for a star cluster of type 2D0 having an infinite number of stars (Eqn. \\ref{edge-effects}), and is included for reference. The dashed line is $p(s) = 2s$ (see text). $s$ is normalized to the overall radius of the cluster, as described in the text.} \\label{Arti/p(s)} \\end{figure*}      \\begin{figure*} \\centerline{\\psfig{figure=fig3a.ps,width=6cm} \\psfig{figure=fig3b.ps,width=6cm} \\psfig{figure=fig3c.ps,width=6cm}} \\vspace{0.5cm} \\centerline{\\psfig{figure=fig4a.ps,width=6cm} \\psfig{figure=fig4b.ps,width=6cm}} \\caption{Distribution functions for separations between randomly chosen stars in five real star clusters. (a) $\\rho$ Ophiuchus, (b) IC2391, (c) IC348, (d) Taurus, and (e) Chamaeleon. The solid line is the value of $p(s)$ for a star cluster of type 2D0 having an infinite number of stars (Eqn. \\ref{edge-effects}), and is included for reference. In (d) and (e), the solid line represents a smoothed version of the raw data, to show the existence of multiple maxima.} \\label{Real/p(s)} \\end{figure*} ",
        "conclusions": "We have explored two statistical measures for analysing objectively the observed (i.e. projected) structures of star clusters. These measures are based on the Mean Surface Density of Companions (MSDC), and the Minimal Spanning Tree (MST). The measures are $\\bar{s}$, the normalised mean separation between stars, and $\\bar{m}$, the normalised mean edge-length of the MST, both of which are independent of the number of stars in the cluster. For artificial star clusters, created with a smooth large-scale radial density profile ($n \\propto r^{-\\alpha}$), and for artificial star clusters created with sub-structure having fractal dimension $D$, $\\bar{s}$ and $\\bar{m}$ both decrease with increasing $\\alpha$ and/or decreasing $D$ -- but at different rates. Hence a cluster with a radial gradient can be distinguished from one with sub-structure by evaluating ${\\cal Q} = \\bar{m} / \\bar{s}$. For a cluster of  uniform volume-density (i.e. $\\alpha = 0$ and $D = 3.0$), ${\\cal Q} \\simeq 0.80$. If the cluster is made more centrally condensed by increasing $\\alpha$, ${\\cal Q}$ increases monotonically, reaching ${\\cal Q} \\simeq 1.50$ at $\\alpha = 2.9$. Conversely, if the cluster is given sub-structure by reducing $D$, ${\\cal Q}$ decreases monotonically, reaching ${\\cal Q} \\simeq 0.45$ at $D = 1.5$.  On the basis of their ${\\cal Q}$ values, $\\rho$ Ophiuchus and IC348 have radial gradients with $\\alpha \\simeq 1.2 \\pm 0.3$, and $2.2 \\pm 0.2$, respectively. Chamaeleon and IC2391 have sub-structure with notional fractal dimension $D' \\simeq 2.2 \\pm 0.2$. Taurus has even more sub-structure, with $D' \\simeq 1.55 \\pm 0.25$, and if the binaries in Taurus are treated as single systems, $D'$ increases to $1.9 \\pm 0.2$. $D'$ is only a notional fractal dimension, because the integral measures we have defined do not give any indication of whether the sub-structure is hierarchically self-similar. (Indeed, for clusters having only $\\sim$ 200 stars the range of separations is too small to possess hierarchical self-similarity.)"
    },
    "0403/astro-ph0403704_arXiv.txt": {
        "abstract": "Primary and atmospheric cosmic-ray spectra were precisely measured with the BESS-TeV spectrometer. The spectrometer was upgraded from BESS-98 to achieve seven times higher resolution in momentum measurement. We report absolute fluxes of primary protons and helium nuclei in the energy ranges, 1--540~GeV and 1--250~GeV/n, respectively, and absolute flux of atmospheric muons in the momentum range 0.6--400~GeV/$c$. ",
        "introduction": "\\label{sec:introduction} The absolute flux and spectral shape of primary cosmic rays are the basis to discuss the origin and the propagation history of the cosmic rays in the Galaxy. The spectrum is also essential as an input to calculate spectra of cosmic-ray antiprotons and positrons which are secondary products of cosmic-ray interactions with the interstellar gas. Recently, the importance of the cosmic-ray spectrum has been emphasized in connection with the study of neutrino oscillation observed in atmospheric neutrinos~\\cite{SuperK}. In contrast with a long-baseline experiment, such as K2K~\\cite{K2Kexp}, the feature of the oscillation study with atmospheric neutrinos lies in its wide energy range. For example, the Super Kamiokande water Cherenkov detector~\\cite{SKdetc} can observe neutrinos from 0.1 to 100~GeV. In order to estimate an accurate flux of atmospheric neutrinos up to around 100~GeV, we need to know the primary cosmic-ray flux well above 100~GeV and the feature of hadronic interactions in that energy region. It is crucial to measure primary cosmic-ray flux by determining the absolute energy of particles. Measurements of atmospheric muon spectrum is also important to check and improve our understanding of hadronic interactions.  The Balloon-borne Experiment with a Superconducting Spectrometer (BESS) \\cite{ASTRMG,BESSd1} has been carried out since 1993 to perform highly sensitive searches for cosmic-ray antiparticles and precise measurements of absolute flux of various cosmic-ray components. The absolute fluxes of primary protons and helium nuclei were precisely measured in the energy ranges, 1--120~GeV and 1--54~GeV/n, respectively, by a balloon observation in 1998~\\cite{BESS98}. The overall uncertainties in the measurements were less than 5~\\% for protons and 10~\\% for helium nuclei. The absolute flux of atmospheric muons was measured from 0.6 to 20~GeV/$c$ at sea level~\\cite{GESS95} and from 0.6 to 100~GeV/$c$ at a mountain altitude~\\cite{MESS99}. The overall uncertainty in the measurements was less than 10~\\%. In order to extend the energy range of the precise measurements of cosmic-ray flux up to higher energy, the BESS-TeV spectrometer was developed. A magnetic rigidity ($R\\equiv pc/Ze$) of an incident particle was accurately measured by a new tracking system. The maximum detectable rigidity (MDR) was significantly improved from 200~GV to 1.4~TV.  Cosmic-ray observations were carried out with the BESS-TeV spectrometer at a balloon altitude and at sea level in 2002. We report precise measurements of absolute fluxes of primary protons and helium nuclei, and atmospheric muons.  As to primary cosmic rays, we focused on a energy range above 1~GeV/n in this paper. The BESS spectrometer has measured primary cosmic-ray flux down to 0.2~GeV/n, which provides significant information on the solar modulation effect and the propagation of cosmic rays~\\cite{JZWang}. Measurements of the low-energy cosmic-ray flux with the BESS-TeV spectrometer will be reported elsewhere together with a series of BESS balloon-flight data~\\cite{Shikaz}. ",
        "conclusions": "\\label{sec:conclusion}  We have measured energy spectra of primary protons in the range 1--540~GeV and helium nuclei in the range 1--250~GeV/n by a balloon observation, and momentum spectrum of atmospheric muons in the range 0.6--400~GeV/$c$ by a ground observation at sea level.  The overall uncertainties were less than 15~\\% for protons, less than 20~\\% for helium nuclei, and less then 20~\\% for muons. Primary cosmic-ray spectra provide fundamental information on the origin and propagation history of the cosmic rays. The results also provide accurate input spectra to predict an atmospheric neutrino flux. The result of atmospheric muon spectrum will improve the accuracy of atmospheric neutrino calculation in the wide momentum range.  \\begin{ack} The authors thank NASA and the National Scientific Balloon Facility for their professional and skillful work in carrying out the BESS flight. They deeply thank ISAS and KEK for their continuous support and encouragement for the BESS experiment. They would especially like to thank Doctor T. Taniguchi for his kind advice and cooperation to develop the drift chambers. This experiment was supported by Grant-in-Aid for Scientific Research (12047206 and 12047227) from the Ministry of Education, Culture, Sport, Science and Technology, Japan, and was supported by NASA, USA. \\end{ack}  \\clearpage"
    },
    "0403/astro-ph0403018_arXiv.txt": {
        "abstract": "{ We report on  bursts from the Soft Gamma-Ray Repeater SGR 1806--20 detected with   INTEGRAL  in October 2003, during a period of moderate activity of the source. The spectral and temporal properties of 21 short bursts are consistent with those found in previous observations, even if these bursts are among the faintest observed in the 15-200 keV range from this source. During some of the bursts a clear spectral evolution is visible. The data also show, for the first time, evidence for a  hardness-intensity anti-correlation within SGR 1806--20 bursts. ",
        "introduction": "Soft Gamma-ray Repeaters (SGRs) are a class of peculiar high-energy sources discovered through their recurrent emission of soft $\\gamma$-ray bursts. These bursts have typical durations of $\\sim$0.1 s and luminosities in the range 10$^{39}$-10$^{42}$ ergs s$^{-1}$ (see \\cite{hurleyrew} for a review of this class of objects). The bursting activity and the persistent emission observed in the $\\sim$0.5-10 keV energy range are generally explained in the framework of the ``Magnetar'' model (e.g. \\cite{dt92}, \\cite{td95}), as caused by a highly magnetized ($B\\sim$10$^{15}$ G) slowly rotating ($P\\sim$ 5-8 s) neutron star.   SGR 1806--20 is one of the most active Soft Gamma-ray Repeaters. Here we report new observations of this source obtained with the INTEGRAL satellite in October 2003 during a period of bursting activity (\\cite{gotza}, \\cite{hurley}, \\cite{merec}, \\cite{gotzb}). These data have two advantages compared to previous observations in the soft $\\gamma$-ray energy range of bursts from this source. First, they have been obtained with an imaging instrument, thus we can exclude that the bursts originate from a different source in the field. Second, they have a good sensitivity and time resolution which allows us to study the spectral evolution of relatively faint  bursts.    \\begin{figure*} \\hspace{0cm}\\psfig{figure=Ga282_f1.ps,width=18cm,height=18cm,angle=90} \\caption[]{IBIS/ISGRI background subtracted light curves of the SGR 1806-20 bursts in the 15-100 keV range. Each panel corresponds to a time interval of one second and the time bins are of 10 ms. Units of the axes are time in seconds and vignetting corrected counts per bin. Time 0 corresponds to the the starting time of the $T_{90}$ computation and is reported on top of each panel together with the total number of net counts.} \\label{lc} \\end{figure*} ",
        "conclusions": "\\label{disc}  Previous studies indicated weak  or no spectral evolution for SGR bursts (e.g. \\cite{fenimore}, \\cite{kouveliotou}). Up to now indication for a  hard-to-soft evolution has been reported only for a single burst from \\src  (\\cite{strohmayer}) and for a precursor to a long  burst (3.5 s) from SGR 1900+14 (\\cite{ibrahim}). The same kind of spectral evolution has also been reported for a $\\sim$9 s long burst from SGR 0526-66 (\\cite{golenetskii}): the softening trend is seen in the first three of the four spectra extracted, while the last one is as hard as the first one. In our sample we do not find evidence for this kind of evolution.  On the other hand, soft-to-hard evolution has been seen with the BATSE instrument for two peculiar bursts very likely originating from SGR 1900+14 (\\cite{woods}). These two bursts were quite different from the usual bursts, both in terms of duration (lasting $\\sim$ 1 s), and spectral  hardness ($kT$ of the order of 100 keV).  In the framework of the magnetar model (\\cite{dt92}), short ($\\sim$0.1 s) SGR bursts are usually  described as the radiation originating from the cooling of an optically thick pair-photon plasma. This plasma is generated in the neutron star  magnetosphere by an Alfv\\'{e}n pulse, which is triggered by a sudden shift in the magnetospheric footpoints driven by a fracture in the neutron star crust (\\cite{td95}). This model is able to explain the time histories and energetics of the typical SGR bursts, and predicts that the effective temperature of the spectra should vary weakly during the bursts, owing to the constancy of the photospheric energy flux. No detailed predictions are available concerning more complex spectral evolution as we observe in some of the bursts. For example the model does not account for the presence of a soft precursor as observed in burst number 2 (see Fig. \\ref{trehr}).  Our results  indicate that a hardness-intensity anti-correlation (which in many bursts manifests itself as a soft-to-hard time evolution) is present in  bursts from \\src which are not particularly long, nor spectrally hard and not at all very energetic. It is interesting to note that this correlation is opposite to what found for the bursts emitted from 1E 2259+586 (\\cite{gavrill}), which although have lower fluences than the ones we measure.  \\begin{figure} \\hspace{0cm}\\psfig{figure=Ga282_f3.ps,angle=90,width=8.5cm} \\caption[]{Hardness-Intensity plot of the time resolved hardness ratios of the 12 bursts with the best statistics. The hardness ratio is defined as $(H-S)/(H+S)$, where $H$ and $S$ are the background subtracted counts in the ranges 40-100 keV and 15-40 keV respectively. The count rates are  corrected for the vignetting and refer to the 15-100 keV range. } \\label{hi} \\end{figure}"
    },
    "0403/astro-ph0403532_arXiv.txt": {
        "abstract": "{We present the broadband phase averaged spectrum of one of the brightest intermediate polars V1223 Sgr, obtained with INTEGRAL and RXTE observatories (3-100 keV). Good statistical quality of the spectrum in a hard X-ray energy band (INTEGRAL/IBIS and RXTE/HEXTE) allowed us to disentangle contributions of a direct optically thin plasma emission and a reflected component to the spectrum of V1223 Sgr. The obtained measurement of the post-shock temperature of the accreting matter give us the information about the mass of the white dwarf and the inclination of the system. ",
        "introduction": "V1223 Sgr is a very well known binary star that belongs to the class of intermediate polars (IPs) -- binary systems with a white dwarf that accrets matter from a Roche lobe filling secondary star (see e.g. comprehensive model and discussion in \\cite{beuermann04}). In the IP systems magnetic field of a white dwarf is strong enough to control the accreting matter at relatively large distances from it and to form a magnetosphere and magnetic funnel (see e.g. \\cite{patterson94} for a review). The matter falls down to the surface of the white dwarf with almost free fall velocity and is heated in a strong shock to temperatures of the order of $\\sim 10$ keV. Falling matter preferentially comes to the magnetic poles and emits X-rays via an optically thin thermal plasma emission (see e.g. \\cite{lamb79}). In spite of this relatively simple scheme the observed emission of IP is much more complicated. First of all its spectrum can contain a set of components with different temperatures (mostly visible at low energies), it often demonstrates an intrinsic photoabsorption, partial covering, different ionization states of an emitting plasma and a reflection from an optically thick relatively medium (see e.g. \\cite{warner95} for a review). Often the spectrum depends on orbital phase and phase of rotation of the white dwarf (e.g. \\cite{norton89}). Different objects of this class were extensively studied by several X-ray observatories over last decades. However, until now there was no study of hard X-ray emission of IP ($E>20$ keV) except for the case of V709 Cas (de Martino et al. 2001).  Energy band $>$ 20 keV is important because it allows to confidently measure the maximal temperature of the shock in the accretion column of the IP. It is especially important if the X-ray spectrum of IP contains strong reflected component, that distorts the observed continuum at energies $\\sim$10-20 keV. In the most cases the temperature of the post-shock region of the accretion column of IP is of the order of $\\sim 10-40$ keV that makes it quite difficult to measure with telescopes working at 0.5-10 keV energy band, like EINSTEIN, ASCA, CHANDRA, XMM. Observations of IP performed with GINGA/LAC (e.g. Ishida 1991, Ishida et al. 1994) practically do not provide statistically significant measurements at  $E\\sim$20 keV, that makes it hard to confidently measure the temperatures higher than $\\sim 10-20$ keV. As the energetics of the X-ray spectrum of IP is dominated by bremsstrahlung emission with maximal temperatures the accuracy of its measurement is of high importance (see. e.g. discussion in \\cite{beuermann04}).  In this work we present results of measurements of the broadband spectrum of the intermediate polar V1223 Sgr with the RXTE and INTEGRAL observatories with an accent on the averaged hard X-ray part of the spectrum. ",
        "conclusions": "In the previous section we analyzed the broadband phase averaged spectrum of one of the brightest intermediate polar V1223 Sgr. We have shown that even apart from detection of fluorescent Fe line at energy $\\sim 6.4$ keV (see \\cite{beardmore00}) there are strong indications on the presence of the reflection component in the spectrum of V1223 Sgr. Because of this component any measurements of the post-shock temperature of the accretion column near the surface of IP that are based on instruments with an effective energy band lower than $\\sim$20-30 keV are biased towards obtaining higher temperatures. Our spectral modeling for the first time includes hard X-ray energy range (20-100 keV) and provides more confident estimation of the post-shock temperature. Using this value we derive some important parameters of V1223 Sgr.  Current models of a shock in the accretion column of IP assume that the accreted matter is heated up to its virial temperature (see e.g. \\cite{aizu73}, \\cite{fabian76}, \\cite{lamb79}). The heated plasma produces mainly bremsstrahlung with possible contribution from a cyclotron emission at low accretion rates (see e.g. \\cite{fabian76}, \\cite{warner95}).  Our new broadband measurements of the X-ray spectrum of V1223 Sgr provide us a set of observables  with higher confidence than that of previous. They are: 1) the post-shock temperature $T_{\\rm s}=29\\pm2$ keV, 2) the broadband X-ray flux (absorption corrected) $F_{\\rm 0.5-100 keV}=4.6\\times10^{-10}$ erg s$^{-1}$ cm$^{-2}$, that in turn gives us the emission measure of one emitting accretion column $E_{\\rm m, 1} =\\int N_e^2 dV$ = $1.07\\times 10^{13}~4\\pi D^2$ cm$^{-3}$, where $N_e$ -- is the number density of electrons in a hot emitting plasma.  The distance to V1223 Sgr is now known from HST measurement of the parallax of the system, $D=527 ^{+54}_{-43}$ pc (Beuermann et al. 2004). Therefore we can evaluate the isotropic luminosity of one accretion column $L_{\\rm x,1} = 4\\pi D^2 F_{\\rm 0.5-100 keV}$ = $(1.5 \\pm 0.3)  \\times 10^{34}$ erg s$^{-1}$, and the volume emission measure E$_{\\rm m,1} \\approx  3.4 \\times 10^{56}$ cm$^{-3}$.  The shock temperature can give the direct evaluation of a white dwarf mass (see e.g. \\cite{fabian76}, \\cite{warner95})  \\[ T_s = 31.9 \\frac{M_{\\rm WD}}{M_{\\odot}} \\frac{10^9 \\rm cm}{R_{\\rm WD}} \\rm keV. \\]  Using an approximation of the white dwarf mass-radius relation from Nauenberg (1972)  \\[ R_9=0.78\\left[\\left({1.44 M_\\odot\\over{M}}\\right)^{2/3}-\\left({M\\over{1.44 M_\\odot}}\\right)^{2/3}\\right]^{1/2} \\]  where $R_9$ -- is the radius of the white dwarf in the units of $10^9$ cm, we obtain for the mass and radius of the white dwarf in the V1223 Sgr:  \\[ M_{\\rm WD}= 0.71 \\pm 0.03 M_{\\odot} \\] \\[ R_{\\rm WD}=(7.8\\pm0.2)\\times 10^8 \\rm cm \\]  The obtained estimation of the white dwarf mass is slightly smaller than that obtained by Beardmore et al. (2000) because of lower post-shock temperature measured by us and is better consistent with the evaluation of white dwarf mass from optical observations: $M_{\\rm WD}=0.4-0.6 M_\\odot$ (\\cite{penning85}).    The mass function of the secondary star in the system V1223 Sgr is $f(M_2)=0.0026 \\pm 0.0015$ $M_{\\odot}$ (Watts et al. 1985), the mass of the secondary star is $M_{2}$ = 0.4 $M_{\\odot}$ (Penning 1985). Using these values we can obtain the inclination angle of the system $i = 22^\\circ \\pm 5^\\circ$. If we will assume that the mass of the secondary star is lower, e.g. $M_{2}=0.25 M_{\\odot}$ (\\cite{beuermann04}) the estimation of the inclination angle will increase to $i=32^\\circ \\pm 7^\\circ$.   As the X-ray emission is the major channel of energy losses in the post-shock region we can derive the mass accretion rate $\\dot M$  \\[ \\dot M \\approx {3\\over{8}} {\\mu m_{\\rm H} 2 L_{\\rm x,1} \\over{k T_{\\rm s}}} \\] where $\\mu$=0.62 is the mean molecular weight of the accreting matter. Here factor 2 stands in order to account for two emitting magnetic poles of IP. Substituting measured values we obtain $\\dot M \\approx 2.54 \\times 10^{17}$ g s$^{-1}$.  Using $M_{\\rm WD} = 0.71 M_{\\odot}$, $R_{\\rm WD} = 7.8 \\times 10^8$ cm and $E_{\\rm m,1} = 4 \\pi R^2_{\\rm WD} f/2 ~ h_s N_e^2 \\sim 3.4 \\times 10^{56}$ cm$^{-3}$, where $f$ -- is the fraction of the white dwarf surface, occupied by two accretion columns, we can evaluate the height of standing shock $h_{\\rm s}$ and the number density of electrons in the post-shock region $N_e$ following Fabian et al. (1976) and Warner (1995).  \\[ N_e = 3.1 \\times 10^{13} \\dot M_{17} \\left(\\frac{M_{\\rm WD}}{M_{\\odot}}\\right)^{1/2} R_{9}^{-3/2} f^{-1} \\] \\[ h_s = 3.3 \\times 10^{4} f \\,\\, \\dot M_{17} \\left({M_{\\rm WD}\\over{M_\\odot}}\\right)^{3/2} R_9^{1/2} \\rm cm \\] where $\\dot M_{17}$ -- is the mass accretion rate in units of $10^{17}$ g s$^{-1}$. We can derive $N_e=9.6\\times10^{16} f^{-1}_{-3}$ cm$^{-3}$ and $h_{\\rm s}=1.1 \\times 10^{7} f_{-3}$ cm, where $f_{-3}=10^{-3} f$ -- is the fraction of the surface of the white dwarf, occupied by two accretion columns in the units of $10^{-3}$. Note that the derived parameters are similar to those obtained by de Martino et al. (2001) for the intermediate polar V709 Cas."
    },
    "0403/astro-ph0403254_arXiv.txt": {
        "abstract": "{We consider a contribution of microlensing to the X-ray variability of high-redshifted QSOs. Such an effect could be caused by stellar mass objects (SMO) located in a bulge or/and in a halo of this quasar as well as  at cosmological distances between an observer and a quasar. Here, we  not consider microlensing caused by deflectors in our Galaxy  since it is well-known from recent MACHO, EROS and  OGLE observations that the corresponding optical depth for the Galactic halo and the Galactic bulge is lower than $10^{-6}$. Cosmologically distributed gravitational microlenses could be localized in galaxies (or even in bulge or halo of  gravitational macrolenses) or could be distributed in a uniform way. We have analyzed both cases of such distributions. As a result of our analysis, we obtained that the optical depth for microlensing caused by stellar mass objects is usually small for quasar bulge and  quasar halo gravitational microlens distributions ($\\tau\\sim 10^{-4}$). On the other hand, the optical depth for gravitational microlensing caused by cosmologically distributed deflectors could be significant and could reach $10^{-2} - 0.1$ at $z\\sim 2$.  This means that cosmologically distributed deflectors may contribute significantlly  to the X-ray variability of high-redshifted QSOs ($z>2$). Considering that the upper limit of the optical depth ($\\tau\\sim 0.1$) corresponds to the case where dark matter forms cosmologically distributed deflectors,   observations of the X-ray variations of unlensed QSOs can be used for the estimation of the dark matter fraction of microlenses. ",
        "introduction": "The X-ray radiation of Active Galactic Nuclei (AGNs), in the continuum as well as in spectral lines, has rapid and irregular variability (see e.g. \\cite{Mar81,Bar86,Law93,Gre93,Tur99,Wea01,Man02}, etc.). X-ray flux variability has long been known to be a common property of active galactic nuclei (AGNs), e.g. Ariel 5 and HEAO 1 first revealed long-term (days to years) variability in AGNs  and by uninterrupted observations of EXOSAT  rapid (thousands of seconds) variability was also established as common in these sources (see, for example reviews by \\cite{Mush93,Ulrich97} and references therein). X-ray flux variations are observed on timescales from $\\sim$1000 s to years, and amplitude variations of up to an order of magnitude are observed in the $\\sim$ 0.1 - 10 keV band. It was first suggested by Barr \\& Mushotzky (1986) that the flux variation of an AGN is inversely proportional to its luminosity. \\cite{Law93} and  \\cite{Gre93} confirmed the variability-luminosity relationship, finding that the variability amplitude ($\\sigma$) varies with luminosity as $\\sigma=L_X^{-\\beta}$ with $\\beta\\approx0.3$. Recently, \\cite{Man02} analyzed the variability of a sample of 156 radio-quiet quasars taken from the ROSAT archive, considering the trends in variability of the amplitude with luminosity and with redshift. They found that there was evidences for a growth in AGN X-ray variability amplitude towards high redshift ($z$) in the sense that AGNs of the same X-ray luminosity were more variable at $z>2$. They explained the $\\sigma$ {\\it vs.} $z$ trend assuming that the high-redshifted AGNs accreted at a larger fraction of the Eddington limit than the low-redshifted ones.  On the other hand, the contribution of microlensing to AGN variability was considered in many papers (see e.g. \\cite{Hawk93,Hawk02, Wamb01,Wamb01b,Zakh97}, and references therein). Moreover, recently X-ray microlensing of AGN has been considered \\citep{Popov01a,Tak01,Chart02a,Popovic03,Dai03}. Taking into account that the X-rays of AGNs are generated in the innermost and very compact region of an accretion disc, the X-ray radiation in the continuum as well as in a line can be strongly affected by microlensing \\citep{Popovic03}.\\footnote{Simulations of X-ray line profiles are presented in a number of papers, see, for example, \\cite{Zak_rep02,Zak_rep02a,Zak_rep02_xeus,Zak_rep03_ASR,Zak_rep03_Su} and references therein, in particular \\cite{ZKLR02} showed that an information about magnetic filed may be extracted from X-ray line shape analysis; \\cite{Zak_rep03_AA} discussed signatures of X-ray line shapes for highly inclined accretion disks.}  Recent observations of three lens systems seem to support this idea \\citep{Osh01,Chart02a,Dai03}.  \\cite{Popovic03}  showed that objects in a foreground galaxy with very small masses can cause strong changes in the X-ray line profile.  This  fact may indicate  that the observational probability of X-ray variation due to microlensing events is higher  than in the UV and optical radiation of AGNs. It is connected with the fact that typical sizes of X-ray emission regions are much smaller than typical sizes of those producing optical and UV bands. Typical optical and UV emission region sizes could be comparable or even larger than Einstein radii of microlenses and therefore microlenses magnify  only a small part of the region emitting in the optical or UV band (see e.g. \\cite{Pop01b,Aba02}, for UV and optical spectral line region). This is reason that it could be a very tiny effect from an observer point of view.  The aim of this paper is to discuss the contribution of microlensing to the relation $\\sigma$ {\\it vs.} $z$ for X-ray radiation considering the recent results given by \\cite{Man02} and \\cite{Popovic03}. In the next section we will consider the optical depth. ",
        "conclusions": "For a discussion of the contribution of microlensing to the X-ray variability of high-redshift QSOs we calculated optical depth considering the density of deflectors in the halo and bulge of the host galaxy as well as for a cosmological distribution of microdeflectors.  From our calculations we can conclude:  i) The optical depth in the bulge and halo of host galaxy is $\\sim 10^{-4}$. This is in good agreement with previous estimates by \\cite{Popovic03}. Microlensing by  deflectors from the host galaxy halo and bulge makes a minor contribution to the  X-ray variability of QSOs.  ii) The optical depth for  cosmologically distributed deflectors could be $\\sim 10^{-2}-0.1$ at $z\\sim 2$ and might contribute  significantly to the X-ray variability of high-redshift QSOs. The value $\\tau\\sim 0.1$ corresponds to the case where compact dark matter forms cosmologically distributed microlenses.  iii) The optical depth for  cosmologically distributed deflectors ($\\tau_L^p$)  is higher for $z>2$ and increases slowly beyond $z=2$. This indicates that the contribution of microlensing on the X-ray variability of QSOs with redshift $z>2$ may be significant, and also  that this contribution could  be nearly constant for high-redshift QSOs. This is in good agreement with the fact that AGNs of the same X-ray luminosity are more variable at $z>2$ \\citep{Man02}.  { iv) Observations of X-ray variations of unlensed QSOs can be used for estimations of matter fraction of microlenses. The rate of microlensing can be used for estimates of the cosmological density of microlenses, and consequently (see Sec. 2.3) the fraction of dark matter microlenses, but the durations of microlensing events could be used for gravitational microlens mass estimations.}"
    },
    "0403/astro-ph0403062_arXiv.txt": {
        "abstract": "We present high-resolution X-ray spectra of the multiple (hierarchical quadruple) weak-lined T Tauri star system HD 98800, obtained with the High Energy Transmission Gratings Spectrograph (HETGS) aboard the Chandra X-ray Observatory (CXO). In the zeroth-order CXO/HETGS X-ray image, both principle binary components of HD 98800 (A and B, separation $0.8''$) are detected; component A was observed to flare during the observation. The infrared excess (dust disk) component, HD 98800B, is a factor $\\sim4$ fainter in X-rays than the apparently ``diskless'' HD 98800A, in quiescence. The line ratios of He-like species (e.g., Ne {\\sc ix}, O {\\sc vii}) in the HD 98800A spectrum indicate that the X-ray-emitting plasma around HD 98800 is in a typical coronal density regime ($\\log{n} \\stackrel{<}{\\sim} 11$). We conclude that the dominant X-ray-emitting component(s) of HD 98800 is (are) coronally active. The sharp spectral differences between HD 98800 and the classical T Tauri star TW Hya demonstrate the potential utility of high-resolution X-ray spectroscopy in providing diagnostics of pre-main sequence accretion processes. ",
        "introduction": "The Einstein and ROSAT missions established the ubiquity of X-ray emission from low-mass, pre-main sequence (PMS) stars. These early X-ray observations, and subsequent observations by the ASCA satellite X-ray observatory, left unsolved the fundamental problem of the physical origin of the X-ray emission, which could be solar-type coronal activity, star-disk interactions, or some combination of these mechanisms. Much of the X-ray data pointed to the likely importance of magnetic activity (Feigelson \\& Montmerle 1999).  To make further progress, it is necessary to determine the temperature distributions, densities, and elemental abundances of the X-ray-emitting plasmas of PMS stars, so as to provide constraints on models of X-ray emission from coronal and star-disk interactions and to compare with, e.g., the physical conditions characterizing well-established stellar coronal X-ray sources (such as RS CVn systems and other close binaries).  With the advent of X-ray gratings spectrometers aboard the Chandra X-ray Observatory and the XMM-Newton satellite observatory, astronomers are now beginning to explore the X-ray spectral characteristics of PMS stars. In this regard, the TW Hya Association (TWA) represents an especially useful young PMS cluster. The TWA is a group of about 30 PMS stars  (Zuckerman et al.\\ 2001 and references therein) located only $\\sim50$ pc from the Sun and far from the nearest star forming clouds (Kastner et al.\\ 1997). Its age ($\\sim5-10$ Myr; e.g., Weintraub et al.\\ 2000, Zuckerman et al.\\ 2001) likely corresponds to the epoch of Jovian planet formation in the early solar system. The TWA's proximity, relatively large ratio of X-ray to bolometric luminosity ($\\log{L_X/L_\\star} \\sim 3$, where $L_X$ is measured within the ROSAT spectral bandpass of 0.1 to 2.4 keV; Kastner et al.\\ 1997), and lack of cloud absorption lead to uniformly high X-ray fluxes among its member stars. In terms of ROSAT X-ray spectral properties, the TWA appears to represent a transition stage between cloud-embedded PMS stars and the zero-age main sequence (Kastner et al.\\ 2003).  The CXO/HETGS spectrum of TW Hya itself is perhaps the most extreme and intriguing of the many Chandra/HETGS spectra of X-ray active stars obtained to date, in several key respects (Kastner et al.\\ 2002): (1) Ne is highly overabundant and Fe severely underabundant, even in comparison with stars exhibiting strong coronal abundance anomalies; (2) the temperature distribution derived from fluxes of temperature-sensitive emission lines is sharply peaked, at $\\log T = 6.5$; (3) perhaps most significantly, density-sensitive line ratios of Ne~{\\sc ix} and O~{\\sc vii} indicate plasma densities $\\log n \\sim 13$. This is more than an order of magnitude larger than density estimates similarly obtained for coronally active late-type stars.  TW Hya remains the TWA's only unambiguous example of an actively accreting (i.e., ``classical'') T Tauri star. Given the evidence that TW~Hya likely is surrounded by a circumstellar disk from which it is still accreting (e.g., Muzzerole et al.\\ 2000), we used the Chandra/HETGS results to explore the hypothesis that X-ray emission from classical T Tauri stars might originate from accretion streams that connect the circumstellar disk to the star (Kastner et al.\\ 2002). Both the density range and the characteristic temperature of X-ray emission obtained from modeling the CXO/HETGS spectrum are consistent with recent models of magnetospheric accretion onto T Tauri stars (Kuker, Henning, \\& Rudiger 2003).  In this {\\it Letter}, we report on the results of CXO/HETGS observations of a second well-studied TWA member, HD~98800 (TV Crt). HD 98800 is a hierarchical quadruple, weak-lined T Tauri Star (wTTS) and is one of the best examples of a solar-mass ``Vega-type'' (infrared excess) system (Zuckerman \\& Becklin 1993; Sylvester et al.\\ 1996; Soderblom et al.\\ 1996, 1998). HD 98800 is a visual double with a separation of $\\sim0.8''$. Each visual component is itself a spectroscopic binary, and one of these (HD 98800B) is double-lined (Soderblom et al.\\ 1996). Although the HD 98800B binary system harbors a dust disk, the system is apparently non-accreting; models indicate that the dust disk has an inner gap extending to $\\sim$2 AU (Prato et al.\\ 2001). The HD 98800 system displays an X-ray luminosity similar to that of its fellow TWA member TW~Hya (Kastner et al.\\ 1997). Given that HD 98800 does not appear to be accreting, it makes an excellent target for further investigation into the origins of X-ray emission from PMS stars with CXO/HETGS. ",
        "conclusions": "It is intriguing that the apparently ``diskless'' component of the HD 98800 system, HD 98800A, appears to be the stronger X-ray source (Fig.\\ 1). Given the relatively small visual extinction toward the system ($E(B-V) = 0.10$; Sylvester et al.\\ 1996) and the inference that the HD 98800B binary is viewed at intermediate inclination angle (Prato et al.\\ 2001), it seems unlikely that this difference in apparent X-ray luminosities is due to differential intervening absorption. Instead, it appears that A is intrinsically brighter in X-rays than B, with a difference of $\\Delta (\\log{L_X/L_{\\rm bol}}) \\approx 0.5$ between components (see \\S\\S 2 \\& 3.1 and Prato et al.\\ 2001). While this result appears consistent with a recent study of Orion showing that disk-enshrouded pre-main sequence (PMS) stars are, in general, weaker X-ray sources than ``diskless'' PMS stars (Flaccomio et al.\\ 2003), one must first demonstrate that the quiescent emission from both components does not show long-term variability, and that other binary TTS resolvable by Chandra show X-ray flux ratios that are likewise anticorrelated with their relative IR excesses.  In terms of the {\\it fir} line ratios of He-like ions, and its Ne to Fe line ratios, the CXO/HETG X-ray spectrum of HD 98800 resembles those of ``classical'' coronal sources, such as II Peg, UX Ari, and HR 1099 (see, e.g., Fig.\\ 6 in Kastner et al.\\ 2002). Hence --- although it remains to determine {\\it fir} line ratios for additional weak-lined T Tauri stars --- {\\it the {\\rm fir} line ratio results for HD 98800 provide some of the strongest evidence to date that the X-ray emission from such stars is coronal in origin.}  There is a marked contrast between the {\\it fir} ratios in the X-ray spectra of HD 98800 and TW Hya, however (Fig.\\ 4). For the latter (classical T Tauri) star, the forbidden line is by far the weakest of the O {\\sc vii} and Ne {\\sc ix} triplets, and the {\\it i:r} ratio is near unity in each case (the Mg {\\sc xi} triplet is anomalously weak in the TW Hya spectrum; Kastner et al.\\ 2002). Fig.\\ 5 shows that, assuming the UV radiation field incident on the X-ray emitting plasma of TW Hya is not strong (see below), the plausible density regimes of Ne {\\sc ix} line formation are non-overlapping, for TW Hya and HD 98800; i.e., the Ne {\\sc ix} {\\it f:i} ratio in the spectrum of TW Hya requires $\\log{n} > 12$.  Furthermore, whereas the differential emission measure (DEM) distribution of TW Hya is sharply peaked at $\\log{T} \\sim 6.5$ (Kastner et al.\\ 2002), the DEM distribution of HD 98800 is relatively flat over the temperature range $\\log{T} \\sim 6.4$ to $\\log{T} \\sim 7.0$. The latter behavior is more like that of coronal sources (Huenemoerder et al.\\ 2003a,b) although --- unlike such sources --- HD 98800 evidently lacks strong emissivity around $\\log{T}=7.2-7.6$. This may be an indication of the evolutionary status of the corona, in that the dynamo is not yet as strong as in the coronally active binaries, or that the flare frequency, which seems to drive the hotter peak, is not as high.  The overall X-ray spectral similarity between HD 98800 and coronally active stars makes the CXO/HETG X-ray spectrum of the classical TTS TW Hya that much more remarkable. As argued by Kastner et al.\\ (2002), both the plasma densities implied by line ratios of He-like ions and its sharply peaked (and rather low) X-ray temperature distribution point to accretion as a likely source of some or all of its X-ray emission. The modeling of Kuker et al.\\ (2003) lends additional credence to this argument. As TW Hya evidently is actively accreting (based on its strong H$\\alpha$ emission and its UV and near-infrared excesses) whereas HD 98800 is not (based on these same accretion indicators), the sharp distinction between their X-ray spectral characteristics appears to further support the hypothesis that the X-ray emission from TW Hya --- and, by extension, other classical TTS --- may be generated, at least in part, via accretion.  On the other hand, these two TWA stars are similar in X-rays, in certain respects. For example, TW Hya was also observed to flare, with flare characteristics (e.g., rise time, decay time, peak flare to quiescent count rates) similar to those of HD 98800 (Kastner et al.\\ 2002). In addition, although TW Hya is by far the most extreme star thus measured by HETG in terms of its Ne/Fe abundance ratio, the X-ray spectrum of HD 98800 shows similar abundance patterns overall (Huenemoerder et al.\\ 2003b).  These similarities would appear to cast some doubt on the accretion hypothesis for TW Hya.  It is possible, for example, that UV radiation, generated in accretion streams onto TW Hya, depletes the populations of atomic levels responsible for the forbidden line component of the He-like triplets, thereby driving the line ratios to their high-density limits (e.g., Ness et al.\\ 2002). Although this phenomenon is reasonably well established in the case of the intense UV fields of X-ray-luminous O stars, it is less clear that it is a viable model in the case of accreting cTTS. The relatively weak UV fields of such stars likely would require that the X-ray-emitting plasma be in very close proximity to the UV source -- effectively placing the point of X-ray generation within (or very near) the accretion stream itself.  Whatever their origin, the contrasting results for the density-sensitive line ratios of He-like ions in the X-ray spectra of TW Hya and HD 98800 (Fig.\\ 5) suggest that fundamentally different physical conditions characterize the X-ray emitting plasmas of cTTS and wTTS. CXO/HETG observations of additional cTTS and wTTS systems, as well as detailed physical models of UV-irradiated coronal plasmas, are now required to establish whether these He-like triplet line ratios are probing X-ray emitting plasma in accretion funnels or are, instead, diagnostic of the intensity of accretion-powered UV and its proximity to the corona of the accreting star."
    },
    "0403/astro-ph0403548_arXiv.txt": {
        "abstract": "We compute the synthetic 2-100 keV spectrum of an actively star forming galaxy. To this aim we use a luminosity function of point sources appropriate for starbursts (based on {\\it Chandra} data), as well as the types of stellar sources and their corresponding spectra. Our estimates indicate that a Compton spectral component - resulting from scattering of SN-accelerated electrons by ambient FIR and CMB photons - could possibly be detected, in deep INTEGRAL observations of nearby starburst galaxies, only if there is a break in the spectra of the brightest ($L \\geq 2 \\times 10^{38}$ erg s$^{-1}$) point sources. ",
        "introduction": "Star formation (SF) leads to X-ray emission on various spatial and temporal scales, including emission from O stars, X-ray binaries, supernovae and their remnants (SNRs), and galactic-scale emission from diffuse hot gas, and Compton scattering of FIR \\& CMB photons by SN-accelerated electrons. Integrated spectra of actively SF galaxies may show all these components as well as emission from an active nucleus (Rephaeli et al. 1991, 1995).  Persic \\& Rephaeli (2002; hereafter, PR02) have quantitatively assessed the roles of the various X-ray emission mechanisms in SFGs. They have used an equilibrium stellar-population synthesis model of the Galactic population of X-ray binaries to deduce birthrates for interacting binaries; these, combined with estimates of the duration of the X-ray bright phase, have allowed PR02 to make realistic estimates of the relative (Galactic) abundances of high-mass and low-mass X-ray binaries (HMXBs and LMXBs, in which the optical companions are main-sequence stars with $M>8\\,M_\\odot$ and $M<M_\\odot$, respectively). The abundance of SNRs (both Type II and Ia) was also consistently estimated. From the literature PR02 derived typical spectra for these source classes. The spectral properties and relative abundances of the various classes of stellar sources allowed PR02 to calculate the composite X-ray spectrum arising from a stellar population of Galactic composition.  In addition to such \"stellar\" components, PR02 considered two diffuse, non-stellar contributions to a starburst (SB) X-ray spectrum.  \\noindent {\\it (i)} Thermal emission from the regions of interaction between SN-powered outgoing galactic wind and denser ambient ISM, mostly at energies $\\leq$1 keV, and thus of little concern to us here.  \\noindent {\\it (ii)} Non-thermal emission due to Compton scattering of the SN-accelerated relativistic radio-emitting electrons off the FIR and CMB radiation fields. The high SN rate in a SB is bound to yield high relativistic electron densities since SN shocks are known to be primary sites of cosmic ray acceleration. Moreover, in a SB environment the mean FIR energy density can be $\\sim 10$ times higher than that of the CMB. Compton scattering of a $100 \\mu$ photon by a $\\sim$1 GeV electron boosts the energy of the photon to $\\sim$10 keV. A high density of electrons in a SB would obviously be reflected also in enhanced radio emission. If the same electron population produces both the radio and hard X-ray emission, then both spectra are expected to have {\\it roughly} the same power-law index (Rephaeli 1979; Goldshmidt \\& Rephaeli (1995). Compton X-ray emission gauges the current SFR since the electron Compton-synchrotron lifetime is approximately $\\leq$10$^8$ yr (i.e. comparable to or shorter than a typical starburst duration).  PR02 determined that the $\\sim$2-15 keV emission is dominated by LMXBs; this is not surprising, given that the synthetic spectrum was calibrated by that of a normal, quietly star forming galaxy like ours. Due to the long times between the formation and the onset of the X-ray phase (the optical companion is a subsolar main-sequence star) in LMXBs, their emission traces the higher, {\\it past} Galactic SFR, and not the lower, {\\it current} SFR. Thus in order to obtain an estimate of the current SFR, the spectral contribution of LMXBs should not be included. As noted by PR02, at energies of approximately $\\geq$30 keV (relevant for INTEGRAL) the SB spectrum is expected to be dominated by the Compton component.  Recent observations of nearby galaxies -- mostly with the {\\it Chandra} satellite -- have revealed several X-ray point sources (XPs) with 2-10 keV luminosities in the range $2 \\times 10^{38}$ erg s$^{-1} \\leq L \\leq 10^{39}$ erg s$^{-1}$ (hereafter dubbed Very Luminous Sources: VLXs), and a few Ultra-Luminous Sources (ULXs), with $L \\geq 10^{39}$ erg s$^{-1}$ (Fabbiano \\& White 2003) \\footnote{These ranges correspond to Eddington luminosities for spherical accretion on, respectively, 2 and $8\\,M_\\odot$ BHs. This is the range of BH masses that are thought to be attainable via ordinary stellar evolution. ULXs are called 'super-Eddington sources'.}. Although their nature is not yet known, VLXs and ULXs are observed to be directly linked to current star formation activity (e.g., Zezas et al. 2002), and when identified, the optical counterpart of ULXs are O stars (Liu et al. 2002), thus establishing a link between HMXBs and ULXs. The spectra of VLXs and ULXs are characterized both as power-law in the range of interest, with index $\\Gamma \\sim 2$ and $\\sim$1.3, respectively. \\footnote{Broadly speaking, the VLXs and ULXs could, in fact, describe different luminosity/spectral states of BH X-ray binaries (BHXBs) as a function of the mass infall rate ($\\dot M$) onto the BH, while the other relevant parameters are constrained within a much smaller range (e.g., BH mass $M \\sim 5-10 M_\\odot$, efficiency $\\eta \\sim 0.3 - 0.6$). In fact, qualitatively speaking, the inverse correlation between luminosity and spectral index (i.e., flatter spectra for higher luminosities) can be understood within the accretion model devised to explain the high X-ray luminosities of accretion-powered binary systems. The increase in luminosity is driven by an increase of $\\dot M$, which in turn means a piling up of material around the emission region that leads to a higher Compton optical depth, and hence to a lower value of $\\Gamma$ i.e. a flatter spectrum. Namely, if a source is embedded in a plasma cloud of optical depth $\\tau$ and temperature $T_{\\rm e}$, the spectral index $\\Gamma$ of the high-frequency wing of the spectrum of the escaping radiation is $$ \\Gamma ~=~ \\biggl[ {9 \\over 4} ~+~ {\\pi^2 m_{\\rm e} c^2 \\over 3 \\, (\\tau + 2/3)^2 \\, kT_{\\rm e}} \\biggr]^{1/2} ~-~ {1 \\over 2} $$ (Sunyaev \\& Titarchuk 1980; Shapiro et al. 1976).} ",
        "conclusions": "The resulting 2-100 keV spectrum is shown in Fig.1. In the synthetic spectrum a conspicuous knee, due to the break in the HMXB spectrum, is immediately recognizable at 20 keV. If the ULX component, which turns out to lie very close to that of HMXBs for energies up to 20 keV, did turn out to have a similar break, then the knee in the synthetic spectrum would be sharper.  \\begin{figure} \\vspace{5.0cm} \\special{psfile=SB_phot.ps hoffset=-6 voffset=160  hscale=33.0  vscale=27.0 angle=-90}  \\caption{The integrated 2-100 keV spectrum of a starburst. Shown in ascending order at 5 keV are the SNR, Compton, ULX, HMXB, and VLX components. The \"stellar\" components are normalized in energy flux in the 2-100 keV band. The Compton component has been normalized at 0.1 of the total flux at 10 keV. The spectrum is photoelectrically absorbed through an HI column density $N_{\\rm H}= 10^{22}$ cm$^{-2}$.} \\end{figure}  An important related consideration is that the Compton component will be swamped by the VLX and ULX components and thus practically undetectable by INTEGRAL, unless both \"stellar\" components exhibit a spectral break similar to that of HMXBs. It can be argued that the HMXB spectral break is related to the magnetic field strength, $B$, on the surface of the accreting neutron star: i) statistically, the cutoff energy, $\\epsilon_c$, is correlated with the observed cyclotron resonance energy, $\\epsilon_0$, by $\\epsilon_0 \\simeq 2\\, \\epsilon_c$: as the latter depends on $B$ so will the former, which will then provide a measure of $B$ (a most helpful feature in those cases where a resonant scattering feature is not observed -- White et al. 1995); and ii) physically, the spectral break could be due to radiative transfer effects in a strong magnetic field (Boldt et al. 1976). If the magnetic fields able to induce $\\sim$20 keV spectral breaks in VLX-ULX spectra have strengths comparable to those observed in HMXBs \\footnote{A crude estimation suggests that for BH masses in the range $2-8 \\,M_\\odot$, which can be produced via ordinary evolution (i.e. through type II SN events) by progenitor stars with $M \\sim 8-40\\, M_\\odot$, the conservation of magnetic flux implies that for $B \\sim 10-100$ G at the surface of the progenitor star, it will be $B \\sim 10^{12} - 10^{13}$ G at the Schwarzschild radius, nearly irrespective of the BH mass. (See Maeder \\& Meynet 1989 for the relevant quantities of the stellar progenitors). These values are consistent with those deduced for the magnetized neutron stars which are members of HMXB systems (see Boldt et al. 1976; White et al. 1995). }, then a $\\sim$20 keV cutoff is expected also in the spectra of VLX and ULX sources. Hence, the knee of the synthetic spectrum will turn out more marked and the Compton component will possibly stand out more clearly (at energies of approximately $\\geq$30 keV). INTEGRAL measurements can possibly determine whether there is a break in VLX/ULX spectra.  On the other hand the Compton emission, in order to be discernible even in the absence of VLX/ULX spectral breaks, has to be stronger than assumed here. In our calculation we have assumed the Compton flux to be 0.1 of the total at 10 keV (which implies that it is $\\sim$10$\\%$ of the total 2-10 keV flux). This is a factor of $\\sim$2 higher than the upper limit set by Griffith et al. (2000) on the Compton flux fraction of M82 observed with {\\it Chandra} (see Persic \\& Rephaeli 2003). It is possible that in SB galaxies high SN rates and dust temperatures enhance the Compton component such that it dominates the spectrum at energies higher than $\\sim$30 keV."
    },
    "0403/astro-ph0403312_arXiv.txt": {
        "abstract": "Precision radial velocity measurements of the $\\gamma$ Cephei (HR8974) binary system suggest the existence of a planetary companion with a minimum mass of 1.7 Jupiter-mass on an elliptical orbit with a $\\sim$2.14 AU semimajor axis and 0.12 eccentricity \\cite{hatzes03}. I present in this paper a summary of the results of an extensive numerical study of the orbital stability of this three-body system for different values of the semimajor axis and orbital eccentricity of the binary, and also the orbital inclination of the planet. Numerical integrations indicate that the system is stable for the planet's orbital inclination ranging from 0 to 60 degrees, and for the binary's orbital eccentricity less than 0.5. The results also indicate that for large values of the inclination, the system may be locked in a Kozai resonance. ",
        "introduction": "The existence of planets in binary star systems is no longer a mere idea. Approximately 35\\% of extrasolar planets discovered till 2002 exist in multiple star systems \\cite{kortenkamp01}. These systems are mostly wide binaries with separations between 500 and 750 AU, and with planets revolving around one of the stars. At such large distances, the perturbative effect of one star on the formation and dynamical evolution of planets around the other star is entirely negligible. A recently discovered Jupiter-like planet around the primary of the $\\gamma$ Cephei binary system is, however, an exception to this rule.  Gamma Cephei is a close spectroscopic binary composed of a 1.59 solar-mass K1 IV subgiant and a 0.4 solar-mass red M dwarf \\cite{fur03}. The semimajor axis of this system has been reported to have a lower value of 18.5 AU \\cite{hatzes03} and an upper value of 21 AU \\cite{griffin02}. Precision radial velocity measurements suggest that a planet with a minimum mass of 1.7 Jupiter-mass revolves around the primary of this system on an elliptical orbit with an eccentricity of 0.12 $\\pm$ 0.05 and a semimajor axis of approximately 2.14 AU. Being the first discovered S-type binary-planetary system\\footnote{As classified by Rabl $\\&$ Dvorak (1988), a binary-planetary system is called S-type when the planet revolves around one of the stars, and P-type when the planet revolves around the entire binary.}, it is quite valuable to investigate whether this system is dynamically stable, and for what values of its orbital parameters its stability will remain. In this paper I present a summary of the results of an extensive numerical study of the dynamical stability of this system for different values of the orbital parameters of the binary and also the orbital inclination of the planet. A more comprehensive study of the dynamics of this system are to be published elsewhere. ",
        "conclusions": ""
    },
    "0403/astro-ph0403638_arXiv.txt": {
        "abstract": "We present measurements of the redshift--space three-point correlation function of galaxies in the Sloan Digital Sky Survey (SDSS).  For the first time, we analyze the dependence of this statistic on galaxy morphology, color, and luminosity.  In order to control systematics due to selection effects, we used $r$--band, volume-limited samples of galaxies, constructed from magnitude-limited SDSS data ($14.5<r<17.5$), and further divided the samples into two morphological types (early and late) or two color populations (red and blue).  The three-point correlation function of SDSS galaxies follows the hierarchical relation well, and the reduced three-point amplitudes in redshift--space are almost scale-independent ($Q_z=0.5\\sim1.0$).  In addition, their dependence on the morphology, color, and luminosity is not statistically significant.  Given the robust morphological, color and luminosity dependences of the two-point correlation function, this implies that galaxy biasing is complex on weakly non-linear to non-linear scales. We show that a simple deterministic linear relation with the underlying mass could not explain our measurements on these scales. ",
        "introduction": "While the first--year WMAP (Wilkinson Microwave Anisotropy Probe) data imply that the primordial density fluctuations at $z\\approx1000$ are virtually Gaussian \\citep{K2003}, the present structure traced by galaxies shows significant non-Gaussianity.  Such non-Gaussian features naturally arise during the course of nonlinear gravitational evolution of dark-matter density fields, the formation of luminous galaxies and their subsequent evolution. Therefore, the degree and nature of non-Gaussian signatures in the galaxy distribution will provide important empirical constraints on the physics of galaxy formation. The three-point correlation function (3PCF) is the lowest-order unambiguous statistic to characterize such non-Gaussianities.  The determination of the 3PCF of galaxies was pioneered by \\citet{PG1975} using the Lick and Zwicky angular catalogs of galaxies. \\citet{GP1977} found that the 3PCF $\\zeta (r_{12}, r_{23}, r_{31})$ obeys the following {\\it hierarchical relation}: \\begin{eqnarray} &&  \\zeta (r_{12}, r_{23}, r_{31}) \\cr && = Q_{\\rm r} \\, [\\xi(r_{12})\\xi(r_{23}) + \\xi(r_{23})\\xi(r_{31}) + \\xi(r_{31})\\xi(r_{12})], \\end{eqnarray} with $Q_{\\rm r}$ being a constant, and $\\xi(r)$ is the two-point correlation function (2PCF). The value of $Q_{\\rm r}$ in real space de-projected from these angular catalogues is $1.29 \\pm 0.21$ for $r \\lesssim 3 h^{-1}{\\rm Mpc}$.  Although subsequent analyses of redshift catalogs confirmed the hierarchical relation, at least approximately, the value of $Q_z$ (in redshift space) appears to be smaller, $0.5\\sim1$ (\\cite{Bean1983}; \\cite{Hale-Sutton}; Jing, B\\\"orner 1998). Recently, \\citet{JB2004} have studied the luminosity dependence of $Q$, and found a small, but significant, trend that brighter galaxies tend to have a lower amplitude than fainter ones.  N-body simulations and perturbative analysis generally predict that the 3PCF should depart from the hierarchical relation, especially on non-linear scales (Fry 1984; Suto, Matsubara 1994; Jing, B\\\"{o}rner 1997; Barriga, Gazta\\~{n}aga 2002). However, Matsubara \\& Suto (1994) demonstrated that redshift distortions substantially reduce the scale-dependence of $Q_{\\rm r}$, resulting in $Q_z$ ($Q$ in redshift space) being nearly constant.  Still, the amplitude of the $Q_z$ of galaxies is roughly (50\\% $\\sim$ 100\\%) smaller than predicted by N-body simulations.  The discrepancy is likely to be the effect of biasing.  Recently, \\citet{TJ2003} proposed a phenomenological model to predict the 3PCF as a function of the galaxy properties, which is based on the halo model.  They argue that the color dependence of $Q_{\\rm r}$ should be strong, i.e., the galaxy biasing severely affects the 3PCFs.  There are several statistics that are closely related to the 3PCF, e.g., the reduced moments of counts in cells, such as skewness and kurtosis (see \\cite{SzapudiEtal2002}, 2004 in preparation). Another similar statistics is the bispectrum, which is the Fourier counterpart of the 3PCF. Recently, \\citet{VerdeEtal2002} computed the bispectrum of the 2dF Galaxy Redshift Survey, and concluded that the non-linear bias is consistent with zero. The other complementary approach to quantify the non-Gaussianity in the SDSS galaxies by means of topological analysis (i.e., in terms of the Minkowski functionals) was already conducted (Hoyle et al. 2002; Hikage et al. 2002, 2003). The above results are largely consistent with that predicted from purely gravitational nonlinearity, while a weak morphological difference is marginally detected (e.g., figure 13 of \\cite{H2003}).  Here, we report on the first results of characterizing the non-Gaussianity in the SDSS galaxy samples using 3PCFs in redshift space. In particular, we consider the dependence of 3PCFs on the morphology, color, and luminosity of galaxies by constructing volume-limited samples so as to avoid possible systematics due to the selection function.  A separate companion paper by Nichol et al. (in preparation) focuses on a detailed comparison of 3PCFs between the 2dF (Jing, B\\\"orner 2004) and the SDSS galaxies.   The paper is organized as follows: section 2 describes our volume-limited sample of SDSS galaxies and the morphology and color classification methods. The measurements of 3PCFs are detailed in section 3 with particular attention given to their dependence on the morphology, color and luminosity of galaxies. Finally, we summarize the results in section 4.  We compare several different estimators of 2PCFs and 3PCFs in Appendix. Throughout the data analysis, we adopt the following set of cosmological parameters: the matter density parameter, $\\Omega_{\\rm m}=0.3$; the dimensionless cosmological constant, $\\Omega_\\Lambda=0.7$; and the Hubble constant in units of 100km/s/Mpc, $h=0.7$. ",
        "conclusions": "We have presented the first detailed study of the three-point correlation function of SDSS galaxies in redshift space.  We examined the dependence of their reduced amplitudes, $Q_z$, [eq.(\\ref{eq:Qdef})] on the scale and shape of the triangles for various volumed-limited samples with different morphologies, colors, and luminosities.  As for equilateral triangles, we basically confirm the hierarchical clustering relation for scales of $1h^{-1}{\\rm Mpc}<s<10h^{-1}{\\rm Mpc}$. On the other hand, $Q_z$ shows significant shape dependence, as expected from perturbation theory and N-body simulations, particularly at large scale, although the amplitude of $Q_z$ is much smaller than the expectation.  The most important finding is that although the two-point correlation functions of those galaxies show a clear dependence on their luminosity, morphology, and color, we did not find any robust dependence of $Q_z$ on those properties; the $Q_z$ results when taken alone imply that galaxies are faithful tracers of the underlying mass distribution, and that their intrinsic properties are independent of the purely nonlinear gravitational evolution.  At the same time, however, their two-point correlation functions strongly indicate that they are biased tracers.  If we attempt to reconcile these two behaviors by using a simple linear bias model, we cannot adjust it to simultaneously explain both the two-point correlation functions and the insensitivity of $Q_z$ on the galaxy properties.  This implies that the galaxy biasing is non-linear and fairly complex, while nonetheless displaying a remarkable constancy in reduced amplitudes.  The above conclusion may appear to be inconsistent with the argument of \\citet{VerdeEtal2002}, who concluded that the non-linearity in 2dF galaxy biasing is negligible from the bispectrum analysis. We note, however, that they did not attempt to divide those galaxies into different subsamples. If we look at the results of all galaxies alone, top-panels of figures \\ref{fig:Q_eq_cin_lum} and \\ref{fig:2ptbias_lum} may indicate the same conclusion as reached by \\citet{VerdeEtal2002}. In other words, our empirical finding indicates that the non-linearity of the biasing, $b_2$, of galaxies vanishes when the corresponding linear biasing parameter, $b_1$, is unity.  The fact that the two parameters, $b_1$ and $b_2$, are highly correlated with each other is qualitatively consistent with some biasing model predictions including halo models or peak models (Mo et al. 1997).  It is remarkable that the current galaxy survey approaches the level of discussing the validity of those models quantitatively.  Future SDSS datasets with more area and a larger number of galaxies will be able to place stronger limits on this observation. This in turn will be useful for a detailed understanding of the nature of galaxy biasing.  \\bigskip We thank Y.P. Jing and Gerhard B\\\"orner for useful correspondences on their measurements of three-point correlation functions of 2dF galaxies and an anonymous referee for his/her useful comments. I. K. gratefully acknowledges support from the Takenaka-Ikueikai fellowship. He enjoyed the hospitality at Carnegie Mellon University and University of Pittsburgh where part of this work was done, and thanks the Hayakawa fund of Astronomical Society of Japan for the travel support.  Numerical computations were carried out at ADAC (Astronomical Data Analysis Center) of the National Astronomical Observatory of Japan, and also at computer facilities at the University of Tokyo supported by the Special Coordination Fund for Promoting Science and Technology, Ministry of Education, Culture, Sports, Science and Technology. IS and JP was supported by NASA through AISR NAG5-11996, and ATP NASA NAG5-12101 as well as by NSF grants AST02-06243 and ITR 1120201-128440.  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, 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, Princeton University, the United States Naval Observatory, and the University of Washington.  \\bigskip  \\appendix"
    },
    "0403/hep-ph0403296_arXiv.txt": {
        "abstract": "We perform a new study of future sensitivities of galaxy redshift surveys to the free-streaming effect caused by neutrino masses, adding the information on cosmological parameters from measurements of primary anisotropies of the cosmic microwave background (CMB). Our reference cosmological scenario has nine parameters and three different neutrino masses, with a hierarchy imposed by oscillation experiments.  Within the present decade, the combination of the Sloan Digital Sky Survey (SDSS) and CMB data from the PLANCK experiment will have a 2$\\sigma$ detection threshold on the total neutrino mass close to $0.2$ eV.  This estimate is robust against the inclusion of extra free parameters in the reference cosmological model. On a longer term, the next generation of experiments may reach values of order $\\sum m_{\\nu} = 0.1$ eV at 2$\\sigma$, or better if a galaxy redshift survey significantly larger than SDSS is completed.  We also discuss how the small changes on the free-streaming scales in the normal and inverted hierarchy schemes are translated into the expected errors from future cosmological data. ",
        "introduction": "Neutrino physics has provided the first clear indication of particle physics beyond the Standard Model, since we have experimental evidences for non-zero neutrino masses. Analyses of data from atmospheric and solar neutrino experiments have shown the allowed regions for the squared mass differences ($\\Delta m_\\nu^2$) at two different scales. Such values will be known with better precision in the next years, in particular for the larger atmospheric $\\Delta m_\\nu^2$ using the results of future long-baseline oscillation experiments.  However, from oscillation experiments no information can be obtained on the absolute values of neutrino masses, since the lightest neutrino mass remains unconstrained. Tritium decay experiments tell us that each neutrino mass cannot be larger than 2.2 eV (95\\% CL) at present \\cite{Bonn:tw}, to be improved to $\\sim$0.35 eV with KATRIN \\cite{Osipowicz:2001sq}. More stringent bounds exist from experiments searching for neutrinoless double beta decay, that will be improved in the near future \\cite{Elliott:2002xe}, but unfortunately they depend on the details of the neutrino mixing matrix.  Cosmology offers several advantages: the cosmic neutrino background provides an abundant density of relic neutrinos with an equal momentum distribution for all flavors (up to 1\\% corrections), which implies that mixing angles have no effect. Although neutrinos cannot be the dominant dark matter component, they can still constitute a small, hot part of the matter density producing an erasure of perturbations at small scales through their free-streaming effect (for a review, see e.g.\\ \\cite{Dolgov:2002wy}). A comparison with data from the large scale structure (LSS) of the Universe is thus sensitive to neutrino masses, as emphasized in \\cite{Hu:1997mj}.  At present, cosmological data allow us to bound the total neutrino mass to values of $\\sum m_\\nu \\alt 0.6-1.0$ eV \\cite{Spergel:2003cb,Hannestad:2003xv,Elgaroy:2003yh,Tegmark:2003ud, Barger:2003vs,Hannestad:2003ye,Crotty:2004gm}, depending on the data and priors used. These ranges already compromise the 4 neutrino scenarios that could explain the additional large neutrino mass difference required by the LSND results (that also imply a fourth, sterile neutrino), but is not yet capable of reaching the necessary $0.1$ eV range in order to test the hierarchical 3 neutrino schemes. But such small masses could be detected in the next future when more precise cosmological data are available, in a parallel effort to those of beta and double beta decay experiments on Earth.  In this paper we analyze the future sensitivities of cosmological data to neutrino masses, extending the pioneering work \\cite{Hu:1997mj} and in particular the detailed analysis in \\cite{Eisenstein:1998hr} (see also \\cite{Lesgourgues:1999ej}), that was more recently updated in \\cite{Hannestad:2002cn}. In contrast to this last work we consider, in addition to ideal Cosmic Microwave Background (CMB) observations limited only by cosmic variance, the experimental specifications of satellite missions such as PLANCK and the mission concept CMBpol (Inflation Probe), as well as ground-based detectors such as ACT or SPTpol, that will extend the PLANCK data to smaller angular scales. We also increase the number of cosmological parameters of previous analyses, including also the helium fraction, extra relativistic degrees of freedom, spatial curvature, dark energy with constant equation of state, or a primordial spectrum with running tilt. Finally, our work is the first one in which it is assumed that neutrinos have three different masses, in order to compute accurately the free-streaming effect associated to the mass schemes allowed by oscillation experiments.  Note that throughout this work, we will assume that the LSS power spectrum is measured solely with galaxy redshift surveys. For complementary constraints based on gravitational lensing, we refer the reader to Refs.\\ \\cite{Abazajian:2002ck,Kaplinghat:2003bh}.  This paper is organized as follows. In Sec.\\ II we review the expected values of neutrino masses and their impact on Cosmology. We describe future CMB experiments and galaxy surveys in Sec.\\ III and the method to forecast the errors on cosmological parameters in Sec.\\ IV. Finally, we present our results in Sec.\\ V, with a summary and conclusions in Sec.\\ VI. ",
        "conclusions": "In this paper we have analyzed the sensitivities of future CMB and LSS data to the absolute scale of neutrino masses, taking into account realistic experimental sensitivities and extending the results of previous works \\cite{Hu:1997mj,Eisenstein:1998hr,Lesgourgues:1999ej,Hannestad:2002cn}.  We have considered the values of neutrino masses distributed according to the presently favored three neutrino mass schemes, that follow either a normal or an inverted hierarchy. As discussed in section \\ref{sec:numasses}, a different distribution of the same total neutrino mass leads to small changes in the cosmological evolution of neutrinos, and in particular in the free-streaming scales (qualitatively discussed, for instance, in \\cite{Hu:1997mj,Hannestad:2002cn}). These changes disappear when the total neutrino mass enters the quasi-degenerate region.  We used the Fisher matrix method to forecast the errors on cosmological parameters that can be extracted from future CMB experiment and redshift survey data, assuming a fiducial 9-dimensional cosmological model close to the currently favored $\\Lambda$CDM model. Our theoretical CMB and matter power spectra were generated with the standard Boltzmann code CMBFAST, modified in order to include three neutrino states with different masses.  In particular, for the case of PLANCK and SDSS we found good agreement with the results of \\cite{Eisenstein:1998hr}, with a 2$\\sigma$-error on the total neutrino mass of $0.2$ eV that will allow us to probe only the quasi-degenerate neutrino mass region. Better sensitivity will be achieved with the combination CMBpol and SDSS, for which we found $0.12$ eV, close to the minimum value of the total neutrino mass in the inverted hierarchy case. These results correspond to a conservative value of $k_{\\rm max} = 0.15\\, h$ Mpc$^{-1}$, the maximal wavenumber on which the LSS data are compared with the predictions of linear theory. We also tested how the errors change when including additional cosmological parameters to our fiducial model. In general, we found that the errors on the neutrino masses are not modified in a significant way.  Our results show that the approach where CMB experiments are only limited by cosmic variance (as in \\cite{Hannestad:2002cn}) is probably too simplistic. However, if a future CMB experiment is capable of getting close to such an ideal limit, then the combination with data from galaxy redshift surveys larger than SDSS would lead to errors on the total neutrino mass comparable to the minimum values of the hierarchical scenarios. In such a case, we have shown that there exist slight differences in the expected errors between the two hierarchical neutrino schemes for the same total neutrino mass.  In conclusion, we consider that cosmological data can provide valuable information on the absolute scale of neutrino masses, that nicely complements the present and future projects of beta decay and neutrinoless double beta decay experiments.  This conclusion is reinforced when one takes into account other cosmological probes of neutrino masses, complementary to the approach of the present paper. We can cite, for instance, studies of the distribution of matter in the Universe through the distortions of CMB maps caused by gravitational lensing (measured from non-gaussianities in the CMB maps) \\cite{Kaplinghat:2003bh} and the weak gravitational lensing of background galaxies by intervening matter \\cite{Hu:1999ek,Hu:2002rm,Abazajian:2002ck,Song:2003gg}.  It is interesting to note that any information on the absolute neutrino masses from cosmology will be interesting not only for theoretical neutrino models, but also for connected baryogenesis scenarios which occur through a leptogenesis process (see e.g.\\ \\cite{Giudice:2003jh,Hambye:2003rt,Buchmuller:2004nz})."
    },
    "0403/astro-ph0403130_arXiv.txt": {
        "abstract": "It is shown here that for redshifts $z < 0.5$ the luminosity distance, which is predicted in author's model astro-ph/0005084 v2, fits well supernova observational data of astro-ph/0402512 by A.Riess et al. Discrepancies for higher $z$ would be explained in the model as a result of specific deformation of SN spectra due to a discrete character of photon energy losses. The model does not require any dark energy; it is based on the conjecture that gravitons are super-strong interacting particles fulfilling a flat non-expanding universe. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403306_arXiv.txt": {
        "abstract": "{We resume the discussion about irradiation-driven mass transfer cycles in semi-detached compact binary systems. The analytical model that describes the onset of these cycles, which occur on a thermal timescale of the donor star, is reexamined. We take into account a contribution of the thermal relaxation which is not related to the irradiation of the donor star and which was neglected in previous studies. Cataclysmic variables (CVs) containing extended giant donors are more stable than previously thought. CVs close to the upper edge of the period gap can undergo cycles for low angular momentum loss rates, as they have been suggested by recent magnetic braking prescriptions, while they are stable for high braking rates.  A model for the irradiation geometry that takes into account surface elements near the terminator of the donor star indicates that possibly also low-mass X-ray binaries (LMXBs) can undergo mass transfer cycles. Regarding the braking rate, which is necessary to drive cycles, basically the same restrictions apply for short period LMXBs as they do for short period CVs. We confirm that LMXBs containing giants can undergo cyles. In terms of an irradiation efficiency parameter \\( \\alpha  \\) CVs are susceptible to the irradiation instability for \\( \\alpha \\gtrsim 0.1 \\) while LMXBs are susceptible for \\( \\alpha \\lesssim 0.1 \\).  The predictions of the analytical model are checked by the first long-term evolutionary computations of systems undergoing mass transfer cycles with full 1D stellar models. For unevolved main sequence (MS) and giant donors the analytic model provides reasonable values for the boundaries of the stable and unstable regions while CVs containing highly evolved MS donors are more stable at high braking rates than expected.  Taking into account irradiation the minimum period of CVs is increased by up to 1-2 minutes, depending on \\( \\alpha  \\). ",
        "introduction": "Observations of close binary systems have shown that a star which is illuminated by its companion can exhibit a hotter, illuminated, and a cooler, unilluminated side  \\citep[see, e.g.,][~and references therein]{Eddington26, Ritter00}. This is the well-known reflection effect  \\citep[for a review see][]{Vaz85}. If the illuminated star has a convective envelope, then irradiation inhibits the energy transport through the illuminated outer layers \\cite{Nordlund85}. In the case of a semi-detached compact binary system the reflection effect can be caused by accretion luminosity that is released nearby the compact star. In this case the energy transport, i.e., the intrinsic flux through the illuminated outer layers of the donor star is coupled to the mass transfer rate. This leads to a feedback, the consequences of which are the subject of this paper.  \\corr{The influence of irradiation onto mass transfer was first discussed by  \\citet{Podsiadlowski91} for the simplified case of symmetrical irradiation of the donor. Later, \\citet{Hameury93} tried to simulate asymmetrical irradiation by time-periodic symmetric irradiation and concluded that a more realistic model is required.}  \\corr{In a series of papers  \\citet{Ritter95, Ritter96b, King95b, King96} developed a feedback model for asymmetrical irradiation} that predicts under which conditions the secular mass transfer rate becomes \\corr{unstable and} the system undergoes mass transfer cycles which occur on a thermal timescale of the convective envelope of the donor star.  We reexamine the above-mentioned irradiation feedback model and take into account a contribution of the thermal relaxation which is not related to the irradiation of the donor star. It turns out that this term, which was neglected in previous studies \\corr{of the above-mentioned analytical model}, can become important, especially for giant donors.  Unfortunately, there have been basically no numerical computations to test the validity of this model and especially the boundaries of the stable and unstable regions. Only few computations have been carried out by using bipolytrope stellar models and in most cases also by using simplified models for the treatment of irradiation \\cite{Ritter95, Frank98, Ritter00}. There has been only one short-term evolution using full 1D stellar models and a realistic treatment of irradiation which was published by  \\citet{Hameury97}. It has served more as a {}``proof of concept{}'' for their model of the reduction of the intrinsic flux by illumination than as a numerical confirmation of the irradiation feedback model.  The main purpose of this paper is to confirm the predictions of the irradiation feedback model by numerical long-term evolutionary computations that are based on full 1D stellar models for the donor star, a realistic model for the irradiation geometry, and tabulated results of \\citet{Hameury97} for the reduction of the intrinsic flux by illumination.  In Sect.~\\ref{Chap:InputPhysics} we mention the input physics that is necessary for the feedback model and for the numerical computations. This includes the mass loss rate prescription, the reflection effect, the irradiation geometry, and the thermal relaxation of the donor star. Afterwards, in Sect.~\\ref{Chap:AnalyticalModel} we develop two ordinary differential equations which describe the system adequately for our purpose, and in Sect.~\\ref{Chap:LinearStability} we derive the conditions which describe the onset of instability by means of linear stability analysis.  We give a short overview in Sect.~\\ref{Chap:Code} about input physics, which we use in our code, and subsequently, we present our numerical results for CVs and LMXBs containing evolved and unevolved donor stars in Sect.~\\ref{Chap:Numerics}. ",
        "conclusions": "\\label{Chap:Summary}  We have described the physics of irradiation driven mass transfer cycles by a two-dimensional system of ordinary differential equations similar to  \\citet{King96}. An additional term of order of several times \\( \\frac{H_{\\text {P}}}{R_{2}} \\) which reflects the usual thermal relaxation of the mass losing star, can not be neglected, because it provides an upper limit for the driving timescale \\( \\tau _{\\text {d}} \\), respectively a lower limit for the timescale ratio \\( \\frac{\\tau _{\\text {ce}}}{\\tau '_{\\text {d}}} \\), where mass transfer cycles can occur. Taking into account this term and also the rather large pressure scale height and the deep superadiabatic convection zone of giants we have concluded that CVs with giants above a certain core mass can not undergo mass transfer cycles unless the efficiency parameter \\( \\alpha  \\) is unexpectedly large. On the other hand LMXBs with giants might be susceptible to the irradiation instability as it has been predicted by  \\citet{King97}.  We agree with results of  \\citet{King96, Ritter00} regarding CVs with unevolved MS stars: For high braking rates \\citep[e.g.,][]{Verbunt-Zwaan}, as they are also required by the period gap model  \\citep[e.g.,][]{Spruit83, Kolb93}, such systems are stable except for the most massive \\( \\left( \\sim 1\\, M_{\\sun }\\right)  \\) donor stars. On the other hand, for low braking rates, as they are proposed by  \\citet{Sills00, Andronov03}, CVs above the period gap might undergo cycles. In this case only novalikes are affected since dwarf novae have rather large amplitude outbursts with a rather small duty cycle and therefore a small \\( \\alpha _{\\text {d}} \\). Even for gravitational braking the susceptability to irradiation ceases within the period gap so that CVs with unevolved MS stars below the period gap are stable.  Numerical evolutionary computations have shown that the uncritical application of the predictions of the analytical model to highly evolved remnants of thermal timescale mass transfer does not yield good quantitative estimates for the boundaries of the unstable region. The timescale on which the stellar radius changes, if irradiation is changed, seems to be significantly underestimated by the bipolytrope model so that such systems are more stable at high braking rates than previously expected. The most evolved systems are susceptible to irradiation for \\( 0.1\\lesssim \\alpha \\lesssim 1 \\) and orbital periods roughly between \\( 2 \\) and \\( 5 \\) hours, or donor masses between \\( 0.1 \\) and \\( 0.3\\, M_{\\sun } \\), respectively. For giants the ratio of thermal energy of the envelope divided by the stellar luminosity seems to give a good estimate for \\( \\tau _{\\text {ce}} \\).  Since in the point source model surface elements near the terminator are irradiated at a high angle of incidence, they are only partially blocked by irradiation even for high fluxes. Thus, other than expected by  \\citet{King96, King97, Ritter00} also LMXBs with unevolved or evolved MS stars might be susceptible to the irradiation instability. Regarding the driving timescale for such systems basically the same restrictions apply as for CVs. LMXBs can undergo mass transfer cycles only for \\( \\alpha \\lesssim 0.1 \\). \\corr{The contribution of non-local effects of irradiation, e.g., circulations, which were not taken into account by our model, might be non-negligible for LMXBs. It seems plausible that this effect tendentially destabilizes the mass transfer.} However, it is not clear \\corr{in all cases} whether mass transfer in LMXBs is driven by Roche lobe overflow or by irradiation-induced winds \\cite{Basko73, Basko77, Iben97}. In the latter case our model could obviously not be applied to these systems.  We conclude that the analytical model gives a qualitatively and also quantitatively suitable description for the onset of mass transfer cycles except for highly evolved remnants of thermal timescale mass transfer. We think that the boundaries of the unstable regions can be determined with an accuracy of a factor of \\( \\lesssim 2 \\).  To answer the final question whether the observable CV or LMXB population or parts of it is undergoing mass transfer cycles we need to know two basic parameters: the driving timescale, respectively the magnetic braking law, and the efficiency \\( \\alpha  \\). Nevertheless, we have shown that irradiation-driven mass transfer cycles in compact binaries are possible for not unreasonable values of \\( \\alpha  \\) and \\( \\tau _{\\text {d}} \\)."
    },
    "0403/astro-ph0403076_arXiv.txt": {
        "abstract": "We consider magnetohydrodynamical  interaction between relativistic pulsar wind and static magnetosphere in  binary pulsar system PSR J0737-3039. We  construct  semi-analytical model  describing the  form of the interface separating the two pulsars. An assumption of  vacuum dipole  spin down for Pulsar B leads to eclipse duration ten times longer than observed. We discuss possible Pulsar B  torque modification and magnetic field estimates due to the interaction with Pulsar A wind. Unless the orbital inclination is $\\leq  86 ^\\circ$, the duration of eclipses is typically shorter than the one implied by the size of the eclipsing region. We  propose  that  eclipses occur due to  synchrotron absorption by mildly relativistic particles in the shocked Pulsar A wind.  The corresponding optical depth may be high enough if  Pulsar A wind density is at the upper allowed limits. We derive jump conditions at oblique,  relativistic, magnetohydrodynamical shocks and discuss the  structure of the shocked  Pulsar A wind. Finally, we speculate on a  possible mechanism of orbital modulation of Pulsar B radio emission. ",
        "introduction": "A recent discovery of  eclipsing binary pulsar system PSR J0737-3039 \\citep{burgay03,lyne} may serve as an important tool in studying  close environment of the neutron stars. In this system a fast recycled $P_A= 22.7$ msec Pulsar A orbits a $P_B=2.77$ sec Pulsar B on a 2.4 hour orbit inclined at $\\sim 87.7^\\circ$ (with uncertanity of the order of unity) to the line of sight. Pulsar A shows $27$ sec {\\it frequency independent} eclipses at the moments of superior conjunction \\citep{kaspi}. In addition, Pulsar B shows orbital-phase dependent variations of intensity. It is also  very weak at the inferior  conjunction, consistent with being  eclipsed as well (S. Ransom, private communication).  In this paper we first  construct a geometrical model of  Pulsar A eclipses. Similarly to Arons \\etal (private communication) we propose that the  interaction of Pulsar A wind with   Pulsar B magnetosphere resembles Solar wind -- Earth magnetosphere interaction (see Fig \\ref{Binarypsr}). As a result, a bow shock forms in Pulsar A wind, separated from Pulsar B magnetosphere by a sheath of shocked material.  Absorption of Pulsar A radio emission  may occur either in the shocked Pulsar A wind or in the magnetosphere of Pulsar B. If   absorption  occurs at cyclotron  resonance with non-relativistic particles present on the  closed field lines, then it is expected to be strongly frequency dependent, which is not observed. Synchrotron absorption by relativistic particles on the open field lines can give a frequency independent eclipse. Still, we disfavor this possibility, First, there are no indication of  modulation of the eclipse by the rotation of  Pulsar B (though the present temporal  resolution may not be sufficient). Secondly, since the line of sight passes at the edge of magnetosphere (see Fig.  \\ref{pulsarB}), and since we see pulses from Pulsar B, it is hard to imagine geometry in which near the superior  conjunction the line of light constantly  passes through open field lines, giving full eclipses. Thirdly, the dimming of Pulsar  B near the inferior conjunction is consistent with absorption being  due to the shocked Pulsar A wind.   An alternative possibility for the location of the absorbing material, which we favor, is that the absorption occurs in the shocked Pulsar A wind, enveloping  Pulsar B. Though the size of the shocked sheath   is expected to vary by about $25\\%$ due to Pulsar B rotation, it  may still have an approximately constant column density and optical depth  regardless of its size.  What determines the size of the eclipsing region? It may be determined either by geometric factors, so that the  beginning and the end  of the eclipse occur at the moment when the line of sight enters the absorbing region, or by microphysical processes in the plasma so that the eclipse  occurs at the moment when  optical depth  $\\tau$ to scattering/absorption becomes of the order of unity. We strongly prefer the geometrical interpretation of  eclipses since  any relevant  microphysical process gives  a frequency dependent absorption coefficient so that the condition $\\tau \\sim 1$ would  occur at different points for different frequencies.  The geometrical (as opposed to microphysical) interpretation  of the eclipse duration seems, at first sight, inconsistent with the data: a simple scaling of the size of Pulsar B magnetosphere would predict an eclipse which is much longer than observed. The light cylinder radius of  Pulsar B is $1.29 \\times 10^{10}$ cm, so that its angular size seen from Pulsar A is $0.14 $ rad. The line  of sight passes $2.3^\\circ=0.04$ rad $=3.6 \\times 10^9$ cm away from Pulsar B. Eclipsing part, corresponding to a 27 sec eclipse duration, is $1.8 \\times 10^9$ cm \\cite{kaspi}, \\footnote{ At this point magnetic field of the B pulsar $B_B = 13$ G, so that the cyclotron frequency $\\nu_B = 4.4\\times 10^7$ Hz, which is too low to produce cyclotron absorption if particles are non-relativistic.} which subtends a total  angle of $0.02$ radians (see Fig. \\ref{geom}). Thus, the size of the eclipsing region is $\\sim 0.04$ rad $= 4 \\times 10^9$ cm, which is more than 3 times smaller than the light cylinder radius of  Pulsar B.  A possible resolution of this disagreement is that the pressure of Pulsar A wind compresses Pulsar B magnetosphere to  lateral sizes much smaller than the light cylinder. In this paper we first calculate semi-analytically the form of the interface between  Pulsar A wind and Pulsar B magnetosphere. We find that if the magnetic field of  Pulsar B is calculated using the  vacuum dipole formula, the  lateral size of the  sheath is still much larger than the one  inferred from eclipse duration. On the other hand, electromagnetic interaction of  Pulsar B with the sheath strongly modifies the structure of the magnetosphere, so that the vacuum dipole formula becomes inapplicable. The modification of the spin down torque leads to an estimate of   Pulsar B magnetic field which is typically 3-5 times smaller. Still, the size of the sheath is much larger than inferred from eclipse duration. A possible resolution is that inclination angle is $\\leq 86^\\circ$, so that the line of sight cuts at the edge of the sheath.  As  independent parameters of the system we chose the  spin down luminosities $L_A = 5.8 \\times 10^{33} $ erg/sec and $L_B = 1.6 \\times 10^{30}$  erg/sec and the separation of $D=9\\times 10^{10}$ cm. These are the most reliable parameters, while  the quantities like the inferred  surface magnetic fields depend on the details of the magnetospheric structure. ",
        "conclusions": "In this paper we considered magnetohydrodynamic interaction between relativistic pulsar wind and static magnetosphere in  double pulsar system PSR J0737-3039. Our main conclusion are: \\begin{itemize} \\item   Electromagnetic torque on  Pulsar B is  increased due to the interaction with Pulsar A wind. Depending on the nature of interaction, the magnetic field of Pulsar B can be as low as $3 \\times 10^{11} $ G. Still, the geometrical  model for the form of the interface   predicts the size of the interface which is much larger than inferred from the duration of Pulsar A eclipses, unless the orbital  inclination angle is $\\leq 86^\\circ$. \\item A likely cause of eclipses is  synchrotron absorption in the dense shocked Pulsar A wind by low energy relativistic electrons.  The density of Pulsar A wind should be at the upper allowed limit. The presence of such electrons cannot be justified from first principles. \\end{itemize}   The major uncertainty of the model is the source of low energy, $\\gamma \\sim 30$, electrons in the sheath. Pulsar A wind cannot be so slow: otherwise induced Compton scattering in the wind will make it unobservable \\citep{sk92}. One  possibility is  that such low energy electrons are accelerate not at the shock, but at the rotational discontinuity during development of dynamic and/or resistive instabilities. Alternatively, mixing of Pulsar A and Pulsar B plasmas, initiated by such instabilities, may populate the sheath with weakly relativistic electrons accelerated in Pulsar B gaps. If  mixing is efficient, particle number density in the sheath may be dominated by  Pulsar B plasmas.  Yet another possibility is that absorption happens on the strongly bend-back open field lines of Pulsar B which asymptotically take the form of the sheath (it would be hard to distinguish this possibility from absorption in the sheath itself).    Our calculations are consistent with the possibility that Pulsar A wind experiences  a strong shock near Pulsar B. This limits the magnetization parameter to $\\sigma \\leq \\Gamma^2$ (this is a condition that the preshock Lorentz factor $\\Gamma$ is larger than the \\Alfven speed in the wind). Thus,  our results  {do not necessarily imply that  Pulsar A wind is kinetically dominated near Pulsar B}. (The energy required to create a population of low energy electrons is a $1/\\sigma$ fraction of the total wind energy; this is also a fraction of energy dissipated in a strongly magnetized shock with $\\sigma \\gg 1$, \\citep{kc84a}.)   A possible measurement of $\\sigma$ may come from high energy observations of the system. Shocks in kinetically dominated winds are expected to be strongly dissipative so that a large fraction of the incoming energy flux may be radiated in X-rays. The expected luminosity is $L_X \\sim L_A \\Delta \\Omega\\sim 3 \\times 10^{30}$ erg/sec, where $\\Delta \\Omega \\sim 5 \\times 10^{-3}$ is the solid angle of  the shocked region seen from Pulsar A.  A recent detection of a weak X-ray emission \\citep{mclaugh04} at a level $2 \\times 10^{30}$ erg/sec is consistent with this scenario. (A confirmation of this result is needed since the count rate was too low and the X-ray emission is also consistent with Pulsar A magnetospheric emission.)   A better understanding of the system should come from the full relativistic MHD modeling of the system. The  relations for oblique relativistic MHD shocks  derived in this paper can be used  as guidelines and a  check to  such simulations. Qualitatively, the flow in  the sheath is expected to become supersonic both due to the changing shock conditions and due to pressure acceleration away from the apex point, so that the flow will form a de Laval nozzle. In addition, centrifugal forces acting on a flow moving along the curved trajectory (so called Busemann correction) need to be taken into account. Another important property of the shocked flow is that even for small magnetization of Pulsar A wind $\\sigma \\ll 1$ {\\it magnetic field plays an important role near the  contact discontinuity} and cannot be neglected.  Thus, any credible simulation of the interaction {\\it  must } use full relativistic MHD formalism."
    },
    "0403/astro-ph0403595_arXiv.txt": {
        "abstract": "We report on systematic radial velocity surveys for white dwarf -- white dwarf binaries (double degenerates -- DDs) including SPY (ESO Supernovae Ia progenitor survey) recently carried out at the VLT. A large sample of DD will allow us to put strong constrains on the phases of close binary evolution of the progenitor systems and to perform an observational test of the DD scenario for supernovae of type Ia. We explain how parameters of the binaries can be derived from various methods. Results for a sample of DDs are presented and discussed. ",
        "introduction": "Supernovae of type Ia (SN\\,Ia) play an outstanding role for our understanding of galactic evolution and the determination of the extragalactic distance scale.  However, the nature of their progenitors is not yet settled (e.g.\\ Livio 2000). SN\\,Ia are observed in all types of galaxies, including elliptical galaxies containing only old stellar populations. The light curves of SN\\,Ia are dominated by the decay of radioactive material synthesised in the explosion (mainly nickel, decaying to iron). The rapid evolution of SN\\,Ia light curves indicates that the precursors of these supernovae must be compact objects of small mass with little mass holding back the gamma-rays produced by the radioactive decay. The only candidate, which can fulfil these observational constraints, is the thermonuclear explosion of a white dwarf (WD).  According to the current consensus this happens when the WD grows to the Chandrasekhar mass of $\\approx$$1.4M_\\odot$. Since no way is known how this can happen to a single WD, this can only be achieved by mass transfer in a binary system. Several channels have been identified as possibly yielding such a critical mass. They can be broadly grouped into two classes. The single degenerate (SD) channel (Whelan \\& Iben 1973) in which the WD is accompanied by either a main sequence star, a (super)giant, or a helium star, as mass donor and the double degenerate (DD) channel where the companion is another WD (Webbink 1984; Iben \\& Tutukov 1984). Close DDs radiate gravitational waves, which results in a shrinking orbit due to the loss of energy and angular momentum. If the initial separation is close enough (orbital periods below $\\approx$10\\,h), a DD system could merge within a Hubble time, and if the combined mass exceeds the Chandrasekhar limit the DD would qualify as a potential SN\\,Ia progenitor. ",
        "conclusions": "The large programme part of SPY has now been completed with some observations underway to complete the observations of the WDs with only one spectrum taken during the survey.  We increased the number of WDs checked for RV variability from 200 to 1000 and multiplied the number of known DDs by more than a factor of six (from 18 to $\\approx$120) compared to the results achieved during the last 20 years.  Our sample includes many short period binaries (Fig.~\\ref{f:pm}), several with masses closer to the Chandrasekhar limit than any system known before, greatly improving the statistics of DDs. We expect this survey to produce a sample of $\\approx$150 DDs.  This will allow us not only to find several of the long sought potential SN~Ia precursors (if they are DDs), but will also provide a census of the final binary configurations, hence an important test for the theory of close binary star evolution after mass and angular momentum losses through winds and common envelope phases, which are very difficult to model. An empirical calibration provides the most promising approach. A large sample of binary WDs covering a wide range in parameter space is the most important ingredient for this task.  Our ongoing follow-up observations already revealed the existence of three short period systems with masses close to the Chandrasekhar limit, which will merge within 4\\,Gyrs to two Hubble times. Even if it will finally turn out that the mass of our most promising SN\\,Ia progenitor candidate system is slightly below the Chandrasekhar limit, our results already allow a qualitative evaluation of the DD channel. Since the formation of a system slightly below Chandrasekhar limit is not very different from the formation of a system above this limit, the presence of these three systems alone provides evidence (although not final proof) that potential DD progenitors of SN\\,Ia do exist."
    },
    "0403/hep-ph0403013_arXiv.txt": {
        "abstract": "Baryon inhomogeneitites formed in the early universe are important as they affect the nucleosynthesis calculations. Since they are formed much before the nucleosynthesis epoch, neutrino inflation plays a crucial role in damping out these fluctuations. Now neutrinos, in turn, are affected by magnetic fields which may be present in the early universe. In this work we study the evolution of baryonic inhomogeneities due to neutrino induced dissipative processes in the presence of a background magnetic field. We find that at higher temperatures the dissipation of the inhomogeneities are enhanced as the magnetic field increases. Our study also shows that at lower temperatures the same magnetic field may produce less dissipation. Though we limit our study to temperatures below the quark-hadron transition we do establish that magnetic fields present in the early universe  affect the dissipation of baryonic inhomogeneities. ",
        "introduction": "Baryon inhomogeneities are formed during various epochs in the early universe \\cite{inhomo}. These inhomogeneities are important as their presence during the nucleosynthesis epoch affects the neutron to proton ratio and thus the calculated abundances of the light elements. These baryon inhomogeneities get dissipated as the universe gets cooler by the combined effects of neutrino inflation (before 1 MeV) and baryon diffusion ( after neutrino decoupling). The baryon inhomogeneities have a higher baryon number density than the surrounding plasma, so to maintain pressure equilibrium with the plasma they have to have a lower temperature than the surrounding plasma.  As long as neutrinos are free streaming (over the lengthscale of the inhomogeneties) they pass from the high temperature plasma to the low temperature inhomogeneities and deposit heat in them. The deposition of heat disturbs the pressure balance and the inhomogeneity increases in volume to achieve pressure equilibrium again. Thus the volume of the inhomogeneity increases while its amplitude goes down due to transfer of heat by the neutrinos. Detailed study of the dissipation  of these baryon fluctuations have been carried out \\cite{jedamzik,heckler} previously. As seen in these studies neutrino inflation reduces the inhomogeneities but in most cases is unable to wipe them out completely. This is especially true if the inhomogeneity is large. Smaller inhomogeneities are however wiped out by neutrino inflation.  The studies done previously have been done without the presence of a background magnetic fields. There are ample evidences to show that the early universe had large magnetic fields \\cite{kronberg}. Recent studies have also shown that neutrinos though uncharged are affected by magnetic fields especially in the early universe plasma \\cite{kaushik}. Magnetic fields increase heat loss from the background plasma by the enhanced creation of neutrino pairs by e+e- annihilation, neutrino synchrotron radiation and other processes. These neutrinos pass through the baryon inhomogeneties and increase the total heat deposited in them. In this work, we study the effect of this on the baryon inhomogeneities.  Heat loss from the background plasma by neutrinos in presence of magnetic field has been studied with respect to  neutron stars and supernovae explosions. These studies are thus for much lower temperatures. ($T= 10^9 K$). Since neutrino inflation of baryon inhomogeneities occur at much higher temperatures($T>>10^9 K$), and  very few studies have been done for such temperatures; we believe that our work will also provide a motivation to carry out further studies of neutrino effects in higher temperature plasmas in the presence of magnetic fields.  In section II we first review heat deposition in baryon inhomogeneities, then in section III, we review neutrinos in high temperature magnetic plasmas. In section IV, we present our calculations for the dissipation of baryonic inhomogeneities, due to neutrinos in the presence of a magnetic field and give our results. In section V, we briefly discuss how we can get such strong magnetic fields before the nucleosynthesis epoch. Section VI presents our conclusions and some discussions regarding the approximations considered in this work. ",
        "conclusions": "In conclusion we have established that strong magnetic fields do affect the inflation of baryon inhomogeneities by neutrino heating. We have seen that this effect is greater at higher temperatures and for non-quantizing fields. Our results seem to indicate that for the same value of the field, a non-quantizing field has more effect than a quantizing field. This is expected, since as mentioned before, for the quantizing field, most of the plasma particles are in the ground Landau level and therefore are incapable of emitting synchrotron radiation. This reduces the number of neutrinos depositing heat in the inhomogenities. Now,  whether the field would be quantizing or not depends on the temperature. Since at higher temperatures, even strong magnetic fields are non-quantizing, hence the effect described here is greater at higher temperatures and stronger magnetic fields.  We have simplified our calculations using certain approximations. The temperature difference between the inhomogeneity is considered to be constant. However the temperature difference is actually related to the amplitude of the inhomogeneity and changes accordingly with it. But this change is very small. A proper investigation would involve a detailed simulation where the small temperature change should be taken into account.  The result of keeping this constant is that we get the largest possible effect. As the amplitude of the inhomogeneity decreases, the temperature difference also decreases but as our results show the inhomogeneity does not decrease very rapidly so we feel that our results will not change very much even if we take the small temperature changes into account. We have also taken the size of the inhomogeneities equal to the mean free path of neutrinos at that temperature. If the inhomogeneities are larger then the neutrinos cannot penetrate the entire inhomogeneity and will affect only the edges of the inhomogeneities. However since neutrino mean free path increases as temperature decreases, most of the inhomogeneities within the horizon will have their sizes either equal to or smaller than the neutrino mean free path. For smaller size inhomogeneities, the heat deposited will be less by a factor given by the ratio of the size of the inhomogeneity to the mean free path of the neutrinos at that temperature.     We have considered baryon inhomogenities present after the quark-hadron transition but there is the possibility of inhomogenities being present much before that. Since these inhomogenities may be created anytime after the electroweak phase transition, a much through investigation of their evolution should include the high temperature zone ($100$ GeV -$200$ MeV) also. For this one has to study neutrino interaction properties at much  higher temperatures, in the presence of a magnetic field. To our knowledge, such a study has not been carried out. We hope that this work will provide a motivation to study neutrino emission in a magnetic field at GeV temperatures and its subsequent effect on the early universe.   \\vskip .2in \\centerline {\\bf ACKNOWLEDGMENTS} \\vskip .1in I would like to thank Ajit Srivastava, Biswanath Layek and Kaushik Bhattacharya for their suggestions and comments. I would also like to thank Charles Horowitz, Adam Szczepaniak and the Nuclear Theory Center, Indiana University, Bloomington for their support and hospitality."
    },
    "0403/astro-ph0403526_arXiv.txt": {
        "abstract": "Good tracking requires that the quintessence energy fraction slowly increase while the roll $\\lambda\\equiv -d\\ln V/\\varkappa d\\phi$ slowly decreases, but is not yet truly slow-rolling.  The supernova bound on the present quintessence equation of state requires either (1) a cosmological constant or other fine-tuned \"crawling quintessence\" or (2) \"roll-over quintessence\" that tracked until recently, but now became slow rolling, because of a sharp increase in potential curvature.  Thus, fine-tuning is required by constant equation of state and inverse power potentials, but can be avoided by the SUGRA and Skordis-Albrecht potentials and other good trackers, provided quintessence energy domination and slow roll {\\em both} began only recently.  This makes the time in which we live special in {\\em two} respects. ",
        "introduction": "The most surprising and important recent discovery in cosmology is that the present universe is flat and dominated by unclustered dark energy, whose energy density $\\rhoq$ is now small and constant or nearly constant.  Indeed, the next few years should show whether this   energy density is now truly static (cosmological constant) or slowly evolving with the logarithm of the cosmological scale factor $N\\equiv \\ln a=-\\ln(1+z)$.  If dynamic, we would like to know something about the dynamical equation of state $\\gamma_Q(N)\\equiv -d\\ln \\rhoq/3 dN \\neq 0$.  In this section, we review the present bounds on dark energy evolution, before reviewing the constraints on {\\em tracking quintessence}, a particularly interesting model for the smooth energy in terms of a very low mass canonical scalar field.   \\subsection{Kinematics of the Accelerating Flat FRW Universe}  Supernovae Ia and cosmic shear directly explore the space-time geometry by measuring the luminosity distance $d_L(z)=(1+z)\\eta$ to individual distant supernovae, chosen to be standard candles, and the angular-diameter distance $d_A(z)=\\eta/(1+z)$ to distant galaxies, from which the comoving distance \\be\\eta\\equiv c\\int_0^z dz'/H(z')=c\\int_0^t dt'/a(t')\\ee  is inferred. This conformal coordinate distance to the horizon, $\\eta$, describes the proper time evolution of the scale factor $a$.  The need for unclustered dark energy derives from general relativity, which is tested to high precision in the classic solar system and binary pulsar observations.  General relativity is also confirmed out to cosmological distances and down to ages about one minute by BBN light-element abundances, and is consistent with the observed spectra of primordial density fluctuations, CBR temperature anisotropies, and large-scale structure power.  We will therefore ignore alternate gravity theories, such as small-curvature modifications of GR, large extra spatial dimensions or brane cosmology, and assume an homogeneous and isotropic (Friedmann-Robertson-Walker) flat universe, the Friedmann expansion rate is \\be 8\\pi G\\rho=3 H^2 , \\qquad \\mbox{where} \\quad H \\equiv \\dot a/a. \\ee Quantum field theory requires that the energy density, and therefore $G$, be positive, so that we can write $\\sqrt{8\\pi G}\\equiv\\varkappa\\equiv 1/M_P$, where $M_P=2.44e18~GeV$ is the reduced Planck mass. The derived quantity, the cosmological fluid pressure $P=-d (\\rho c^2 a^3)/d a^3 $, may be positive or negative.  In terms of geometrical quantities, \\be\\varkappa^2 P/c^2=-(2\\dot H+3 H^2) ,\\ee the enthalpy is \\be \\varkappa^2(\\rho+P/c^2)=-2\\dot H=-d H^2/d N=-(\\varkappa^2 /3)(d\\rho/dN), \\ee and the over-all barotropic index is \\be\\gamma \\equiv -d\\ln\\rho/3 dN=(\\rho+P/c^2)/\\rho= -\\frac{2}{3}(d\\ln H/dN).\\ee Here the logarithm of the cosmological scale factor $N\\equiv \\ln a=-\\ln(1+z)$, so that $dN=H dt$. This equation of state and its quintessence component depends on the Hubble time $H^{-1}=d\\eta/dz$, and its derivative $d H^{-1}/d z=d^2\\eta/dz^2$.  In flat FRW cosmology, the space-time curvature (Ricci scalar) is \\be R\\equiv -6(\\dot H+2 H^2)=\\varkappa^2 (3 P-\\rho),\\ee while the acceleration, $\\ddot a/a=-\\varkappa^2 (\\rho+3 P)/6$, so that \\be \\ddot a a/{\\dot a}^2=1- d(H^{-1})/d t=-(1+3 w)/2 \\ee ranges from $-2$ to $1$, when the overall equation of state $w\\equiv P/\\rho$ ranges from the stiff value $1$ down to the ultra-soft value $-1$ . We now know that, until about red-shift $z\\sim 0.5$, attractive gravity dominated the cosmological fluid so that large scale structures could form. Only recently, after the expansion rate in equation (7) outpaced the growth of the Hubble radius $H^{-1}$, did $P/\\rho=<-1/3$, and the expansion become accelerated.  \\subsection{Present Cosmological Parameters}  Supernova observations were the first evidence for smooth energy and remain the direct observation of acceleration of the present universe.  This supports the possibility of an inflationary early universe, which is otherwise not directly observable.  Because the barotropic index (5) and its quintessence component (13) depend on the first and second derivatives of the comoving distance $\\eta$, the quintessence evolution $\\dub(z)$ depends on first and second derivatives of the observed luminosity distances \\cite{Moar} or on $H(z),~d H(z)/dz$. In practice, quintessence is appreciable only for small red-shift. This means that, before $\\dub(z)$ can be extracted from the supernova data, the inherently noisy luminosity distance $d_L(z)$ data must be parametrized. For this, and other reasons, along with a large number of high red-shift supernovae, precise knowledge of other cosmological parameters will be needed \\cite{Naka,Hut,Frieman}, and can still determine only one or two parameters characterizing the potential, such as $w_{Q0},~(d\\dub/dz)_0$.  A worrisome feature $^{5)}$ of the present supernova data is that, taken separately, the low $z<0.34$ and high $z>0.34$ observations imply best-fit supernovae Ia absolute magnitudes differing by 0.5 magnitude and are each compatible with an Einstein-Lemaitre $\\Omega_{Q0}=0$.  Indeed, the combined data actually favors a supernova-averaged quintessence equation of state \\[\\widetilde{\\dub}(N)\\equiv\\int_N^0 \\dub(N')\\,d N'/N <-1.\\]  While programs to measure luminosity and angular diameter distances are underway, we already know that we live at a time when \\be\\Omega_{Q0}=0.71\\pm 0.07 , \\qquad \\widetilde{\\dub}<-0.78~(95\\%~CL), \\qquad h\\equiv H_0/100=0.72\\pm 0.05 ,\\ee that the radiation/matter equality took place at red-shift $z_{eq}=3454^{+385}_{-392}$ \\cite{Spergel,Tonry}, dark energy began dominating over matter $\\gtrsim 6.3~Gyr$ ago, and the cosmological expansion has been accelerating since red-shift $z\\sim 0.7$ \\cite{Tonry,Blud}.  This fit to $\\widetilde{\\dub}$ derives from SNIa, 2dFGRS and Ly$\\alpha$ observations, and assumes $\\dub$ constant and $>-1$.  Because the observations average over a small range in $z$, in which $\\dub (z)$ changes relatively little, $\\widetilde{\\dub}<-0.78$ is also a rough upper bound to the present value $w_{Q0}$.  The radiation/matter background density, \\be \\rho_B=(11.67 a+0.003378)/a^4\\qquad meV^4 ,\\ee is now $\\rho_{B0}=11.67 ~meV^4$ and was $\\rho_{Bi}=0.003378~GeV^4$ at fiducial red-shift $z=10^{12}$.  The supernova observations fit an average \\[\\widetilde{\\dub}(N)\\equiv\\int_0^N \\dub(N')\\,d N'/N ,\\] The CBR anisotropy and mass fluctuation spectrum, on the other hand, depend on \"early quintessence\" \\cite{Cald}, back to the last scattering surface $z\\sim 1100$. Where needed, we will fix $h^2=1/2$, so that the present critical density and smooth energy density are $\\rho_{cr0}=40.5~meV^4,~\\rho_{Q0}=28.8~meV^4=(2.3~meV)^4$.  This completes our discussion of the geometry of an accelerated flat universe.  The next two sections review scalar field dynamics and the two Attractor Conditions, extending the many earlier optimistic treatments of the Tracker Condition \\cite{CDS,SWZ,ZWS,Fer,Liddle,Cope,Ng} and of inverse power potentials (Section IV) \\cite{Ratra,Wett}, by considering (1) poor trackers, (2) post-tracker behavior in the present quintessence-dominated era, (3) the range of initial conditions leading to tracking, and (4) the numerical problems encountered in cosmological dynamics, particularly in the freezing and tracking epochs.  Sections IV and V conclude that the nearly static upper bound (8) requires that any quintessence must be either finely-tuned (crawling quintessence) or have large current curvature (cross-over quintessence \\cite{Wett2}) . ",
        "conclusions": ""
    },
    "0403/astro-ph0403460_arXiv.txt": {
        "abstract": "{We study the possibility of a cool disk existing in the Galactic Center in the framework of the disk-corona evaporation/condensation model. Assuming an inactive disk near the gravitational capture distance left over from an earlier evolutionary stage, a hot corona should form above the disk since there is a continuous supply of hot gas from stellar winds of the close-by massive stars. We study the interaction between the disk and the corona. Whether the cool disk can survive depends on the mass exchange between disk and corona which is determined by the energy and pressure balance. If evaporation is the dominant process and the rate is larger than the Bondi accretion rate in the Galactic Center, the disk will be depleted within a certain time and no persistent disk will exist. On the other hand, if the interaction results in hot gas steadily condensing into the disk, an inactive cool disk with little gas accreting towards the central black hole might survive in the Galactic Center. For this case we further investigate the Bremsstrahlung radiation from the hot corona and compare it with the observed X-ray luminosity. Our model shows that, for standard viscosity in the corona ($\\alpha=0.3$), the mass evaporation rate is much higher than the Bondi accretion rate and the coronal density is much larger than that inferred from Chandra observations. An inactive disk can not survive such strong evaporation. For small viscosity ($\\alpha \\la 0.07$) we find condensation solutions. But detailed coronal structure computations show that in this case there is too much X-ray radiation from the corona to be in agreement with the observations. From this modeling we conclude that there should be no thin/inactive disk presently in the Galactic Center.  However we do not exclude that the alternative non-radiative model of Nayakshin (2004) might instead be realized in nature and shortly discuss this question. ",
        "introduction": "The case for a massive black hole in our Galactic Center coincident with the radio source Sgr A$^*$ represents a unique opportunity to probe the dynamics of gas accreting onto a massive black hole. Dynamical measurements of stellar velocities within the central 0.1 parsec of the Galactic Center from 10 years of high resolution imaging indicate a central mass of $2.6\\times10^6M_\\odot$ (Genzel et al. 2000, Ghez et al. 2000) or $3.3\\times10^6M_\\odot$ (Sch\\\"odel et al. 2003), recent measurements of stellar orbits (Genzel et al. 2003) provide evidence for a $3.6\\times10^6M_\\odot$ black hole. The Eddington luminosity for $3.3\\times10^6M_\\odot$ is $L_{\\rm Edd}\\approx 10^{44.6}{\\rm erg\\, s}^{-1}$. Observations show that Sgr A$^*$ is an extremely dim galactic  nucleus. The luminosity in the submillimeter/far infrared region is $L_{\\rm submm}\\sim 10^{-8.6}L_{\\rm Edd} $, even less in the infrared band with $L_{\\rm IR}< 10^{-9.6}L_{\\rm Edd} $, also very low is the quiescent X-ray luminosity, $L_{\\rm X}< 10^{-11}L_{\\rm Edd}$ (Narayan 2002, where $L_{\\rm Edd}=10^{44.6}{\\rm erg\\, s}^{-1}$). Chandra observations directly image the hot X-ray-emitting thermal gas in the vicinity of the Bondi accretion radius where the surrounding gas is captured by the gravitational pull of the central black hole, and determine temperatures and densities that allow to estimate a mass accretion rate of Sgr A$^*$  of $\\dot M_{\\rm{Bondi}}\\sim (0.3-1)\\times 10^{-5}M_\\odot$/yr, equivalent to$\\sim 10^{-4}\\dot M_{\\rm Edd}$ (with $\\dot M_{\\rm Edd}\\equiv L_{\\rm Edd}/0.1 c^2$) (e.g. Baganoff et al. 2003). If mass flows steadily through a Shakura-Sunyaev thin disk (Shakura \\& Sunyaev 1973) to the central black hole at this accretion rate, the thin disk model would predict a luminosity $L_{\\rm disk}\\sim 0.1 \\dot M_{\\rm Bondi} c^2\\sim 10^{40.8}{\\rm erg\\, s}^{-1}\\sim 10^{-4}L_{\\rm Edd} $, much higher than the luminosity in any band observed. Thus, a standard thin disk model appears to be ruled out.   The low luminosity was explained by an advection-dominated accretion flow (ADAF), with a spectral fit first presented by Narayan et al. (1995). In the following years important observational results on the emission of Sgr A* and theoretical work lead to an improved model for radiatively inefficient accretion flows (RIAFs). RIAF models for Sgr A* are discussed in the recent reviews of Yuan, Quataert \\& Narayan (2003) and Quataert (2003). According to these models, most of the thermal energy released by viscosity and increased by compression is retained in the gas and advected to the central black hole. The RIAF model in addition assumes that very little mass from large radii actually accretes onto the black hole while a large part is lost through outflows during the accretion. These models naturally yield the observed spectra of Sgr A$^*$. A key constraint on these models is that the fraction of gravitational energy heating the electrons must be very small and hence a two-temperature treatment of the plasma is required.   Another possibility to explain the low luminosity of Sgr A* might come from the existence of cold molecular gas in the parsec region of the Galactic Center in the form of an inactive/dead thin disk without accretion. Falcke \\& Melia (1997) had suggested such a ``fossil'' disk (for a review see Melia \\& Falcke 2001). The model assumes that gas captured at Bondi accretion rate condenses onto such an inactive disk without mass accretion onto the black hole. This runs into difficulties because the inflowing gas produces a fair amount of luminosity in the infrared (Narayan 2002) as it clashes onto the disk and loses its thermal and kinetic energy. This radiation is not seen. Nayakshin (2004) revisits the concept of an inactive disk. He suggests that, in the case of an extraordinarily long mean free path (larger than the pressure scale of  the corona) and extremely low  viscosity of the hot gas, the energy can be conducted to a very thin transition layer by free streaming electrons and is then radiated in infrared to UV wavelengths. As this radiation from the thin layer would be observed edge-on a discrepancy between predicted and observed luminosities could be avoided. This interesting suggestion deserves further analysis of the transition between the hot gas and the cool layer and the coupling between the hot ions and the energy transferring electrons.  As a further contribution to the issue of a cool disk around the Galactic Center we here study the vertical structure of a hot corona in interaction with a disk below. We assume a cool disk in the outer region around the circularization radius where the free fall of a Bondi type accreting hot gas goes over into a circular motion around the gravitational center due to its specific angular momentum. Such a disk might not be unreasonable since the system must have been quite bright during an earlier evolutionary stage and angular momentum that was released by the high accretion rate should have moved disk mass into outward regions. With winds from young stars being captured by the gravitational field of the black hole, a corona unavoidably forms above the cool disk. The question is, how do disk and corona evolve? Does the hot gas condense to the disk with little mass actually accreting to the black hole, or does mass rather evaporate from the disk to the corona overwhelming the incoming hot gas and even finally depleting the cool disk underneath?   The answer depends mainly on the rate at which the hot gas is supplied from the capture radius. If gas is supplied to the corona at a sufficiently high rate coronal gas condenses to the cool disk. If no gas is supplied from the outside or if the outside gas supply is too small mass instead evaporates from the disk into the corona. Both processes are the consequence of pressure and energy equilibrium between the disk and the corona (Meyer et al. 2000, Liu et al. 2002). Here we study in detail the structure of the corona for the case of the Galactic Center in order to see what the dominant process between disk and corona is, condensation or evaporation.  Can the cool disk survive if evaporation is dominant?  Furthermore, the detailed computation allows to calculate the Bremsstrahlung luminosity of the corona and compare it with the observed X-ray luminosity. We show that these results exclude the existence of any cool disk in our Galactic Center.  In Sect.2 we describe the physics of the interaction between the disk and the corona. In particular we discuss how a radial inflow of mass from the outside affects the mass and energy balance in such a corona. In Sect.3 we present numerical results and show how the value of the viscosity affects evaporation or condensation. We discuss several aspects of our results in Sect. 4, including a comparison with a non-radiative condensation model of Nayakshin (2004) and the question of a past disk being evaporated now in the Galactic Center. A conclusion follows in Sect. 5. ",
        "conclusions": "We have developed a model for a cool disk around the Galactic Center that has a corona above it, allowing for a coronal accretion of gas captured at the Bondi radius from stellar winds of massive stars. Our models incorporate hydrostatic layering, thermal heat conduction, friction, and radiative energy loss and allow for mass exchange between disk and corona by condensation or evaporation of gas, as well as mass gain or loss by radial flow in the corona (the latter a generalization of the former one-zone model).   We find that the solutions we obtain depend on the value of the viscosity parameter $\\alpha$. For standard values, $0.1\\la \\alpha \\la 0.3$, so much mass evaporates from the disk into the hot corona that the additional mass flow from the outside is a negligible contribution. The evaporation then with time could lead to the complete disappearance of the original disk. Only for values $\\alpha \\la 0.07$ there might be solutions in which the incoming accreted gas condenses into the disk. However, the vertical structure computations show that in all these solutions the calculated Bremsstrahlung of the corona in the X-ray band by far exceeds the observed luminosity.  From this we conclude that, if our modeling is correct and applicable, at present no cool disk around the Galactic Center exists. We shortly discussed what limit this puts on an inert disk that might have originally remained from a phase of star formation in which the young bright stars presently seen close to the Galactic Center were formed. We also discuss the difference to the alternative non-radiative condensation model of Nayakshin (2004)."
    },
    "0403/astro-ph0403183_arXiv.txt": {
        "abstract": " ",
        "introduction": "Deep imaging in the rest-frame UV has constrained both the evolution of the cosmic star formation rate density \\cite{mada96} and its time integral, the growth of the cosmic stellar mass density \\cite{dick03}.  Short-wavelength studies give an incomplete picture, however, since an important population of high-redshift galaxies is heavily dust-obscured.  The strength of the extragalactic mid- and far-IR/submillimeter background indicates that about half of the cosmic energy density comes from dusty luminous and ultra-luminous infrared galaxies (LIRGs/ULIRGs: $L_{\\rm IR} \\sim 10^{11.5}$ to $10^{13.5}\\,L_\\odot$) at $z \\geq 1$ \\cite{puge96,pei99}.  Because the brightest of these submillimeter galaxies (SMGs; see \\cite{blai02} and references therein) tend to lack strong X-ray emission \\cite{alma03}, their large IR luminosities probably correspond to high star formation rates \\cite{barg00}.  As the strikingly different appearances of the Hubble Deep Field at $0.83\\,{\\rm \\mu m}$ \\cite{will96} and $850\\,{\\rm \\mu m}$ \\cite{hugh98} exemplify, SMGs are rarer and forming stars much more intensely than typical optically selected systems \\cite{will96,hugh98}.  Here we discuss new observations that shed light on the importance of SMGs in the history of galaxy mass assembly (all numbers assuming a flat $\\Omega_\\Lambda = 0.7$ cosmology with $H_0 = 70\\,{\\rm km\\,s^{-1}\\,Mpc^{-1}}$). ",
        "conclusions": ""
    },
    "0403/astro-ph0403656_arXiv.txt": {
        "abstract": "We present a weak lensing analysis and comparison to optical and X-ray maps of the close pair of massive clusters A222/223. Indications for a filamentary connection between the clusters are found and discussed. ",
        "introduction": "\\label{sec:introduction} $N$-body simulations of cosmic structure formation predicts that matter in the universe should be concentrated along sheets and filaments and that clusters of galaxies form where these intersect \\citep{1999MNRAS.303..188K,1996Nature..380..603B}. This filamentary structure, often also dubbed ``cosmic web'', has been seen in galaxy redshift surveys \\citep{1994ApJ...420..525V} and X-rays \\citep{2002A&A...394....7Z,2001ApJ...563..673T}.  Because of the greatly varying mass--to--light $(M/L)$ ratios between rich clusters and groups of galaxies \\citep{1998Proc..astro-ph/9810298T} it is problematic to convert the measured galaxy densities to mass densities without further assumptions. Weak gravitational lensing, which is based on the measurement of shape and orientation parameters of faint background galaxies (FBG), is a model--independent method to determine the surface mass density of clusters and filaments. Due to the finite ellipticities of the unlensed FBG, every weak lensing mass reconstruction is unfortunately an inherently noisy process, and the expected surface mass density of a single filament is too low to be detected with current telescopes \\citep{2000ApJ...530..547J}.  Cosmic web theory also predicts that the surface mass density of a filament increases towards clusters \\citep{1996Nature..380..603B}. Filaments connecting neighboring clusters should have surface mass densities high enough to be detectable with weak lensing \\citep{1998wfsc.conf...61P}. Such filaments may have been detected in several recent weak lensing studies.  \\citet{1998astro.ph..9268K} found a possible filament between two of the three clusters in the $z=0.42$ super-cluster MS0302+17, but the detection remains somewhat uncertain because of a possible foreground structure overlapping the filament and possible edge effects due to the gap between two of the camera chips lying along the filament. Also, \\citet{2004astro.ph..1403G} could not confirm the detection of a filament in this system.  \\citet{2002ApJ...568..141G} claim to have found a filament extending between two of the three clusters of the Abell~901/902 super-cluster, but the significance of this detection is low and subject to possible edge effects, as again the filament is close the intersection of four chips of the camera. ",
        "conclusions": "\\label{sec:discussion} We presented weak lensing, optical galaxy density, and X-ray maps of the massive, close pair of galaxy clusters A~222/223. All maps show a connection between the clusters in the intercluster region. The density distribution of color-selected objects, supplemented by the spectroscopic confirmation of galaxies at the cluster redshift in the intercluster region by P00 and D02, and the significant connection between both clusters in the $0.5 - 2.4$~keV band of ROSAT's PSPC establish a secure filamentary connection between the clusters. The weak lensing indications for a dark matter bridge are much less secure. The signal-to-noise of the structure extending between the clusters is low but the structure itself is very robust and never disappears when the selection criteria for the lensing catalog are varied.  We note that the luminosity density of color-selected galaxies in the intercluster region is only by a factor of $\\sim 2$ lower than in the peaks of the galaxy clusters. If the early type galaxies trace the mass well and a constant $M/L$ ratio is assumed, we would expect a clearly detectable lensing signal along the galaxy density contours. This is obviously not the case. A more detailled comparison of the observed and expected surface mass density, also addressing the apparent misalignment of mass and light in A~223, will be done elsewhere (Dietrich et al. 2004, in preparation). We also note that a misalignment of mass and light is present in the filament candidate of \\citet{2002ApJ...568..141G}.  \\vspace{0.5cm}\\noindent \\emph{Acknowledgment}: This work has been supported by the German Ministry for Science and Education (BMBF) through DESY under the project 05AE2PDA/8, and by the Deutsche Forschungsgemeinschaft under the project SCHN 342/3--1."
    },
    "0403/astro-ph0403327_arXiv.txt": {
        "abstract": "In a recent paper Pell\\'o et al. have reported a candidate $z=10$ galaxy, A1835\\#1916, which was found in a near-infrared survey of the central regions of the gravitational lensing cluster A 1835. If this detection is confirmed and the detection rate turns out to be typical, then the volume averaged ultraviolet emissivity must be rising rapidly with increasing redshift. For a magnification due to gravitational lensing by a factor ${\\cal M}\\ga 25$ estimated by Pell\\'o et al., the inferred star formation rate at $z=10$ would be about one order of magnitude higher than estimates of the star formation rate density at $z=6$.  Objects at $z=10$ would contribute substantially to the total source counts at $1.6 \\mu {\\rm m}$ and the estimated space density of sources may exceed the space density of dark matter haloes in a $\\Lambda$CDM model.  We therefore argue that if A1835\\#1916 is indeed at $z=10$ then either the magnification factor may have been overestimated or the galaxy has a top-heavy initial mass function.  Sources with the UV flux and space density of A1835\\#1916 produce $\\sim 33 f_{\\rm esc} ({\\cal M}/25)$ hydrogen ionising photons per hydrogen atom per Hubble time, where $f_{\\rm esc}$ is the escape fraction of ionising photons.  This rate should be sufficient to reionise most of the diffuse hydrogen in the Universe at redshift ten.  We further use a correlation between the equivalent width and the redshift of the \\Lya emission line with respect to the systemic redshift observed in Lyman break galaxies to obtain constraints on the ionisation state of the surrounding intergalactic medium (IGM) from the Gunn-Peterson absorption. These constraints also argue in favour of the surrounding IGM being fully ionised.  Pell\\'o et al.  may thus have detected a population of sources which is responsibly for the large electron scattering optical depth indicated by the cross-power spectrum of the temperature and polarisation fluctuations of the cosmic microwave background as measured by WMAP. ",
        "introduction": "The unravelling of the reionisation history of the Universe has been the focus of much recent research mainly due to the surprising detection of a large Thompson electron optical depth of $\\tau =0.17\\pm 0.04$ by WMAP \\citep{Kogut:03}.  If correct, this optical depth requires a substantial ionised fraction of hydrogen at redshift $z=10-20$. This result came somewhat as a surprise as the optical depth for \\Lya scattering increases rapidly in the highest redshift QSOs \\citep{Fan:03, Songaila:04} indicating a drop in the emissivity of hydrogen ionising photons \\citep{Miralda:03b}. This has led to the suggestion that the emissivity of ionising photons peaked at high redshift due to a population of early stars or mini-AGN and that the reionisation may have been complicated, with an extended epoch of partial reionisation \\citep{MadauR:03, RicottiO:03} and/or the possibility that hydrogen was reionised twice \\citep[\\eg,][]{Cen:03a, WyitheL:03, CiardiFW:03, RicottiOI:03,Sokasian:03}. The possible detection of a $z=10$ \\Lya emitting star-forming galaxy by \\cite{Pello:04} thus offers exciting prospects to further constrain the reionisation history.  This result, if confirmed by further observations, pushes back the epoch when ``galaxies'' were already shining in Ly$\\alpha$ emission by $\\sim 0.5$\\,Gyr relative to previous detections at $z \\simeq 5 - 6$ \\citep{Rhoads:03, Kodaira:03, Lehnert:03, Hu:04, Stanway:04}. In this letter, we briefly discuss the implications of detecting such a source for the inferred space density of star-forming galaxies, for the emissivity of hydrogen ionising photons and the ionisation state of the IGM at $z=10$.  We assume throughout the cosmology to be the concordance $\\Lambda$ cold dark matter model with ($\\Omega_\\Lambda, \\Omega_{\\rm m}, \\Omega_{\\rm b}, {\\rm h}) = (0.7, 0.3, 0.04, 0.7$) and primordial scale-invariant power spectrum ($n_s=1$) with rms amplitude of the mass fluctuations on scale of 8 \\Mpc, $\\sigma_8=0.91$. ",
        "conclusions": "The implications of the possible detection of a redshift 10 galaxy by \\cite{Pello:04} depend crucially on the assumed magnification factor and on the assumption that the detection of one source of this kind in the effective survey volume is representative. Our estimate of the space density is consistent with the predicted number densities of DM haloes in a LCDM model. For a Salpeter IMF, it requires, however, a rather extreme (close to $100\\%$) star formation efficiency. For a ``top-heavy'' IMF a star formation efficiency $\\ge 10\\%$ would be sufficient. If stellar feedback is important in these objects then the slope of their mass function is expected to be shallower than the mass function of DM haloes in a $\\Lambda$CDM model. A somewhat smaller magnification and/or space density would then be required to be consistent with the model prediction. It seems thus worthwhile to investigate if a smaller magnification is consistent with the uncertainties in the model of the gravitational lens and the assumed cosmological model.  If the magnification was indeed smaller, then the location of the yet missing counter-image would also be much less well constrained.  More objects of this kind are clearly needed for a more solid assessment of their space density and the implied emissivity.  A large star formation rate density and a moderate amplification factor of this source would obviously be good news for ongoing surveys for objects at $z>7$, both behind lensing clusters and in the field.  For the fiducial magnification ${\\cal M} \\sim 25$ and assuming Salpeter IMF, the evolution of the star formation rate density would show a very rapid decrease between $z=10$ and $z=6$ followed by a increase between $z=6$ and $z=4$ which would suggest that A1835\\#1916 is part of a separate population of objects. This raises the question what could have led to the rapid decline of such a population.  A number of authors have made the suggestion that negative feedback due to star formation in shallow potential wells could lead to a rapid decline of a first generation of star-forming objects \\citep[\\eg,][]{Efstathiou:92b, Barkana:00, Ricotti:02}. If instead the IMF becomes top-heavy at $z \\simgt 10$, then the star formation rate density is consistent with the decreasing values observed between $z=4$ and $z=6$. For the fiducial magnification ${\\cal M} \\sim 25$ the emissivity of hydrogen ionising photons emissivity is large enough that the post-overlap state of the reionisation process should have been reached and the neutral fraction of hydrogen should be small. Sources like A1835\\#1916 may thus well be responsible for the large electron scattering optical depth measured by WMAP.  If the magnification is ${\\cal M} \\sim 5$ and/or the space density is overestimated the star formation rate density at $z=10$ could be consistent with a smooth continuation of the observed evolution at lower redshift.  These sources, however, would not be sufficient to reionise at $z=10$ and the low neutral fraction inferred by the detection of the \\lya line would have to be explained in another way, for instance by a fainter population of galaxies or by partial ionisation by X-rays produced by accretion onto intermediate mass black holes in mini-AGNs \\citep{MadauR:03, RicottiO:03}.   The observed \\Lya flux gives independent constraints on the ionisation state of the surrounding IGM.  If the surrounding IGM were not ionised, the strength of the line requires that its centre is redshifted by 650 km s\\m with respect to the velocity of the neutral IGM. This could occur due to the absorption and resonant scattering of the \\Lya photons by a galactic wind, but such a large offset appears unlikely considering the rather shallow potential well that may host this galaxy. The constraints tighten further if the observed correlation between transmission and velocity shift of \\Lya emission in LBGs also holds for A1835\\#1916.  In this case the minimum size of the ionised region consistent with a neutral surrounding IGM is $R_{\\rm S} \\sim 5$ Mpc (comoving), independent of the velocity shift.  We agree with \\cite{Loeb:04} that the source itself is unlikely to ionise such a large region on its own.  However, the large emissivity of ionising photons which is implied by the small effective survey volume, if confirmed, would make the lack of a suppression of the \\Lya emission due to the Gunn-Peterson absorption by the surrounding IGM less surprising.  Clearly, a small neutral fraction at $z=10$ in the diffuse IGM would be further good news for surveys of high-redshift objects which would strongly benefit from \\Lya emitters with large equivalent widths.  A space density as large as inferred here would also mean that a significant fraction of the faintest objects detected at $1.6 \\mu {\\rm m}$ may be at $z\\sim 10$."
    },
    "0403/astro-ph0403261_arXiv.txt": {
        "abstract": "We consider the prospects for a complete survey of the 18,811 sources of the \\rosat\\ All-Sky Survey Bright Source Catalog (BSC) with NASA's \\swift\\ gamma-ray burst (GRB) mission.  By observing each BSC source for 500\\,s with the satellite's imaging X-ray and UV/optical telescopes, this ``Swift Bright (Source) Catalog Survey'' (\\swiftbscs) would derive $<$5\\arcsec\\ source positions for all BSC sources that have not faded substantially from their \\rosat-era flux levels, and provide 0.2--10 keV X-ray and UV or optical flux measurements.  The improvement by a factor of 10 to 30 in the two-dimensional localization for these sources will enable optical identifications or deep limits in nearly every case, fulfilling the promise of the BSC as a multiwavelength catalog, and allowing the full enumeration of its rare source populations.  Since the \\swiftbscs\\ can be accomplished with 10\\% of the time on-orbit, for a three-year mission, and since its targets will be of lower priority than active afterglows, it will not conflict with the primary \\swift\\ mission of GRB follow-up observations.  Moreover, the BSC targets can be scheduled so as to advance the secondary goals of \\swift: Daily monitoring of the full sky in the hard X-ray band (source fluxes of $\\simgt$20\\,mCrab, 10--100\\,keV) with the wide-field Burst Alert Telescope (BAT); and a two-year all-sky BAT survey down to $\\simgt$1\\,mCrab.  The resulting expansion of the catalog of identified X-ray sources from 2000 to 18,000 will provide a greatly-enriched set of targets for observation by \\xmmlong, \\chandra, and future high-energy observatories. ",
        "introduction": "\\label{sec:intro}  NASA's \\swift\\ Midex mission is in its final ground-based testing phases and is scheduled for launch in September 2004\\footnote{\\swift\\ mission web site: \\texttt{http://swift.gsfc.nasa.gov/}}.  Designed for and dedicated to the detection and automated follow-up of gamma-ray bursts (GRBs), \\swift\\ has a broad-band complement of instruments: a coded-mask CdZnTe Burst Alert Telescope (BAT; \\citealt{barthelmy00}) to localize the prompt hard X-ray emission of gamma-ray bursts, and for pointed observations of the subsequent GRB afterglows, a nested-mirror grazing-incidence X-ray telescope (XRT; \\citealt{bhn+03}) and a UV/Optical telescope (UVOT; \\citealt{rtn+00}) which is a slightly-modified version of the Optical Monitor on \\xmmlong\\ \\citep{mbm+01}.  The XRT, with an EPIC MOS CCD as its focal plane detector, has an effective area of 135 \\cmsq\\ at 1.5\\,keV, making it the seventh-largest focusing X-ray telescope ever flown -- including the three telescopes on \\xmm\\ -- and its 18\\arcsec\\ (half-power diameter) point-spread function compares favorably with all previous missions except the \\einstein\\ and \\rosat\\ HRI instruments, \\chandra, and \\xmm.  The UVOT provides UV/optical (200\\,nm to 600\\,nm, depending on the filter) sensitivity over most of the field-of-view of the XRT, and can reach $B = 24$~mag in a 1000\\,s exposure.  Uniquely among missions of its size, \\swift\\ possesses large reaction wheels which give it a slew rate of one degree per second, enabling prompt follow-up of GRB positions and a continuous sequence of pointings over the course of each orbit, scheduled so that the narrow-field instruments (XRT and UVOT) are never Earth-occulted.  In combination, as we explore below, these capabilities make it attractive and feasible to execute a complete \\swift\\ survey of the \\rosat\\ Bright Source Catalog.  The \\rosat\\ All-Sky Survey (RASS) Bright Source Catalog (BSC; \\citealt{vab+99}) contains 18,811 sources, covering 92\\% of the sky, down to count rates of 0.05 PSPC \\ctsec, corresponding to a flux limit of roughly \\tee{7}{-14} \\ergcms\\ (0.5--2.0 keV) for an unabsorbed blackbody spectrum with $kT=0.1$~keV.  It is the most sensitive X-ray all-sky survey ever performed, reaching fluxes 100~times fainter than the \\einstein\\ survey, and with a greater emphasis on the soft X-ray band.  Typical flux-dependent positional uncertainties for BSC sources are 10--20\\arcsec\\ (radius, one-sigma).  These positional uncertainties are large enough that they have prevented confident optical identification -- and hence, physical classification -- of the great majority of BSC sources.  Firm identifications have been derived for a subset of roughly 2000 of the brightest sources by the \\rosat\\ Bright Survey group \\citep{fhs+98,shl+00}. Suggested identifications have been proposed for perhaps 4000 more in wide-area optical surveys \\citep{zeh+03,avm+03}, but the confidence level of these is such that most have not been subject to subsequent follow-up observations.  Among the $\\simgt$12,000 unidentified BSC sources that remain there must exist, undiscovered and in quantity, members of rare and interesting X-ray source populations; members of common source populations that exhibit distinctive features making them worthy of further study; and quite possibly, entirely new classes of Galactic or extragalactic X-ray source that have so far escaped notice.  These rare and interesting source populations are best identified from within the BSC rather than by other methods, for at least two reasons. First, any BSC source has a sufficiently high X-ray flux to permit detailed follow-up observations by \\xmm, \\chandra, and future high-energy missions.  Second, as we discuss below, the optical counterparts of BSC sources are bright enough, in most cases, to be accessible to optical spectroscopy from modest-sized (\\mbox{4-m} class) facilities, which will allow extensive ground-based exploration of the physical properties of the sources.  In the sections that follow, we make the case for a complete survey of BSC sources with \\swift, the \\swiftbscs.  With a \\mbox{500-s} pointing on each source, sources near the brightness threshold of the BSC will be re-detected with 25 counts in the \\swift\\ XRT, yielding a $<$5\\arcsec\\ position from the X-rays alone.  The simultaneous 500-s UVOT exposure will have a detection limit of $B=23.5$ in the $B$-filter, with comparable limits in other UV/optical filters when they are used.  In most cases this combination of observations will enable the identification of a unique UV/optical counterpart for the source, allowing flux estimates or upper limits for the source to be derived from archival surveys spanning the electromagnetic spectrum. In all, the required 9.4~Msec of observations can be accomplished with 10\\% of the time on-orbit over a three-year portion of the \\swift\\ mission.  It will thus be possible to schedule \\swiftbscs\\ observations so as to advance the secondary goals of the mission: Daily monitoring of the full sky to fluxes of $\\simgt$20\\,mCrab (10--100\\,keV) with the BAT, and a two-year all-sky BAT survey to $\\simgt$1\\,mCrab. ",
        "conclusions": "\\label{sec:conclusions}  We have investigated the prospects for a complete survey of the 18,811 sources of the \\rosat\\ All-Sky Survey Bright Source Catalog (BSC) with the \\swift\\ gamma-ray burst satellite, the ``\\swiftbscs,'' and find them to be good.  With 10\\% of the time on-orbit, over a three-year period, each BSC source can be observed for 500~s with the satellite's narrow-field X-ray (XRT) and UV/Optical (UVOT) telescopes.  The XRT observation will provide an improved $<$5-arcsec position for every BSC source that has not faded by more than a factor of 5 from its \\rosat-era flux level, and in addition, provide a 0.2--10 keV flux for the source.  The localization will represent an improvement by a factor of 10 to 30 in two-dimensional precision over the BSC positions (Fig.~\\ref{fig:bsc-gray}).  The simultaneous UVOT observation, taken in one of six broadband filters in the 200~nm to 600~nm wavelength range, will reach a limiting magnitude of $B=23.5$~mag or its rough equivalent, yielding candidate UV/optical counterparts or deep limits in every case.  Given the brightness of the BSC sources, and the \\lxlopt\\ properties of known source populations, the improved positions of the \\swiftbscs\\ can be expected to yield counterpart identifications in the great majority of cases.  Indeed, \\swiftbscs\\ sources without plausible optical (\\usno\\ or UVOT) counterparts will be immediate candidate X-ray binaries or isolated neutron stars (INSs; Fig~\\ref{fig:bsc-hist}).  Deriving counterpart identifications for the great majority of BSC sources will realize the promise of the BSC as a multiwavelength catalog, and is the best means to enumerate the rare source populations in the BSC and discover any new populations that may exist.  Without improved positions these efforts will continue to be frustrated by confusion problems -- with multiple \\usno\\ sources in the typical 60\\arcsec-diameter 2-sigma error circle, establishing a unique identification is nontrivial \\citep{rbp+00,rfb+03} -- and affected by unknown levels of contamination -- most efforts have associated each BSC source with the most promising candidate in its error circle, without deriving a quantitative estimate of the likelihood of association.  Each counterpart identification, then, will enable flux estimates or upper limits across the electromagnetic spectrum via current optical (DSS, DPOSS, SDSS), near-infrared (2MASS), and radio (NVSS, FIRST) surveys.  Any interesting BSC source that can be identified in this manner will be readily accessible to high signal-to-noise, high-resolution X-ray spectroscopy by \\xmmlong\\ and \\chandra.  In addition, those sources with optical counterparts will for the most part be accessible to optical-spectroscopic studies from modest-sized (\\mbox{4-m} class) ground-based facilities.  We have investigated \\swift\\ mission constraints on the \\swiftbscs\\ and do not find them to be severe.  The BSC targets will be of lower priority than new GRBs or active GRB afterglows, so that the survey will not conflict with the primary \\swift\\ goal of GRB observations. Moreover, the \\swiftbscs\\ observations can be scheduled so as to advance the secondary science goals of a daily all-sky hard X-ray (10--100 keV) survey to $\\simgt$20\\,mCrab, and a two-year all-sky survey to $\\simgt$1\\,mCrab, executed with the wide-field Burst Alert Telescope (BAT).  The 10\\% commitment of satellite time corresponds to roughly one BSC target per 95-min orbit.  Each such orbit will typically involve observations at 4 to 6 distinct pointings, some 3 to 5 of which will be active GRB afterglows.  Execution of the \\swiftbscs\\ will surely represent an operational challenge for the mission.  The progress of the survey will have to be monitored on an orbit by orbit basis, with new targets chosen from the unobserved portions of the BSC, and target lists updated as necessary when new GRB triggers result in the automated rescheduling of satellite observations.  Moreover, some level of global optimization of target choices will be desired, to retain isotropy of the remaining targets, as much as possible, until the very end of the survey.  However, the resulting capability for routine observations of large numbers of distinct targets will prove valuable should the \\swift\\ team choose to pursue complete surveys of other target catalogs.  In particular, the positions of sources identified in the two-year BAT all-sky survey will only be known to few-arcmin precision, so that a uniform campaign of follow-up observations with the XRT might be worth considering.  Complete XRT+UVOT surveys of other target populations -- for example, targets from the forthcoming \\textit{GALEX} all-sky imaging survey -- might also be contemplated.  Execution of these or similar projects will require a capability like that needed for the \\swiftbscs.  Ultimately, the choice to pursue the \\swiftbscs\\ is a choice to pursue the science we have described in preference to other non-GRB observations that \\swift\\ might carry out.  Beyond the first 1.5~years of on-orbit operations, the \\swift\\ Guest Investigator program will enter Cycle~2, and may solicit proposals from the community for pointed XRT and UVOT observations of designated non-GRB targets. Continued execution of the \\swiftbscs\\ into this era would imply that less time would be available for these programs, for two or three cycles, until the \\swiftbscs\\ is completed.  In a generic sense, the choice to make short observations of a large number of targets, as we propose here, may be contrasted with the alternative of making deeper exposures of some smaller set of targets, for example, monitoring X-ray binaries, soft gamma-ray repeaters, anomalous X-ray pulsars, or selected X-ray bright AGN at regular intervals.  The variety of science which might be accomplished with such campaigns is surely rich, and beyond our capability to summarize or even hint at in this context.  However, we believe that the science gains of a complete \\swiftbscs\\ are of sufficient value to warrant at least a similar level of consideration.  Most importantly, the science of the \\swiftbscs\\ is science that only \\swift\\ can achieve.  \\swift's combination of multiwavelength instrumentation, fast-slew capability, and flexibility in target scheduling make it the only present or planned satellite mission capable of executing a complete BSC survey without significant operational overheads.  The \\swiftbscs\\ thus represents the best opportunity for many years to come to expand the number of securely-identified X-ray sources by an order of magnitude, from $\\sim$2000 to $\\sim$18,000.  Declining to take advantage of this opportunity will restrict the observations of the current and next-generation X-ray missions to more or less the known population of X-ray sources that exists today.  On the other hand, prompt execution of the \\swiftbscs\\ will allow immediate follow-up studies with \\xmm\\ and \\chandra\\ of the most intriguing new members of all X-ray source populations, from stars to compact objects to AGN.  This will renew the promise of these observatories, for all areas of X-ray astronomy, as they enter their extended mission phase."
    },
    "0403/astro-ph0403507_arXiv.txt": {
        "abstract": "{We analyze the consequences of the hypothesis that the formation of the Broad Line Region is intrinsically connected to the existence of the cold accretion disk. We assume that the Broad Line Region radius is reliably estimated by the formula of Kaspi et al. (2000). We consider three models of the disappearance of the inner disk that limit the existence of the Broad Line Region: (i) the classical ADAF approach, i.e. the inner hot flow develops whenever it can exist (ii) the disk evaporation model of Meyer \\& Meyer-Hofmeister (2002) (iii) the generalized disk evaporation model of R\\' o\\. za\\' nska \\& Czerny (2000b). For each of the models, we determine the minimum value of the Eddington ratio and the maximum value of the broad line widths as functions of the viscosity parameter $\\alpha$ and the magnetic field parameter $\\beta$. We compare the predicted parameter space with observations of several AGN. Weak dependence of the maximum value of the FWHM and minimum value of the Eddington ratio on the black hole mass in our sample is noticeable. It seems to favor the description of the cold disk/hot inner flow transition as in the classical ADAF approach rather than with the model of disk evaporation due to conduction between the disk and accreting corona. ",
        "introduction": "One of the most important issues in studying accretion onto compact objects is to understand the process that causes some fraction of the accreting material to be heated sufficiently strongly to form a hot optically thin, predominantly thermal plasma in the vicinity of a black hole (see e.g. Narayan, Mahadevan \\& Quataert 1998;  Collin et al. 2001;  Czerny 2003).  Observations of the X-ray emission from this plasma do not uniquely constrain the geometry of the plasma. Several phenomenological models are under consideration;  those include magnetic flares above an accretion disk extending always down to the marginally stable orbit (Galeev, Rosner \\& Vaiana 1979; Haardt, Maraschi \\& Ghisellini 1994), or the solutions of accretion disk structure with inner optically thin (possible advection-dominated) flow (Ichimaru 1977; Narayan \\& Yi 1994, Abramowicz et al. 1995). The geometry possibly depends on the luminosity to the Eddington luminosity ratio in a given object. Detailed XMM-Newton and Chandra studies of the X-ray reflection (both continuum and iron $K_{\\alpha}$ line) suggest that at high Eddington ratios the disk extends to the marginally stable orbit and its surface appears strongly ionized (e.g. \\Agata et al. 2004 for Ton S180; Janiuk et al. 2001 and  Pounds et al. 2003 for PG1211+143; Sako et al. 2003 for MCG -6-30-15 and Markarian 766) while in low Eddington ratio sources the spectra are well-described as being due to cold reflection from the outer part of the accretion disk (e.g. O'Brien et al. 2003).  Those observations motivate the development of models in which the material slowly evaporates from the cold disk into a hot corona as the black hole is approached (Liu et al. 1999, 2002, \\Agata \\& Czerny 2000b, Meyer \\& Meyer-Hofmeister 2002, Meyer-Hofmeister \\& Meyer 2003). All such models predict that at high Eddington ratios the evaporation effect is relatively inefficient so the cold optically thick disk survives all the way down to the black hole while for low Eddington ratio sources the evaporation leads to a complete evacuation of the disk and formation of an inner optically thin, hot flow.  However, the physical description of the evaporation rate involves a number of complicated processes, and various models rely on several underlying assumptions so it is very important to verify the models against observational constraints. In this paper we use the constraint based on the measured width of the broad emission lines. The exact nature of the broad line region (BLR) is not known but most models connect the formation of low ionization lines like $H_{\\beta}$ with the wind developing close to the disk surface (Collin-Souffrin et al. 1988; Dultzin-Hacyan et al. 2000, Laor 2003a). If such an approach is adopted, the existence of the broad lines at a certain range of radii is possible only if the cold disk extends also at least to the same radial range, and regions of the efficient line/wind formation are related to specific phenomena taking place within the disk. For example, such an approach led Elvis (2000) to propose a general, consistent scheme for quasar structure.  The details of the processes, including disk/BLR coupling, are by no means clear but the large amount of the existing data opens possibilities to test various hypotheses of the nature of this relation. Laor (2003b) noticed that there appears to be an upper limit to the broad line width of about 25 000 km/s which apparently sets a lower limit for the luminosity of the broad line object. He postulated that a significant population of AGN exists that intrinsically do not have broad lines. Further developing this idea, Nicastro et al. (2003) found the division between the AGN with and without broad lines occurs at the Eddington ratio of $\\sim 0.001$, and interpreted this result as an argument for a connection between the onset of a wind from the disk surface and the transition from a gas dominated to a radiation pressure dominated part of the cold disk.  In this paper we assume that the outer cold disk transitions to an inner hot, optically thin and radiatively inefficient flow.  We further assume that the BLR may form above a cold disk, in its atmosphere or outflowing wind, but not above a hot inner flow. These assumptions allow us to qualitatively test the theoretical predictions for the location of such a transition. We consider three specific models: the classical condition for the existence of ADAF (Abramowicz et al. 1995), the evaporation model of Meyer \\& Meyer-Hofmeister (2002) and the evaporation model of \\Agata and Czerny (2000) generalized for the case of a non-negligible magnetic field. ",
        "conclusions": "Tran (2001, 2003) and Laor (2003b) argue that in a fraction of AGN, broad lines are not detected - but not because they are hidden from our line of sight, but because those objects do not develop a BLR.  Recent X-ray observation of NGC 6251 supports this view (Gliozzi et al. 2004).  Our study shows that the existence of this class of objects is a natural consequence of the evaporation of the cold disk close to a black hole and the transition of the accretion flow into a hot optically thin phase, if the BLR formation is intimately connected with the cold disk. The transitions happens far away for low Eddington ratio objects, at distances where the ionization conditions would be correct for broad line formation, and the BLR apparently lose the needed disk 'support'.  Quantitatively, the allowed parameter space (in log $L/L_{Edd}$ vs. log $FWHM$ and  log $M$ vs. log $FWHM$; see Figs.~\\ref{fig:data1} and \\ref{fig:data2}) for broad line objects in our sample seems to be better represented by the {\\it strong ADAF principle} than by the more physical disk evaporation model."
    },
    "0403/astro-ph0403677_arXiv.txt": {
        "abstract": "We study the peak energy ($E_{\\rm{p}}$) distribution of the $\\nu F_{\\nu}$ spectra of gamma-ray bursts (GRBs) and X-ray flashes (XRFs) with a sample of 57 bursts observed by {\\em High Energy Transient Explorer 2} ({\\em HETE-2}) French Gamma Telescope and discuss its implications for the jet structure models. Combining the observed $E_{\\rm{p}}$ distribution of {\\em HETE--2} GRBs/XRFs with that of BATSE GRBs, we find that the observed $E_{\\rm{p}}$ distribution of GRBs/XRFs is a bimodal one with peaks of $\\lesssim 30$ keV and $\\sim 160-250$ keV. According to the recently-discovered equivalent-isotropic energy-$E_{\\rm{p}}$ relationship, such a bimodal distribution implies a two-component structure of GRB/XRF jets. A simple simulation analysis shows that this structured jet model does roughly reproduce a bimodal distribution with peaks of $\\sim 15$ and $\\sim 200$ keV. We argue that future observations of  the peak of $\\sim 15$ keV in the $E_{\\rm{p}}$ distribution would be evidence supporting this model. {\\em Swift}, which covers an energy band of 0.2--150 keV, is expected to provide a key test for our results. ",
        "introduction": "% X-ray flashes (XRFs) have being gotten a lot of attention in the last 2 years (Heise et al. 2001; Kippen et al. 2003). They are thought to be a lower energy extension of the known gamma-ray burst (GRB) population, based on the fact that their spectral behaviors are similar to those of GRBs (Kippen et al. 2003; Barraud et al. 2003; Sakamoto et al. 2004; Lamb et al. 2003a, b, 2004). The nature of a narrow cluster of the observed $E_{\\rm{p}}$ distribution of BATSE GRBs remains poorly understood, which might be related to the jet structure of GRBs. XRFs broaden the energy coverage of prompt GRB emission and may bring more signatures of the jet structure of GRBs (Lamb et al. 2003a, b, 2004).  The jet structure models are currently under heavy debate. Any model should present a unified description for GRBs and XRFs. Two currently competing models are the structured jet model (M\\'{e}sz\\'{a}ros, Rees \\& Wijers 1998; Dai \\& Gou 2001; Rossi, Lazzati, \\& Rees 2002; Zhang \\& M\\'{e}sz\\'{a}ros 2002; Granot \\& Kumar 2003; Kumar and Granot 2003; Panaitescu \\& Kumar 2003; Wei \\& Jin 2003) and the uniform model (e.g., Rhoads 1999; Frail et al. 2001). Zhang et al. (2004) show that the current GRB/XRF prompt emission/afterglow data can be described by a quasi--Gaussian-type (or similar structure) structured jet with a typical opening angle of $\\sim 6^{o}$ and with a standard jet energy of $\\sim 10^{51}$ ergs. Alternatively, based on the {\\em High Energy Transient Explorer 2} ({\\em HETE-2}) observations, Lamb et al. (2003a, b, 2004) propose that the uniform jet model can reasonably describe the unified scheme of GRBs/XRFs. Very recently, the two-component jet model was advocated by Berger et al. (2003) based on observations of GRB 030329, which has two different jet breaks in an early optical afterglow light curve (0.55 days, Price et al. 2003) and in a late radio light curve (9.8 days). Millimeter observations of this burst further support the two--component jet model (Sheth et al. 2003). Numerical calculations of such a model were performed by Huang et al. (2004 ). This model suggests that a GRB/XRF jet has two components: a narrow, highly relativistic one, and a wide, mildly relativistic one. When the line of sight of an observer is within the narrow component, the observed burst is a typical GRB, but when the line of sight is pointing to the wide--component, it is an XRF.  A broad spectral energy distribution could constrain the jet structure models. A low peak energy of the $\\nu F_{\\nu}$ spectrum ($E_{\\rm{p}}<50$ keV) and weak gamma--ray fluxes ($F<0.2$ photons cm$^{-2}$ s$^{-1}$, 50--300 keV energy range) distinguish XRFs from typical GRBs (Kippen et al. 2003; Mochkovitch et al. 2003). It is well known that the observed $E_{\\rm{p}}$ distribution of BATSE GRBs is narrowly clustered. Does the observed $E_{\\rm{p}}$ distribution of XRFs exhibit a similar feature? In this Letter, we focus on this question. We analyze the observed $E_{\\rm{p}}$ distribution with a sample of 57 bursts observed by {\\em HETE-2}/ French Gamma Telescope (FREGATE). Combining the observed $E_{\\rm{p}}$ distribution of {\\em HETE-2} GRBs/XRFs with that of BATSE GRBs, we find that the observed $E_{\\rm{p}}$ distribution of GRBs/XRFs is a bimodal one peaking at $\\lesssim 30$ keV and $\\sim 160-250$ keV. With respect to this result, we suggest that the two--component jet model is a reasonable candidate model for GRB/XRF jets. A simulation analysis confirms this suggestion. ",
        "conclusions": "We have studied the observed $E_{\\rm{p}}$ distribution of 57 {\\em HETE-2}/FREGATE bursts, and discuss its implications for the jet structure models. Combining the observed $E_{\\rm{p}}$ distribution of {\\em HETE-2} GRBs/XRFs with that of BATSE GRBs, we suggest that the observed $E_{\\rm{p}}$ distribution of GRBs/XRFs is bimodal with peaks of $\\lesssim 30$ keV and $\\sim 160-250$ keV. According to the recently--discovered equivalent-isotropic energy--$E_{\\rm{p}}$ relationship, we find that the bimodal distribution can be explained by the two--component model of GRB/XRF jets. A simple simulation analysis shows that this structured jet model does roughly reproduce the bimodal distribution with peaks of $\\sim 15$ keV and $\\sim 200$ keV.  The peak of $\\sim 15$ keV in the simulated $E_{\\rm{p}}$ distribution is key evidence for the two-- component jet model. It is near the lowest end of the energy bandpass of {\\em HETE-2}/FREGATE. Fortunately, {\\em HETE-2} provides a weak clue to this peak. A more sensitive instrument than {\\em HETE--2} with an energy bandpass $1-50$ keV is required to further confirm this peak. {\\em Swift}, which covers an energy band of 0.2--150 keV (we mark this region in Figure 2 with diagonal lines)\\footnote{ http://swift.gsfc.nasa.gov/science/instruments/} is expected to provide a key test for it.  Simulations of the propagation and eruption of relativistic jets in massive Wolf-Rayet stars by Zhang, Woosley, \\& Heger (2004b) show that an erupting jet has a highly relativistic, strongly collimated core, and a moderately relativistic, less energetic cocoon. The cocoon expands and becomes visible at larger angles. The energy ratio of the cocoon to the core in their simulation is about one order. From our simulation results, we find that it is $\\sim (E_{\\rm{p,GRB}}/E_{\\rm{p,XRF}})^2(\\theta_{1}/\\theta_2)^2\\sim 40$, being roughly consistent with their results. Their simulations seem to support the two-component jet model. We have noted that the ability of the cocoon to cause an XRF depends sensitively on its Lorentz factor, which is determined by the degree of mixing between the jet and envelope material. Matzner (2003) argued that this mixing might be difficult to resolve in numerical simulations.  A two-component jet was suggested to be universal for GRB/XRF phenomena in this Letter, based on the multi-wavelength observations of GRB 0303029 (Berger et al. 2003; Sheth et al. 2003) and the bimodal distribution of $E_{\\rm {p}}$. It should be pointed out that other jet models such as uniform jets and single-component-universal jets were proposed to explain numerous observations on the afterglows and some correlations (e.g., Lamb et al. 2003b; Lloyd-Ronning et al. 2004; Lloyd-Ronning \\& Zhang 2004). Thus, one would expect strong evidence showing which jet model is more reasonable.  We are very grateful to the anonymous referee for his/her valuable suggestions. We also sincerely thank Bing Zhang and Johan P.U. Fynbo for their helpful comments. These suggestions and comments have enabled us to improve  the manuscript greatly. This work is supported by the National Natural Science Foundation of China (grants 10233010 and 10221001), the National 973 Project (NKBRSF G19990754), the Natural Science Foundation of Yunnan (2001A0025Q), and the Research Foundation of Guangxi University."
    },
    "0403/astro-ph0403394_arXiv.txt": {
        "abstract": "I review the basic phenomenology and theory of the millisecond brightness oscillations observed during thermonuclear X-ray bursts from 13 of $\\approx 70$ accreting neutron stars in low-mass X-ray binaries. Compelling observations indicate that the oscillations are produced by surface brightness patterns on the rapidly rotating neutron stars. However, it remains to be understood (1) why the brightness patterns producing them persist for up to 15 s during an X-ray burst, whereas the burning should cover the entire surface in less than 1 s, and (2) why the frequencies drift upward by $\\approx 5$ Hz during the course of the burst. These peculiarities can probably be explained by taking into account the expansion of the surface layers caused by the burning, zonal flows that form due to pressure gradients between the equator and poles, and Rossby-Alfv\\'{e}n modes that are excited in the surface ocean. Further progress toward understanding how burning progresses on the surface of the neutron star can be made with a next-generation X-ray timing mission, which would provide a larger sample of sources with oscillations, detect sideband signals produced by the spectrum of modes that should be excited in the neutron star ocean, and measure harmonic structure in the profiles of the oscillations. These observations would be crucial for measuring the distribution of the rotation rates of neutron stars, the progression of unstable nuclear burning in the accreted ocean, and the curvature of the space-time around the neutron star. ",
        "introduction": "Thermonuclear X-ray bursts are produced when helium on the surface of an accreting neutron star ignites unstably \\citep[see][for a review]{sb03}. The accreted material itself is usually hydrogen. The helium is produced because the accreted hydrogen is compressed and heated by the column of material above it, and begins to burn steadily via the CNO cycle. Eventually, the critical temperature and density are reached such that He can burn via a triple-$\\alpha$ process. The burning is unstable, and engulfs the entire surface of the neutron star in less than a second. This results in a $10^{38}$ erg s$^{-1}$ flash of X-rays that out-shines the emission from the accretion flow for tens of seconds. These bursts recur on time scales of hours to days, and so multiple bursts have been observed from the neutron stars in $\\approx 65$ low-mass X-ray binaries.  The unstable burning is likely to begin in a small region on the neutron star, and so it has long been expected that the resulting hot spot should produce modulations in the burst flux at the spin period of the neutron star. Indeed, ``burst oscillations'' have now been observed with {\\it RXTE} from 13 neutron star LMXBs (Table~\\ref{tab:init}). There are several reasons to believe that the burst oscillations occur at the spin frequencies of the neutron stars. First, and most importantly, two of these sources are X-ray pulsars, which exhibit periodic modulations in the persistent emission between bursts at the same frequency as the burst oscillations \\citep{cha03,str03}. The frequencies of all of the oscillations are characteristic to each source, and are distributed uniformly between 270 and 620 Hz. Second, once one accounts for a small frequency drift (described below) the burst oscillations are nearly coherent \\citep{str96,sm99}. In one case, the oscillations are observed to be coherent for $10^5$ cycles during a carbon superburst \\citep{sm02}. Third, the maximum frequencies of the oscillations are stable to within a few parts in a thousand in bursts separated by several years \\citep{str98, sm99, gil02, mun02a}. Fourth, the oscillations are strongest in the rises of bursts, when the nuclear burning is likely to be confined to small areas on the surfaces of the neutron stars \\citep{str98}. Finally, in the tails of the bursts, the amplitudes of the oscillations as a function of energy are consistent with those expected from temperature variations of $\\approx 0.2$ keV across the surface of the neutron star \\citep{moc03}. As signals from the surfaces of neutron stars, these oscillations can be used to study the evolution of the spin frequencies of accreting neutron stars, the spacetime around the star, and how thermonuclear burning proceeds on the stellar surface. \\begin{figure}[thb] \\resizebox{.7\\textwidth}{!} {\\includegraphics{munom_f1.ps}} \\caption{Dynamic power spectra illustrating the frequency evolution of burst oscillations. Contours of power as a function of frequency and time were generated from power spectra of 2 s intervals computed every 0.25 s. A Welch function was used to taper the data to reduce sidebands in the power spectrum due to its finite length. The contour levels are at powers of $0.02$ in single-trial probability starting at a chance occurrence of $0.02$. The PCA count rate is plotted referenced to the right axis. \\label{fig:freqev}} \\end{figure}  However, the simplest models of inhomogeneous burning fail to explain two aspects of the oscillations, which are illustrated in Figure~\\ref{fig:freqev}. First, the oscillations persist for up to 15 s during a burst, long after the burning should have engulfed the entire surface of the neutron star. Second, they drift upward in frequency by up to 5 Hz during the course of a burst, which suggests that the brightness pattern moves opposite the sense of the rotation, such that $\\omega_{\\rm obs} = \\Omega_{\\rm NS} - \\omega_{\\rm pattern}(t)$ \\citep{str97a}. Several models have been proposed to explain one or both of these aspects of the oscillations. ",
        "conclusions": ""
    },
    "0403/astro-ph0403657.txt": {
        "abstract": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% A new classification of neutron star cooling scenarios, involving either ``minimal'' cooling or ``enhanced'' cooling is proposed.  The minimal cooling scenario replaces and extends the so-called standard cooling scenario to include neutrino emission from the Cooper pair breaking and formation process.  This emission dominates that due to the modified Urca process for temperatures close to the critical temperature for superfluid pairing.  Minimal cooling is distinguished from enhanced cooling by the absence of neutrino emission from any direct Urca process, due either to nucleons or to exotica such as hyperons, Bose condensates or deconfined quarks.  Within the minimal cooling scenario, theoretical cooling models can be considered to be a four parameter family involving the equation of state (including various compositional possibilities) of dense matter, superfluid properties of dense matter, the composition of the neutron star envelope, and the mass of the neutron star.  The consequences of minimal cooling are explored through extensive variations of these parameters.  The results are compared with the inferred properties of thermally-emitting neutron stars in order to ascertain if enhanced cooling occurs in any of them.  All stars for which thermal emissions have been clearly detected are at least marginally consistent with the lack of enhanced cooling, given the combined uncertainties in ages and temperatures or luminosities.  The two pulsars PSR 0833-45 (Vela) and PSR 1706-44 would require enhanced cooling in case their ages and/or temperatures are on the lower side of their estimated values whereas the four stars PSR 0656+14, PSR 1055-52, Geminga, and RX J0720.4-3125 may require some source of internal heating in case their age and/or luminosity are on the upper side of their estimated values.  The new upper limits on the thermal luminosity of PSR J0205+6449 (in the supernova remnant 3C58) and RX J0007.0+7302 (in CTA 1) are indicative of the occurrence of some enhanced neutrino emission beyond the minimal scenario.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "introduction": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  Within the last several years, several candidates for thermally-emitting neutron stars have been discovered (see, e.g., \\cite{PZ02} for a short review).  These stars are presumably cooling through the combination of neutrino emission from the interior and photon cooling from the surface, the latter responsible for their observed thermal emissions.  Their temperatures have been deduced by fitting atmosphere models to their spectra whereas ages can be inferred from kinematics or from the associated pulsar spin-down timescale. Neutron star cooling depends upon the equation of state (EOS) of dense matter as well as the neutron star mass and their envelope composition.  It has been hoped that comparing theoretical cooling curves, i.e., the temperature-age or luminosity-age relation for neutron stars, with observations could yield information about their internal properties.  A goal of this paper to explore how observations of thermal emission from neutron stars might be able to constrain the equation of state of dense matter. We will treat theoretical neutron star cooling trajectories as a four parameter series of models. The parameters are: \\\\  \\noindent (1) the equation of state (including various compositional possibilities), \\\\ % \\noindent (2) superfluid properties of the relevant components, \\\\ % \\noindent (3) the envelope composition, and \\\\ % (4) \\noindent the stellar mass.  \\\\  In nature, only one equation of state and one set of superfluid properties is realized, but at the present time, the theoretical range of superfluid properties for a given equation of state is so broad that these must be treated as an independent parameter.  In the future, as more observations become available, it should be possible to eliminate some combinations of EOS and superfluid parameter sets, if not entire families of possibilities.  It is also not known if a neutron star's envelope composition is unique (at least for a given mass) or if it varies from star to star or as a function of time.  The neutron star mass is constrained to lie between the maximum mass (a function of the equation of state) and a minimum mass of about 1.1 M$_\\odot$ (set by theoretical considerations of neutron star birth (see, e.g., \\cite{BL86}).   Historically, theoretical neutron star cooling models have fallen into two categories, ``standard'' cooling or enhanced cooling.  The so-called ``standard cooling'' scenario has no ``enhanced cooling'' which could result from any of the direct Urca processes involving nucleons, hyperons, meson condensates or quark matter (see, e.g. \\cite{P92} and \\cite{Pr98}).  Until recently, ``standard'' cooling has been treated as being dominated by the modified Urca process \\citep{FM79}.  However, in the presence of superconductivity or superfluidity in the neutron star interior, an additional source of neutrino emission, Cooper pair breaking and formation, occurs  \\citep{FRS76,VS87}. For temperatures near the associated gap energies, Cooper pairs, in fact, dominate the neutrino emissivities.  Although the magnitude of the superfluid gap energies as a function of density is somewhat uncertain at present, it is generally accepted that superfluidity occurs in neutron star matter.  For this reason, we embark on the ``minimal cooling'' scenario in which ``standard cooling'' is extended to include the effects of superfluidity, including Cooper pair breaking and formation.  The purpose of this paper is to explore in as complete a fashion as possible the consequences of this minimal cooling paradigm, employing the four kinds of parameters described above, and to compare our results with the inferred properties of cooling neutron stars.  In this way, it will become apparent to what extent one or more of the so-called enhanced cooling mechanisms might be necessary to understand the observations.  A future paper will explore in a similar fashion the consequences of enhanced cooling.  One consequence of the minimal cooling paradigm, that enhanced cooling will not occur, is a restriction upon the equation of state involving the symmetry energy.  It is well-known (see, e.g., \\cite{LPPH91}) that the density dependence of the nuclear symmetry energy controls the charge fraction in uniform beta-equilibrium matter.  Since the direct Urca process occurs in uniform beta-equilibrium matter when the charge fraction exceeds 1/9 (in the absence of muons or hyperons), minimal cooling thus restricts the density dependence of the nuclear symmetry energy: the critical density for the onset of the direct Urca process must remain above the star's central density.  When muons are considered, this critical density, for a given equation of state, is slightly lowered.  The appearance of hyperons can also trigger other direct Urca processes~\\citep{PPLP92}.  Thus, to the extent that minimal cooling can explain existing observations, a constraint on the equation of state could be inferred.  In \\S~\\ref{Sec:Data}, observations of cooling neutron stars are reviewed.  The input physics, including the equation of state, superfluid properties, and neutrino emissivities are discussed in \\S~\\ref{Sec:Physics}.  The influence of the neutron star envelope is briefly discussed in \\S~\\ref{Sec:Envelope}.  The results of cooling calculations for minimal cooling models are extensively discussed in \\S~\\ref{Sec:General}. The coldest stars possible within the minimal cooling scenario are identified in \\S~\\ref{Sec:Try-the best} and \\S~\\ref{Sec:Minimal-Data} contains a summary of the confrontation of the minimal cooling paradigm with existing data.  A comparison with other studies is performed in \\S~\\ref{Sec:Other}.  Conclusions are offered in \\S~\\ref{Sec:Conclusion}. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{Sec:Conclusion} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  We have presented a detailed study of the thermal evolution of an isolated neutron star using what we term as the minimal cooling scenario.  This scenario is an extension of the well-known ``standard cooling'' scenario to include the effects of nucleon pairing and complements neutrino emission by the modified Urca and nucleon bremsstrahlung processes with the pair breaking and formation (PBF) process.  We have confirmed the results of previous works by others that for many models of nucleon pairing, the PBF process actually dominates the cooling of the star (see \\S~\\ref{Sec:Coop}) and hence is an essential ingredient of the minimal cooling scenario.  Among the four parameters we proposed for an overall classification of neutron star cooling models, we found that the EOS at supranuclear densities is well constrained by the requirements of the minimal cooling scenario. Moreover, we showed that the stellar mass has little effect on the results.  We emphasize that for scenarios beyond the minimal one, i.e., when new particles and neutrino emission processes appear, these two parameters definitely gain importance.  The other two parameters we considered, pairing properties of the nucleons and chemical composition of the envelope, introduce the largest uncertainties in our theoretical predictions.  We singled out three subclasses of scenarios due to uncertainties in the size and extent of the $n$ $^3P_2$ gap.  For this gap, we considered three different cases: a vanishingly small gap, a somewhat small gap (our model ``a'') and a relatively large gap (model ``b'').  Within each of these cases, variations of the $n$ and $p$ $^1S_0$ gaps covering the published ranges of these gaps were also considered.  With respect to the chemical composition of the envelope, we singled out two extreme cases: an envelope consisting of heavy iron-like elements and an envelope containing essentially only light elements. For each choice of the $n$ $^3P_2$ gap, we obtained two families of closely packed cooling curves representing each of the two extreme envelope cases, with some spreads due to variations of the $n$ and $p~^1S_0$ gaps.  Intermediate envelope chemical compositions, or its possible temporal evolution, result in cooling trajectories intermediate between the extremities (see Figures~\\ref{Fig:Cooling-eta} and \\ref{Fig:Decay-Env}).  Comparing our results with observationally inferred temperatures $T_\\infty$ and luminosities $L_\\infty$ of eleven isolated neutron stars (Figures~\\ref{Fig:plot_cool_lum_data} and \\ref{Fig:plot_cool_data_11}), we found that the observations were in overall good agreement with the minimal cooling scenario taking into account the uncertainties of envelope and pairing properties as well as those of the ages and inferred $T_\\infty$ and $L_\\infty$ for these stars.  It is probably not possible to understand these data within a single model with a unique envelope chemical composition.  Considering that the compositions of the upper layers are strongly dependent on poorly understood processes which occurred during and soon after the birth of the star, and  possibly later due to the dynamics, it appears likely that different stars have envelopes with different chemical compositions.  The Vela pulsar 0833-45, and possibly, but with much larger uncertainties, PSR 1706-44, are marginal candidates for enhanced cooling as their inferred temperatures and luminosities are lower than most of our models.  Nevertheless, we found that low mass neutron stars, $\\sim 1.1 - 1.2$ \\Msun, with a light element envelope and extensive nucleon pairing covering essentially the entire core, could reach the inferred values, $T_\\infty \\simeq 10^{5.8}$ K at an age of $\\sim 20,000$ years (Figure~\\ref{Fig:plot_cool_data_11}) which is within the range of estimated age of the associated supernova remnant.  An essential component of the minimal cooling scenario is neutrino emission by the PBF process, which leads, in the presence of appropriate nucleon gaps, to more rapid cooling than possible in the standard cooling scenario.  The low observed temperatures of the two pulsars, Vela and 1706-44, can be accomodated by the PBF process with a $n$ $^3P_2$ gap of sufficient size.  However, if this gap were vanishingly small \\citep{SF04}, then the temperature and thermal luminosity measurements of these two objects would be evidence for the presence of processes beyond the minimal cooling paradigm.  The five older stars, PSR 0656+14, PSR 1055-52, PSR 0633+1748 (Geminga), and RX J0720.4-3125, unfortunately have such large uncertainties on both their ages and thermal luminosities that their interpretation is delicate.  They do not require the occurence of enhanced neutrino emission and can be accomodated within the minimal cooling scenario when the lowest values for $T$ and $L_\\infty$ are chosen.  But, should the upper limits for $T$ and $L_\\infty$ prevail, they would be good candidates for the occurrence of some ``heating mechanism'', i.e., dissipative processes, which inject heat into the star (see, e.g., \\cite{USNT93} \\& \\cite{SSWW99}).  In the case of RX J1856.5-3754, which has a much more tightly constrained age and luminosity, the agreement with the minimal scenario is excellent.   The two objects standing apart from the other observed neutron stars are J0205+6449 (in the supernova remnant 3C58) and RX J0007.0+7302 (in CTA 1). Upper limits on their luminosities are well below the predicted values for any of our models.  These two objects are the best candidates, to date, for the necessity to go beyond the minimal scenario.    Finally, the four upper luminosity limits on the undetected objects in shell supernova remnants (G084.2-0.8, G093.3+6.9, G127.1+0.5, and G315.4-2.3) recently found by \\cite{Ketal2004} are so low that they will constitute the strongest evidence for enhanced neutrino emission well beyond the minimal scenario if it can be demonstrated that they actually correspond to neutron stars and not quiescent black-holes."
    },
    "0403/astro-ph0403441_arXiv.txt": {
        "abstract": "{ We investigate the influence of different magnetic field configurations on the temperature distribution in neutron star crusts. We consider axisymmetric dipolar fields which are either restricted to the stellar crust, ``crustal fields'', or allowed to penetrate the core, ``core fields''. By integrating the two-dimensional heat transport equation in the crust, taking into account the classical (Larmor) anisotropy of the heat conductivity, we obtain the crustal temperature distribution, assuming an isothermal core. Including quantum magnetic field effects in the envelope as a boundary condition, we deduce the corresponding surface temperature distributions. We find that core fields result in practically isothermal crusts unless the surface field strength is well above $10^{15}$ G while for crustal fields with surface strength above a few times $10^{12}$ G significant deviations from isothermality occur at core temperatures inferior or equal to $10^8$ K. At the stellar surface, the cold equatorial region produced by the quantum suppression of heat transport perpendicular to the field in the envelope, present for both core and crustal fields, is significantly extended by the classical suppression at higher densities in the case of crustal fields. This can result, for crustal fields, in two small warm polar regions which will have observational consequences: the neutron star has a small effective thermally emitting area and the X-ray pulse profiles are expected to have a distinctively different shape compared to the case of a neutron star with a core field. These features, when compared with X-ray data on thermal emission of young cooling neutron stars, will open a way to provide observational evidence in favor, or against, the two radically different configurations of crustal or core magnetic fields. ",
        "introduction": "%  The presence of strong magnetic fields in neutron stars is one of their distinctive characteristics. In typical neutron stars the observed and/or inferred surface fields are of the order of $10^{12...13}$ G ; for magnetars they even reach $10^{14...15}$ G. The strength, structure and time evolution of the magnetic field is intimately related to its origin, which is still an open problem. Basically, two qualitatively different types of field structures are conceivable, one having an initial field penetrating the whole star while the other is characterized by having the field and its supporting currents restricted to the stellar crust (see, e.g., Chanmugam \\cite{C92}). To date, there is still no compelling observational evidence in favor of or against either of these two hypotheses but recently Link \\cite{L03} argued that long period pulsar precession, as observed in PSR B1828-11 (Stairs et al. \\cite{SLS00}), may be impossible if the magnetic field penetrates regions of the core where neutrons are superfluid and proton superconducting.  Any observed magnetic field, either based on crustal or core currents, has to penetrate the crust matter, thereby affecting its transport properties. Moreover, the structure of the field in the shallow spherical shell layer below the surface has a direct  effect on the surface temperature and leads to a non-uniform distribution with observational consequences (Page \\cite{P95}) if its strength is above $10^{10}$ G. Roughly, the effects of the magnetic field onto the transport processes can be divided into classical and quantum ones (see, e.g., Yakovlev \\& Kaminker \\cite{YK94} for a review). The classical effects are due to the Larmor rotation of the electrons, the main carriers of charge and heat, and are determined by the magnetization parameter $\\omega_{\\scriptscriptstyle B} \\tau$ where $\\omega_{\\scriptscriptstyle B} \\equiv e B/m_e^* c$ is the gyrofrequency of the electrons, $\\tau$ being their relaxation time and $m_e^*$ their effective mass. Quantized motion of the electrons transverse to the magnetic field causes the quantum effects which are of importance only if few Landau levels are occupied, requiring thus densities $\\rho < \\rho_B = 7045 (B/10^{12} \\mathrm{G})^{3/2} \\frac{<A>}{<Z>}$ g cm$^{-3}$ and temperatures $T \\ll T_B \\approx 1.34 \\cdot 10^8 (B/10^{12} \\mathrm{G})$ K (Potekhin et al. \\cite{PYCG03}) where $<A>$ and $<Z>$ are the average mass and charge numbers of the ions. From these numbers it becomes clear that quantum effects will play an important role for  strong magnetic fields in the outermost layers of the neutron star crust - the thin low density shell of the envelope - but are negligible in the deeper layers. The thinness of this envelope justifies a study of heat transport in a plane parallel, one dimensional approximation and many such calculations have been performed (see, e.g., for the most recent one, Potekhin \\& Yakovlev \\cite{PY01}, hereafter PY01, and references therein). Two-dimensional calculations of heat transport with magnetic field have been presented by Schaaf (\\cite{S90a},\\cite{S90b}) who however restricted himself to the thin envelope and an uniform magnetic field. Tsuruta (\\cite{T98}) has presented results of two-dimensional cooling calculations of neutron stars which included the quantizing effect of a dipolar magnetic field in the envelope. These 2D calculations showed that the 1D approximation is indeed very good when the field affects heat transport only in the thin envelope.  The deeper layers of the crust can, however, be affected by classical effects in case of strong fields, i.e., large $\\omega_{\\scriptscriptstyle B}$, and/or low temperature, i.e., high $\\tau$. In such a case the usual assumption of an isothermal crust is questionable and this is the issue we want to address in this paper. In particular, in the case of a crustal magnetic field its strength in the crust unavoidably exceeds its surface value by one to two orders of magnitude (see, e.g., Page et al. \\cite{PGZ00}) \\ and its effect can naturally be expected to be stronger than in the case of a field permeating the whole star. This may open a way to study the structure of the magnetic field in the crust and provide observational features which will allow us to discriminate between the above mentioned two types of hypothesized field structure, crustal vs. core.  The paper is organized as follows: In the next Sect. ~\\ref{sec:input} the basic equations are introduced which describe the magnetic field and the heat transport influenced by it. The components of the heat conductivity tensor are given and the outer boundary condition is discussed. The physical input as well as the numerical method are shortly described. In Sect. \\ref{sec:results} the results of the  numerical calculations are presented. For different core temperatures, magnetic field strengths and geometries the crustal temperature profiles, the surface temperature distributions and the corresponding luminosities are calculated. Sect.~\\ref{sec:discon} is devoted to the discussion of the consequences of the magnetic field effects on to the crustal and surface temperature distribution. ",
        "conclusions": "\\label{sec:discon}  The insulating effect of a quantizing magnetic field component perpendicular to the radial heat flux is well known and its role in the low density layers of the envelope of a neutron star has been extensively studied (e.g., Hernquist \\cite{H85}; Schaaf \\cite{S90a}; Heyl \\& Hernquist \\cite{HH98}; YP01). Here, importance of the classical (Larmor) effect for the heat transport through the whole crust is demonstrated, provided the neutron star possesses a sufficiently strong magnetic field maintained by electric currents circulating in the crust but which does not penetrate the core of the star. Besides its insulating effect the tangential crustal magnetic field creates a meridional heat flux from the equatorial regions towards the high latitude ones. It transports the heat, which is dammed in the equatorial region, towards the poles where it can be much more easily irradiated away. Eventually this leads to an equatorial belt which is much cooler than the poles. These are a direct consequence of the confinement of the magnetic field lines to the crust which result in the presence of a large meridional field component in most of the crust.  Recently, Svidzinsky (\\cite{S03}) argued that the accumulation of magnetic field lines along the proton superconductor at the crust-core boundary, due to the Meissner-Ochsenfeld effect, produces an insulating barrier preventing heat to flow between the crust and the core. Our results are in the same line of thought but they do not confirm his claims of insulation of the crust from the core. The field geometries, i.e., the Stoke's stream functions (Eqs.~\\ref{equ:Stokes-r},~\\ref{equ:Stokes-theta}), we used in our calculations come from field evolution models (Page et al \\cite{PGZ00}) of crustal fields in which the migration, and accumulation, of the currents and the field lines toward the crust-core superconducting boundary was modeled in detail and, as illustrated in Fig.~\\ref{fig:heat_field}, they allow heat diffusion through the crust-core boundary. Nevertheless, other field evolution scenarios, e.g. expulsion of the magnetic flux from the core by a proton type I superconductor (Link \\cite{L03}; Buckley et al \\cite{BMZ04}) may produce a much stronger piling up of field lines tangentially to the crust-core boundary and result in more efficient thermal insulation.  In case the magnetic field is allowed to penetrate the core of the star, and assuming a star-centered dipolar geometry in the crust, we have shown that the stationary thermal state of the crust is very close to isothermality thus confirming, for this field geometry, the assumptions of the models of surface temperature distribution of magnetized neutron stars (e.g., Greenstein \\& Hartke \\cite{GH83} ; Page \\cite{P95}; Page \\& Sarmiento \\cite{PS96}; Shibanov \\& Yakovlev \\cite{SY96}) which considered only the insulating effect of a quantizing field in the envelope and assumed the rest of the crust to be isothermal.   The obvious difference between the isothermal crust and our crustal field results is demonstrated in Fig.~\\ref{fig:b3e12_cru_dens} and \\ref{fig:b3e12_10_8_dip_dens}. For a crustal field these differences cannot be neglected for highly magnetized neutron stars.  Since the warm polar region in case of a strong crustal field is much smaller than for star-centered and/or weak magnetic fields (see Fig.~\\ref{fig:surtemp_iso_crust_dip}) this may open a new way to distinguish between crustal and core magnetic fields: {\\em A strong crustal magnetic field implies a smaller effective area for thermally emitting cooling neutron stars}. This has consequences which can potentially be observed in X-ray in cases where fits of the thermal component of the spectrum of a cooling neutron star result in effective emitting radii $R_\\mathrm{eff}$ which are significantly smaller than the expected radius of a neutron star. For example, the ``Three Musqueteers'' (PSR 0656+14, PSR 1055-52 and Geminga: Tr\\\"umper \\& Becker \\cite{TB98}) as well as RX J185635-3754 (Pons et al \\cite{Petal02}) all have $R_\\mathrm{eff} \\sim 5-7$ km when their thermal spectra are fitted with blackbody spectra. If these radii would coincide with the radius of the neutron star, the equation of state describing the state of the core matter would have to be extremely soft. A relatively small warm polar region, created by a strong crustal field and emitting almost all the thermal radiation would be a reasonable explanation for such small  $R_\\mathrm{eff}$.   The differences in the photon luminosities for a star-centered or a crustal field will also affect the long term cooling of neutron stars. Future cooling calculations as well as the comparison of X-ray spectra and pulse profiles with model calculations assuming different field structures will open a way to discriminate the two basic scenarios: crustal or core magnetic field. Finally we mention the consequences the non-isothermality of the crust may have for the crustal field itself. Since the electric conductivity in the hot polar region is much smaller than in the equatorial layer, the field decay will be affected too and may cause differences in the field structure. This ``back reaction'' of the field onto its own decay via a field driven non-spherical symmetric crustal temperature distribution is subject of future investigations."
    },
    "0403/astro-ph0403488_arXiv.txt": {
        "abstract": "Powerful Type I X-ray burst with strong radial expansion  was observed from the low mass X-ray binary 4U 1820-30 with {\\it Rossi X-ray Timing Explorer}  on May 2, 1997. We investigate closely the flux profile during the burst expansion stage. Applying a semi-analytical model  we are able to uncover the behavior of a photospheric radius and to simulate the evolution of neutron star (NS)-accretion disk system.  We argue that although the bolometric luminosity is always the Eddington value $L_{Edd}$, the photon flux at the bottom of the expanded envelope can decrease during the expansion stage. In fact, at the initial moment of explosion when the bottom burning temperature is  $\\sim 2\\times10^9$ K,  the bottom flux $L_{bot}$ is a few times  the Eddington limit, because the electron cross-section is a few times less than the Thomson cross-section at such a high temperatures. The surplus of energy flux with respect to the Eddington, $L_{bot}-L_{Edd}$, goes into the potential energy of the expanded envelope. As cooling of the burning zone starts the surplus decreases and thus the envelope shrinks while the emergent photon flux stays the same $L=L_{Edd}$. At a certain moment the NS low-hemisphere, previously screened by the disk, becomes visible to the observer. Consequently, the  flux detected by  the observer increases.  Indeed, we observe to the paradoxical situation when the  burning zone  cools,  but the apparent flux increases because of the  NS-accretion disk geometry. We demonstrate a strong observational evidence of NS-accretion disk occultation in the behavior of the observed bolometric flux. We  estimate the anisotropy due to geometry and find that the system should have a high inclination angle. Finally, we apply an analytical model of X-ray spectral formation  in the neutron star atmosphere during burst  decay stage to infer NS mass-radius relation. ",
        "introduction": "\\src is one of the brightest low mass X-ray binaries (LMXB). It resides in the globular cluster NGC 6624. Extremely short 685-second orbital period discovered by \\citet{stella} from {\\it EXOSAT} observations implied that the system is ultra-compact and the secondary is a low-mass helium reach degenerate star. Using the analysis of UV diagrams of NGC 6624 \\citet[][ VLP hereafter]{vlp} evaluate the distance to be \\src by $6.4\\pm0.6$ kpc, while the optical observations give a distance estimate of 7.6 kpc \\citep{rml}. \\rxte observations revealed \\src as a prominent source of kilohertz quasi-periodic oscillations \\citep{smale}. Type I X-ray bursts \\citep{lpt,str2003} from \\src were discovered by \\citep{gg}. It is now well established that these short ($\\sim 10 - 15$ s) outbursts of energy are due to the unstable thermonuclear burning of hydrogen/helium mixture at the bottom of the NS atmosphere accumulated by accretion process. SAS-3 observations showed strong evidence that X-ray bursts are seen only in its low-intensity  state \\citep{clark}. All bursts observed from 4U 1820-30 show the radial expansion with an apparent photospheric radius increase by a factor of $\\sim 20$. Such an expansion is accompanied by spectral softening which moves the spectrum completely out of X-ray bandwidth and results in a so called ``precursor'' effect \\citep[see][and references therein]{str2003}. A several hour long ``superbursts'' was observed from \\src on September 9, 1999. It is now believed that superbursts are caused by the burning in the carbon ashes produced by Type I bursts \\citep{str2002}.   \\citet{ehs83} were first to study hydrodynamics of an illuminated atmosphere of X-ray  burst at the peak (Eddington) flux. Subsequently, the  X-ray spectral formation along with hydrodynamics during expansion and decay  stage of Type I bursts  were  studied  in \\citet{t94a}, hereafter T94. Further development of the expansion stage solution is presented in \\citet{st2002}, hereafter ST02. In these papers the important effect of the radiative expansion was pointed out. Specifically, while the emergent photon luminosity is always the Eddington,  $L_{Edd}$, the photon flux at the bottom of the expanded envelope can reach a value of a  few times greater than the Eddington limit. Indeed, at the initial moment of explosion when the bottom burning temperature is  $\\sim2\\times10^9$ K the bottom flux $L_{bot}$ is a few times more than Eddington luminosity, because the electron cross-section is a few times less than the Thomson cross-section for such a high temperatures (see below Eq. 1). The surplus of energy flux with respect to the Eddington limit, $L_{bot}-L_{Edd}$ goes into the potential energy of the expanded envelope. As cooling of the burning zone starts, the surplus decreases and thus the envelope shrinks (contraction stage of the radial expansion phase) while the emergent photon flux stays the same $L=L_{Edd}$. Shaposhnikov, Titarchuk \\& Haberl (2003), hereafter STH,  draw attention to the fact that at the certain moment of the contraction stage the NS low-hemisphere becomes visible to the observer, having previously been screened by the disk. In this paper we provide more details into the physical insight of this phenomenon and its relation to the data.  The spectral formation theory, employed in the present work, accounts for the effects of Comptonization, free-free absorption and emission. Analytical expressions are derived for color factors, and the spectral shapes are presented as functions of the input parameters: atmospheric chemical composition, NS mass, radius and the source distance. This technique was first applied to EXOSAT data from \\src and 4U 1705-44 in \\citet{ht}, hereafter HT95. In case of  4U 1820-30 data from seven bursts was used. For fixed valued of NS mass within the range 0.8-1.8 $M_\\odot$ authors constrained NS radius using assumed source distance of 6.4$\\pm0.6$ kpc. In particular, for $d=6.4$ kpc and $M_{NS}=1.4 M_\\odot$ they obtained the radius value of 7$\\pm0.4$ km. In \\citet{cygx2}, hereafter TS02 and STH where the technique is presented in detail, authors apply it to \\rxte data from Cygnus X-2 and 4U 1728-34 respectively. In this  Letter we  apply this methodology in analyzing the burst from 4U 1820-30 observed by RXTE on May 2, 1997. High spectral and statistical quality of the data allows us to infer tight constrains on the NS parameters. We find strong a evidence of accretion disk-star geometry evolution, that closely resembles the behavior of 4U 1728-34. We compare our results for 4U 1820-30 with equations of state (EOSs) of NS matter.     We present the Letter in the following manner. A brief description of the data  used in the analysis is given in \\S 2. In  \\S3 we consider the  general burst phenomenology. We specifically address its behavior during the expansion stage in \\S3.1 and calculate photospheric radius and bottom temperature profiles. We discuss in detail different processes which may affect the behavior of the source flux during the burst expansion. We present the model and the results of its application to the burst data of \\src in \\S 3.2. Specifically, we obtain the dependence of the NS mass on the radius as error contours, calculated for a set of distances to the system taken from the interval obtained by \\citet{vlp}. We discuss  our results and come to conclusions in \\S 4. ",
        "conclusions": "In this {Letter} we report the results of application of the theory of spectral formation during expansion and decay stage of a Type I X-ray burst using  data from LMXB 4U~1820-30 collected with \\rxte.  We investigate in detail the physics governing the burst properties during the expansion stage. The super-Eddington energy excess is deployed into gravitational potential energy of the expanded atmosphere which is a result of the dependence of the electron opacity on temperature. The emergent luminosity (cooling rate of the expanded atmosphere) is at the Eddington limit. When the bottom temperature drops and photospheric contraction occurs,  one can intuitively expect the  overall flux decrease. However the exact contrary situation is observed. The photospheric temperature increases and the bolometric flux grows presumably due to a geometrical NS recovery from behind the accretion disk inner edge. Simple arguments leads to the value of geometrically induced anisotropy of  $\\xi_b\\approx1.55-1.71$ which corresponds to the local system inclination angle $i\\approx73^\\circ- 80^\\circ$ (see Table 1).  We infer the mass-radius relationship and  error contours for a given distance by modeling the spectral temperature dependence on the bolometric flux. We include the effect of the radiation anisotropy due to change of the NS-disk geometry into consideration and we estimate the anisotropy coefficient from expansion stage behavior that is similar to 4U 1728-34 (STH). Because of the superior instrumental capabilities of {\\it RXTE} mission with respect to {\\it EXOSAT} satellite, the data even for  one burst observed with {\\it RXTE}  provides us more stringent constrains on the mass-radius relation than several bursts from {\\it EXOSAT}. Both statistical behavior of the model and Eddington flux limit strongly suggests that the distance to 4U 1820-30  close to 5.8 kpc, for which we obtain $M_{NS}=1.29^{+0.19}_{-0.07}M_\\odot$ and $R_{NS}=11.2^{+0.4}_{-0.5}$ km."
    },
    "0403/astro-ph0403455_arXiv.txt": {
        "abstract": "We measure the morphology--density relation of galaxies at $z{=}1$ across the full three orders of magnitude in projected galaxy density available in low--redshift studies. Our study adopts techniques that are comparable with those applied at lower redshifts, allowing a direct investigation of how the morphological segregation of galaxies has evolved over the last 8\\,Gyr. Although the morphology--density relation, as described by the fraction of early--type (E${+}$S0) galaxies, was in place at $z{=}1$, its form differs from that observed at both $z{=}0$ and $z{=}0.5$. In the highest density regions the early--type fraction has increased steadily with time from ${\\feso}{=}0.7{\\pm}0.1$ at $z{=}1$ to ${\\feso}{=}0.9{\\pm}0.1$ at the present epoch. However, in intermediate density regions corresponding to groups and the accretion regions of rich clusters, significant evolution appears to begin only after $z{=}$0.5. Finally, at the lowest densities, no evolution is observed for the early type fraction of field galaxies which remains constant at ${\\feso}{=}0.4{\\pm}0.1$ at all epochs. We examine a simple picture consistent with these observations where the early--type population at $z{=}1$ is comprised largely of elliptical galaxies. Subsequent evolution in both intermediate and dense regions is attributed to the transformation of spirals into lenticulars. Further progress in verifying our hypothesis may be achieved through distinguishing ellipticals and lenticulars at these redshifts through resolved dynamical studies of representative systems. ",
        "introduction": "In the local universe the fraction of galaxies with elliptical and lenticular (i.e.\\ early--type) morphologies is higher in clusters of galaxies than in less dense environments (Hubble 1926; Oemler 1974; Melnick \\& Sargent 1977; Dressler 1980). To first order, this morphology--density relation appears to be a universal characteristic of galaxy populations (e.g.\\ Postman \\& Geller 1984; Helsdon \\& Ponman 2003). In quantitative terms, morphological fractions correlate over three orders of magnitude in projected galaxy density ($\\Sigma$), thereby linking the properties of cluster galaxies (${\\Sigma}{\\simeq}1000\\,{\\rm Mpc^{-2}}$) with those of the field galaxy population (${\\Sigma}{\\ls}10\\,{\\rm Mpc^{-2}}$) (Dressler 1980).  The morphological segregation of galaxies is a generic prediction of cold dark matter simulations of large scale structure formation (Frenk et al.\\ 1985, 1988), and more recent semi--analytic galaxy formation models (Kauffmann 1995; Baugh et al.\\ 1996; Benson et al.\\ 2001; Diaferio et al.\\ 2001).  In that context, the observed morphology--density relation is interpreted as the combination of two mechanisms. Firstly, the local density of galaxies and dark matter is a proxy for the epoch of initial collapse of a given structure; the most massive structures at any epoch represent the earliest that collapsed. Secondly, interactions between galaxies, dark matter and the intra--cluster medium (i.e.\\ environmental processes) are likely to transform in--falling field galaxies from gas--rich spirals to gas--poor lenticular galaxies. The exact balance between these two mechanisms (i.e.\\ nature versus nurture) and the detailed physics of the environmental processes have yet to be identified unambiguously, and are the focus of much ongoing research (e.g.\\ Balogh et al.\\ 2001; Kodama \\& Smail 2001;  Treu et al.\\ 2003).  An important element of investigating the physics of morphological transformation is to trace the cosmic evolution of the morphology--density relation over the full range of projected density available locally. The timescales on which the relation evolves in different density regimes will hold important clues to the physical processes responsible.  To that end, Dressler et al.\\ (1997) used high--resolution imaging with the \\emph{Hubble Space Telescope (HST)} to measure the morphology--density relation in the core regions of a sample of rich clusters at $z{\\simeq}0.5$. Dressler et al. found that the fraction of lenticular galaxies in clusters declined by a factor of 2--3 between $z{=}0$ and $z{=}0.5$ and this evolution was accompanied by a corresponding increase in the fraction of star--forming spirals (see also Andreon 1998; Couch et al.\\ 1998; Fasano et al.\\ 2000; Treu et al.\\ 2003).  At higher redshifts, the distinction between elliptical and lenticular morphologies becomes increasingly difficult to draw (Smail et al.\\ 1997; Fabricant et al.\\ 2000).  Nevertheless, several authors have measured the total early--type fraction ${\\feso}$ in individual clusters at $z{\\simeq}1$ (e.g.\\ van Dokkum et al.\\ 2000, 2001; Lubin et al.\\ 2002).  These authors find ${\\feso}{=}0.5$ in clusters at $z{\\simeq}1$, i.e.\\ a smaller fraction than that found in the densest environments at $z{=}0$.  However, as van~Dokkum \\& Franx (2001) caution, these estimates are preliminary because they are based on a very small number of clusters.  In this paper we measure the morphology--density relation at $z{=}1$ across the full three orders of magnitude in galaxy density spanned in local samples.  We compare our results with those obtained at lower and intermediate redshifts (Dressler 1980; Dressler et al.\\ 1997; Treu et al.\\ 2003) and thus chart, for the first time, the form of the morphology--density relation over a cosmologically significant time interval (${\\sim}8\\,{\\rm Gyr}$).  A plan of the paper follows. In \\S\\ref{data} we develop a strategy for measuring the morphology--density relation at $z{=}1$ and summarize the data used for this purpose. Then in \\S\\ref{analysis} we describe the analysis, focusing separately on high-- and low--density environments.  The main results, the morphology--density relation at $z{=}1$ and its evolution to the present--day are presented in \\S\\ref{results}. In \\S\\ref{discuss} we discuss a possible interpretation, including how it relates to previous measurements of $\\feso$ in high--redshift clusters. We summarize our conclusions in \\S\\ref{conc}.  We parameterize the Hubble expansion as $h{=}H_0/100{\\rm km s^{-1}Mpc^{-1}}{=}0.65$, and adopt the currently favored values of $\\Omega_{\\rm M}{=}0.3$ and $\\Omega_\\Lambda{=}0.7$ when our analysis requires us to make distance estimates.  In this cosmology $1''{\\equiv}8.63{\\rm kpc}$ physical size at $z{=}1$. Unless otherwise stated, all error bars are stated at 1--$\\sigma$ significance.  All magnitudes are quoted in the Vega system. ",
        "conclusions": "\\label{conc}  We have used 52 individual \\emph{HST}/WFPC2 observations through the F814W filter, supplemented by panoramic ground--based imaging to measure the morphology--density relation of galaxies at $z{=}1$.  Our study adopts analysis methods similar to those developed at lower redshifts (e.g.\\ Dressler 1980) and our principal achievement is to span, at $z\\simeq$1, the full three orders of magnitude range in the projected number density of galaxies encompassed by the low redshifts studies. We choose to make a like--for--like comparison of the early--type fractions spanning field (${\\Sigma}{\\ls}$10\\,Mpc$^{-2}$), group (${\\Sigma}{\\simeq}$100\\,Mpc$^{-2}$) and rich cluster (${\\Sigma}{\\simeq}$1000\\,Mpc$^{-2}$) environments.  We briefly summarize our findings as follows:  \\begin{list}{(\\roman{fred})}{\\usecounter{fred}\\setlength{\\labelwidth}{5mm}\\setlength{\\itemindent}{0mm}\\setlength{\\labelsep}{1.5mm}\\setlength{\\leftmargin}{6.5mm}\\setlength{\\itemsep}{2mm}}  \\item Morphological segregation remains a prominent feature of the galaxy population at $z{=}1$, although the slope of the ${\\feso}$--${\\log}{\\Sigma}$ relation is ${\\sim}3$ times shallower than observed locally.  \\item The morphology--density relations at $z{=}1$ and $z{=}0.5$ are remarkably similar, with a significant difference only detected in the highest density bin. Most of the evolution producing the locally--observed relation occurred in the redshift interval 0$<z<$0.5.  \\item At low densities, the early--type fraction is roughly constant at ${\\feso}{=}0.4{\\pm}0.1$ across the full redshift range ($0{<}z{<}1$).  \\end{list}  These trends suggest to us a simple model whereby most cluster ellipticals formed at high redshift ($z{>}2$) with the bulk of the density--dependent growth arising from the environmental transformation of in--falling disk galaxies into lenticulars, and possibly merging of cluster galaxies at later times. This is motivated by the suggestive agreement (within the uncertainties) between the early--type fraction at $z{=}1$ in high density regions with the elliptical fraction observed at $z{=}0.5$.  Within the observational uncertainties, the majority of the model scenarios that we have explored are consistent with a negligible lenticular fraction at $z{=}1$, ${\\fso}{\\ls}0.1$. It is therefore possible that all cluster early--types at $z{=}1$ are ellipticals. To test this suggestion, resolved dynamical data is needed for a large sample of early--type cluster and field galaxies whose environmental densities can be measured."
    },
    "0403/astro-ph0403380_arXiv.txt": {
        "abstract": "We present a reanalysis of the X-ray data for RX~J0720.4--3125 presented in Zane et al.~(2002), using more data recently available from \\xmm\\/ and \\chandra. This analysis also corrects the \\ros\\/ data used in that paper to the TDB time system, incorporates the revised \\xmm\\/ barycentric correction available since then, and corrects the definition of the instantaneous period in the maximum likelihood periodogramme search. However, we are now unable to find a single coherent period that is consistent with all \\ros, \\chandra\\/ and \\xmm\\/ datasets.  From an analysis of the separate datasets, we have derived limits on the period change of $\\dot P = (1.4 \\pm 0.6) \\times 10^{-13}$~s/s at 99\\% confidence level. This is stronger than the value presented in Zane et al.~(2002), but sufficiently similar that their scientific conclusions remain unchanged. We examine the implications in more detail, and find that RX J0720.4--3125 can have been born as a magnetar provided that it has a young age of $\\sim10^4$ yr. A more conservative interpretation is that the field strength has remained relatively unchanged at just over $10^{13}$ G, over the $\\sim10^6$ yr lifetime of the star. ",
        "introduction": "\\label{int}  RXJ0720.4--3125 is a nearby, isolated neutron star (INS) originally discovered with \\ros\\/ during a systematic survey of the Galactic plane by Haberl et al. (1997). A clear modulation of the X-ray intensity is detected at a period of 8.391 s, which identifies the spin velocity of the neutron star.  RX J0720.4--3125 is, so far, the isolated neutron star with best studied timing properties. Its proximity and relatively high brightness made possible not only to measure the spin period, but also to study the period changes over a long-term. This is crucial, since the positive detection of spin-up or spin-down phases can shed light on the mechanisms that regulate the neutron star interaction with its surroundings. A large positive spin-down can indicate magneto-dipolar breaking, leading in turn to an estimate of magnetic field strength and star spin down age. In the attempt to measure a secular period change, Zane et al.~(2002, hereafter Z02) recently undertook a comprehensive timing analysis of \\ros, \\sax, \\chandra\\/ and \\xmm\\/ data spanning a total period of $\\sim 7$~yrs. Independently, a similar analysis based on an additional \\chandra\\/ dataset, but without the \\xmm\\/ data, was presented by Kaplan et al. (2002, hereafter K02). Although the two studies agreed in their main conclusion that the period change was less than $\\sim 3\\times10^{-13}$ s/s, it was evident that there were inconsistencies in the details of the analysis which have implications for further studies. The arrival times computed by the two groups were different. K02 also raised theoretical concerns about the validity of the {\\it coherent} analysis of the entire dataset carried out by Z02. That was the only method that, at that time, permitted the (positive) sign of the period change to be constrained. In fact, the phase-incoherent analysis presented by Kaplan et al.~(2002) still leads to a large uncertainty in $\\dot P$, and does not permit the discrimination between positive and negative spin-down.  The above-mentioned inconsistencies have motivated us to undertake a complete revision of the solution. At the time Z02 was published, timing inconsistencies between \\xmm\\/ and radio data on other pulsars had come to light (Kuster et al.~2001). This was traced to an error in the \\xmm\\/ Science Analysis System (SAS) task {\\it barycen} before version 13.1 [the spacecraft position vector was constructed as (x, y, y) instead of (x, y, z)]. Although we argued in Z02 that any effect on the coherent analysis was small, we nevertheless performed a reanalysis of the \\xmm\\/ data to check, and also to investigate whether this corrected the inconsistencies with the K02 analysis. We found indeed that the effect was negligible. The implication is therefore that the inconsistencies between the Z02 and K02 analyses remained.  We have therefore undertaken a complete reanalysis of the dataset to investigate and eliminate these inconsistencies. This has brought to light that the times used for the \\ros\\/ data in both the Z02 and K02 studies are in UTC rather than TDB as used for \\chandra\\/ and \\xmm; the times of arrival (ToAs) in K02 are incorrectly calculated; and that the formula used to calculate the instantaneous phase was incorrect in Z02. Substantial new datasets are also available from both \\xmm\\/ and \\chandra\\/ which permit a more comprehensive investigation into the period evolution. We report here the results of this reanalysis. ",
        "conclusions": "\\label{sect:con}   We have eliminated the discrepancies between the calculated ToAs in the K02 paper and those derived from the same data analysed by Z02. This invalidated our derived $[P_0, \\dot P]$ solutions in Z02 (K02 have also updated their ToAs). In order to correct these solutions, we have reanalysed the datasets and taken the opportunity to incorporate significant new datasets from both \\chandra\\/ and \\xmm. At the same time, we have revised the model used by Z02 to be fitted to the data to retain the correct phase reference.  Taken individually, the combination of old and new datasets allow a significantly more accurate $[P_0, \\dot P]$ solution to be derived than that in K02 or Z02. This {\\it establishes that RX J0720.4--3125 is spinning down}.  We have also attempted a phase-coherent analysis of the combined datasets, as in Z02. However, we have been unable to find a consistent solution, even when we analyse the data from different satellites separately. In this case the \\xmm\\/ data are not consistent with the other data. We have explored the possible reasons for this: they may be intrinsic to RX J0720.4--3125 (in which case the constant $\\dot P$ model is not appropriate), or they may result from continuing timing errors in the data we have not succeeded in identifying. We present the solutions from the combined \\ros\\/ and \\chandra\\/ data and the \\xmm\\/ data separately.  The individual (incoherent) $[P_0, \\dot P]$ solution is sufficiently accurate to allows us to place strong constraints on the mechanism causing the spin-down. The large $\\dot P \\sim 10^{-13}$~s/s rules out accretion and propeller origins and points toward an interpretation in terms of magneto-dipolar breaking. This in turn constrains the age of RX J0720.4--3125 and its magnetic field strength.  For a magnetic field as strong as $B = (2.8-4.2) \\times 10^{13}$~G, the electron cyclotron line is unaccessible (it falls at about 350 keV).  However, a relatively broad proton cyclotron absorption feature is predicted at $\\approx 0.2-0.3$~keV (Zane et al.,~2001). The presence of a cyclotron absorption line at the low energy of the sensitive energy band of the \\xmm\\/ EPIC instruments has been indeed recently reported (see Haberl et al.,~2003a, 2004). The feature may explain the observed deviations from a Planckian shape at these energies, and possibly also the anticorrelation of the modulation of hardness ratio and total X-ray intensity detected by {\\xmm}\\/ (Cropper et al.,~2001).  Our study allows a possible evolutionary link between dim INS and AXPs (which relies on the interpretation of the former as warm-out magnetars) but only if the INS have an age of $\\sim10^4$ yr.  A more conservative interpretation is that RX J0720.4--3125 was born with a strong, but not superstrong, field [$ B = (2.8-4.2) \\times 10^{13}$~G] which is compatible with those of the canonical radio-pulsars. Similar conclusions have been argued also by K02, however their measure of $\\dot P$ has a large error and does not permit the discrimination between this and the separate possibilities that RX J0720.4--3125 has a more conventional field ($B\\sim 10^{12}$~G) and lower spin-down rate ($\\dot P \\sim 10^{-15}$s/s) or is even spinning-up. Comparing with the radio-active pulsar population, objects with $B>10^{13}$~G are rare, but their evidence is rapidly growing (Camilo et al.,~2000). The parameters of RX J0720.4--3125 we have derived  are not too dissimilar from those of PSR~J1814--1744, which has $P\\sim 4$~s and $\\dot P \\sim 7.4 \\times 10^{-13}$~s/s.  One of the most striking mysteries about \\ros \\/ isolated neutron stars, AXPs, soft-$\\gamma$ repeaters and objects like Geminga is why they do not exhibit radio emission.  As suggested by Motch et al.~(2003), who derived very similar conclusions basing on the proper motion measurement, it could thus be that RX J0720.4--3125 and several others apparently radio-quiet neutron stars are just radio pulsars with off-axis beam that does not cross the Earth. The radio beam narrows with increasing periods (Briggs 1990), making this explanation even more plausible. On the other hand, evidence for a population of genuinely radio-silent young neutron stars arises from population synthesis (Neuh\\\"auser \\& Tr\\\"umper,~1999; Gotthelf \\& Vasisht, 2000; Popov et al.,~2000a,b and references therein). The group of ROSAT INSs is very homogeneous and they are all relatively close-by (within 300-400 pc), making unplausible that all their beams are misaligned. At present, these two possibilities cannot be distinguished. Detailed timing analysis of other \\ros \\/ INSs is crucial to reach a more comprehensive understanding of the entire population."
    },
    "0403/astro-ph0403663_arXiv.txt": {
        "abstract": "We have performed Chandra observations during the past 3 years of 5 of the M31 supersoft X-ray sources discovered with ROSAT. Surprisingly, only one of these sources has been detected, despite a predicted detection of about 20-80 counts for these sources. This has motivated a thorough check of the ROSAT M31 survey I data, including a relaxation of the hardness ratio requirement used to select supersoft sources. This increases the number of supersoft sources identified in survey I by 7. We then carried out a comparison with the ROSAT M31 survey II dataset which had hitherto not been explicitly investigated for supersoft X-ray sources. We find that most of the ROSAT survey I sources are not detected, and only two new supersoft sources are identified. The low detection rate in the ROSAT survey II and our Chandra observations implies that the variability time scale of supersoft sources is a few months. If the majority of these sources are close-binary supersoft sources with shell hydrogen burning, it further implies that half of these sources predominantly experience large mass transfer rates. ",
        "introduction": "Observations during the past decade have suggested the definition a new class of sources. Luminous supersoft X-ray sources (SSSs) have luminosities in the range $10^{35}-10^{38}$ erg/s and $kT$ in the range $20-80$ eV, with no hard X-ray component of comparable luminosity. Some SSSs are simply hot white dwarfs (e.g., post novae), or pre-white-dwarfs (in planetary nebulae). What is most intriguing  about SSSs, however, is the fact that the physical nature of a majority of the sources with optical IDs is not yet understood. These more mysterious sources include the prototypes CAL 83 and CAL 87, discovered with Einstein (Long, Helfand, \\& Grabelsky 1981), and more numerously with ROSAT (e.g. Greiner 2000). The most promising explanation for the majority of the sources invokes quasi-steady nuclear burning of matter accreting onto the surface of a white dwarf (WD) to generate these systems' prodigious flux (see, e.g., van den Heuvel \\etal\\ 1992). There is indirect evidence in favor of these models for several of the sources. The binary sources which are so luminous that nuclear-burning models seem to be required, will be referred to as close-binary supersoft sources (CBSS).  Observing supersoft sources  in M31 has the advantage that several questions can be attacked more easily as compared to the local sources (including those in the Magellanic Clouds): (i) What is the spatial distribution over the galaxy and possible correlations with different environment? (ii) What is the size of the population including the ratio of SSS to other types of low-mass X-ray binaries? (iii) What is the variability pattern (if any) and duty cycle? All these questions can help in providing clues to the nature of the sources.   \\begin{deluxetable}{rlccccc} \\tablewidth{16cm} \\tablecaption{New SSS from the ROSAT PSPC} \\tablehead{ No &  ~~~~~Name & Coordinate  & Error    & count rate & HR1 & HR2 \\\\ &        &   (2000.0)  & (\\asec)& (cts/ksec) & &      } \\startdata \\noalign{\\smallskip} \\multicolumn{7}{c}{First observation during PSPC Survey I} \\\\ 4& RX J0039.3+4047~       &   00\\h39\\m21\\fss4 +40\\degs47\\amin41\\asec & 42 &  0.26$\\pm$0.23 & $-$0.89$\\pm$0.10 & $-$0.29$\\pm$0.65  \\\\ 6& RX J0039.7+4030~       &   00\\h39\\m47\\fss1 +40\\degs30\\amin05\\asec & 15 &  2.03$\\pm$0.30 & $-$0.85$\\pm$0.10 & $-$0.83$\\pm$0.53  \\\\ 7& RX J0039.8+4053~       &   00\\h39\\m50\\fss4 +40\\degs53\\amin38\\asec & 23 &  1.07$\\pm$0.25 & $-$0.75$\\pm$0.18 & ~~0.44$\\pm$0.97   \\\\ 9& RX J0040.4+4013~       &   00\\h40\\m28\\fss6 +40\\degs13\\amin44\\asec & 23 &  0.50$\\pm$0.27 & $-$0.85$\\pm$0.14 & ~~0.72$\\pm$1.00  \\\\ 14& RX J0042.7+4107~       &   00\\h42\\m44\\fss9 +41\\degs07\\amin18\\asec & 22 &  1.04$\\pm$0.31 & $-$0.89$\\pm$0.16 & $-$0.65$\\pm$1.00  \\\\ 17&  RX J0044.2+4117~       &   00\\h44\\m14\\fss0 +41\\degs17\\amin57\\asec & 34 &  0.95$\\pm$0.35 & $-$0.97$\\pm$0.25 & $-$0.58$\\pm$0.53  \\\\ 23&  RX J0047.6+4159~       &   00\\h47\\m42\\fss3 +41\\degs59\\amin59\\asec & 36 &  1.23$\\pm$0.44 & $-$0.82$\\pm$0.28 & $-$0.20$\\pm$1.00  \\\\ \\noalign{\\smallskip} \\multicolumn{7}{c}{First observation during serendipituous Pointing} \\\\ 21&  RX J0047.4+4157~       &   00\\h47\\m27\\fss2 +41\\degs57\\amin34\\asec & 25 &  0.60$\\pm$0.18 & $-$0.98$\\pm$0.21 & ~~0.00$\\pm$0.30  \\\\ \\enddata \\vspace{-0.3cm} \\label{newsssi} \\end{deluxetable}    ROSAT has observed the full disk of the M31 galaxy (about 6.5 deg$^2$) twice. A ROSAT PSPC mosaic of 6 contiguous pointings with an exposure time of 25 ksec each was performed in July 1991 (first M31 survey; Supper \\etal\\ 1997). A second survey was made in July/August 1992 and January 1993 (Supper \\etal\\ 2001). Until now, only the first survey has been investigated systematically for SSS (Greiner \\etal\\ 1996a).  This paper is the second in a series dealing with SSSs in M31, and in particular with their variability properties. The first (Di Stefano \\etal\\ 2004) concentrated on the analysis of several sets of {\\it Chandra} data: (1) $3$ separate $15$ ksec observations of each of $3$ disk fields, and (2) a $40$ ksec ACIS observation of the bulge with the backside-illuminated (BI) chips, combined with information gleaned from $2$ years of regular ACIS front-side illuminated (FI) chips monitoring of the bulge. In fact the disk fields were observed such that the locations of 5 ROSAT SSSs (\\#2, 3, 12, 19, 20) were covered by the BI chips which exhibit enhanced sensitivity for very soft X-rays. Another four SSS (\\#1, 14, 24, 25) are covered by chance coincidence with the FI chips due to the field rotation between the different epochs. The results of that paper relevant to this second paper can be summarized as follows: \\begin{itemize} \\item That paper established that only one of the ROSAT-discovered sources, RX J0038.6+4020, was detected by {\\it Chandra}.  \\item No new supersoft source obeying the same criteria as those applied for the selection in ROSAT data, has been found in any of these Chandra pointings. However, with a modified hardness ratio criterion a total of $16$ new SSSs that are not associated with foreground or background objects, and which are therefore likely members of M31, were discovered in the disk fields. Not all of these $16$ were luminous enough to have been detected by ROSAT; $6$ provided fewer than $20$ counts. Furthermore, some appear to be hard enough not to have been selected as SSSs using the procedures applied to the ROSAT M31 survey data. Nevertheless, at least $3$ of the sources with more than $20$ counts would have likely been selected as ROSAT SSSs. Interestingly enough, it could be established that $2$ of these $3$ sources are transient by comparing the flux at these positions among different {\\it Chandra} pointings or by studying data taken with {\\it XMM-Newton}.  \\item The bulge of M31 is rich in high-luminosity SSSs. By comparing among different {\\it Chandra} pointings or by studying data taken with {\\it XMM-Newton} it has been found that 12 of 16 bulge sources are transient, and one additional source is highly variable. \\end{itemize}  The non-detection of 4 out of 5 ROSAT-discovered supersoft X-ray sources, combined with the failure of finding any new sources with similarly soft X-ray spectra appears to be puzzling for at least two reasons. First, if one assumes sources with constant brightness, could it be possible that the non-detection with Chandra is due to spurious detections with ROSAT? Second, if one assumes sources with variable X-ray emission, why do we not detect as many new sources with Chandra as we miss because they faded away between the ROSAT detection and the Chandra observation?  To answer these questions, we have embarked on a comprehensive re-analysis of the ROSAT data. In particular, in this paper we slightly revise the hardness ratio criterion used to select supersoft sources in the ROSAT survey I data (\\S 2), we analyze the second ROSAT PSPC survey for SSS with the same criteria (\\S 4) plus serendipitous PSPC observations (\\S 5), and we present the first survey of ROSAT supersoft sources with the {\\it Chandra} Observatory (\\S 3), and for completeness also include the  public \\xmm\\ observations (\\S 6). We finally discuss the variability in (\\S 7). ",
        "conclusions": "The evidence from ROSAT, Chandra, and XMM is that SSSs tend to be highly variable, perhaps more variable than any other class of X-ray binary, most notably the hard sources comprising  a substantial number of X-ray binaries (Trudolyubov \\etal\\ 2002b). A large fraction (30\\%) of SSS are transients with turn-off or turn-on times on the order of a few months. The majority of the sources that have fallen below detectability limits have not been detected again. This may argue that the duty cycle is low, while activity times are on the order of months or years. With an on-time duration of months, and a duty cycle of, e.g., 40\\%, we have only a ten percent chance of detecting a source which was \"on\" during one observation, in a second uncorrelated observation, and a 16\\% chance of detecting any given SSS in 2 unrelated observations, consistent with our finding.  In addition, the spatial coverage of M31 with Chandra was small (less than 5\\%), since the coverage for SSS by Chandra is primarily given by the S3 chip. Thus, the likelihood of detecting new SSS with Chandra was small. Population studies have estimated the total SSS population in M31 to be $\\sim 1000$ (Di\\thinspace Stefano \\& Rappaport 1994). With an assumed duty cycle of 10\\% this would correspond to a density of active SSS of 3$\\times$10$^{-3}$ arcmin$^{-2}$, or 0.02 per S3 chip.   While it is unlikely that novae are responsible for the strong X-ray variability in SSS, its physical cause  remains to be explained. Both, better sampling of the light curve as well as optical identifications and subsequent optical monitoring seem to be required to deduce insight into the variability mechanism(s).  We finally note that the more frequent Chandra and XMM observations over the last three years have revealed a number of supersoft  X-ray transients (e.g. Shirey 2001, Trudolyubov \\etal\\ 2002c). While their nature remains to be established as well, they support the notion of the strong variability of SSS."
    },
    "0403/astro-ph0403513_arXiv.txt": {
        "abstract": "We use {\\it Far Ultraviolet Spectroscopic Explorer} (FUSE) and {\\it Space Telescope Imaging Spectrograph} (STIS) spectra to study intergalactic absorption towards the quasar PG\\,1259+593 ($z_{\\rm em}=0.478$), with a particular emphasis on the warm-hot intergalactic medium (WHIM). The combined FUSE and STIS spectrum of PG\\,1259+593 covers the full wavelength range between $905$ and $1730$ \\AA\\, at a spectral resolution of $\\sim 25$ km\\,s$^{-1}$ for the FUSE bandpass ($\\lambda \\leq 1180$ \\AA) and $\\sim 7$ km\\,s$^{-1}$ for the STIS range ($\\lambda > 1150$ \\AA). The signal-to-noise ratios (S/N) per resolution element are $\\sim 10-30$ (FUSE) and $\\sim 7-17$ (STIS). We identify 135 intergalactic absorption lines with equivalent widths $\\geq 10$ m\\AA, tracing 78 absorption components in 72 Ly\\,$\\alpha$/Ly\\,$\\beta$ absorption-line systems. Metal-line absorption by species such as C\\,{\\sc iii}, C\\,{\\sc iv}, O\\,{\\sc iii}, O\\,{\\sc iv}, O\\,{\\sc vi}, and Si\\,{\\sc iii} is clearly detected in four systems, and is possibly seen in four additional cases. We study the distribution and physical properties of the WHIM as sampled by O\\,{\\sc vi} and intrinsically broad Ly\\,$\\alpha$ lines. The number of intervening O\\,{\\sc vi} absorbers for equivalent widths $W_{\\lambda} \\geq 24$ m\\AA\\, is $3-6$ over an unobscured redshift path of $\\Delta z\\approx 0.368$. This implies a number density of O\\,{\\sc vi} systems, $dN_{\\rm OVI}/dz$, of $\\sim 8-16$ above this equivalent width limit along this sight line. A seventh intervening O\\,{\\sc vi} absorber is possibly detected with $W_{\\rm r}($O\\,{\\sc vi})$\\approx15$ m\\AA. The range of $dN_{\\rm OVI}/dz=8-16$ for $W_{\\lambda} \\geq 24$ m\\AA\\, is consistent with estimates from other sight lines, supporting the idea that intervening intergalactic O\\,{\\sc vi} absorbers contain an substantial fraction of the baryonic mass in the low-redshift Universe. We identify a number of broad Ly\\,$\\alpha$ absorbers with large Doppler parameters ($b\\sim40-200$ km\\,s$^{-1}$) and low column densities ($N$(H\\,{\\sc i})$< 10^{14}$ cm$^{-2}$). For pure thermal broadening, these widths correspond to temperatures of $\\sim 1\\times10^5-3\\times10^6$ K. While these broad absorbers could be caused by blends of multiple, unresolved lines, continuum undulations, or by kinematic flows and Hubble broadening, we consider the possibility that some of these features are single-component, thermally broadened Ly\\,$\\alpha$ lines. These systems could represent WHIM absorbers that are too weak, too metal-poor, and/or too hot to be detected in O\\,{\\sc vi}. If so, their widths and their frequency in the PG\\,1259+593 spectrum imply that these absorbers trace an even larger fraction of the baryons in the low-redshift Universe than the O\\,{\\sc vi} absorbing systems. A thermal Doppler broadening explanation for one of these broad features is supported by the probable detection of O\\,{\\sc vi} near the velocity of a broad Ly\\,$\\alpha$  and Ly\\,$\\beta$ absorber with an O\\,{\\sc vi} line width $\\sim 4$ times smaller than for H\\,{\\sc i}. ",
        "introduction": "The analysis of absorption lines in the local and distant intergalactic medium (IGM) towards extragalactic background sources such as quasars (QSOs) and active galactic nuclei (AGNs) has become an important research area to investigate the gaseous environment of galaxies and galaxy-clusters, and to study the large scale structure, the chemical evolution and the baryonic content of the IGM. The most sensitive tracer to follow the distribution and evolution of the IGM from high to low redshifts is the Ly\\,$\\alpha$ line of neutral hydrogen. A large number of Ly\\,$\\alpha$ lines occur along a typical QSO sight line, resulting in the ``Ly\\,$\\alpha$ forest'' (e.g., Rauch 1998; Lu et al.\\,1996). For the low-redshift Universe, the observation of intergalactic Ly\\,$\\alpha$ absorption requires space-based instrumentation, such as the former {\\it Goddard High Resolution Spectrograph} (GHRS) and the more recently deployed {\\it Space Telescope Imaging Spectrograph} (STIS), both installed on the {\\it Hubble Space Telescope} (HST). For higher redshifts ($z>1.6$), ground based 8\\,m-class telescopes such as Keck and the {\\it Very Large Telescope} (VLT) provide a wealth of information about the Ly\\,$\\alpha$ forest in the early Universe (e.g., Kim et al.\\,2002). It is generally accepted that at $z=3$ almost all of the baryons in the Universe are located in the highly-ionized Ly\\,$\\alpha$ forest (Rauch et al.\\,1997; Weinberg et al.\\,1997), whereas at lower redshifts an increasing fraction of the baryons may reside in a low-density, shock-heated component at temperatures between $10^5$ to $10^7$ K, the so-called warm-hot intergalactic medium (WHIM). In the local Universe (i.e., near $z=0$), the WHIM may contain as much as 30 to 40 percent of the baryons (e.g., Cen \\& Ostriker 1999). In comparison, the Ly\\,$\\alpha$ forest at $z=0$ probably contains about 20 percent of the baryons (Penton, Shull \\& Stocke 2000). Detecting the WHIM phase at low redshifts is a challenging task. Absorption spectroscopy of the O\\,{\\sc vi} doublet ($\\lambda\\lambda 1031.9, 1037.6$) currently provides one of the best diagnostics to study the warm phase ($10^5-10^6$ K) of the WHIM at low redshift. However, observing O\\,{\\sc vi} in the low-redshift IGM requires space-based high-resolution spectrographs such as STIS and the {\\it Far Ultraviolet Spectroscopic Explorer} (FUSE) together with a substantial amount of observing time (e.g., Tripp, Savage \\& Jenkins 2000; Savage et al.\\,2002). Due to these instrumental restrictions, the number of sight-lines that have been analyzed for O\\,{\\sc vi} absorption at $z<1$ unfortunately is still very limited. So far, intervening WHIM O\\,{\\sc vi} absorption for $z<1$ in FUSE and STIS data has been found towards H\\,1821+643 (Tripp, Savage \\& Jenkins 2000; Oegerle et al.\\,2000) , PG\\,0953+415 (Tripp \\& Savage 2000; Savage et al.\\,2002), and PKS\\,0405$-$123 (Chen \\& Prochaska 2000). Together with the non-detections or marginal and weak detections reported for the sight lines towards PG\\,0804+761 (Richter et al.\\,2001a) and 3C\\,273 (Sembach et al.\\,2001), the current data imply a number of intervening O\\,{\\sc vi} absorbers per unit redshift of $dN_{\\rm OVI}/dz=14^{+9}_{-6}$ for restframe equivalent widths $W_r \\geq 50$ m\\AA\\, at $\\langle z \\rangle =0.09$ (Savage et al.\\,2002). Assuming that the average metallicity of these absorbers is 0.1 solar, the cosmological mass density in the units of the current critical density is $\\Omega_b($O\\,{\\sc vi}$)\\geq 0.002 h_{75}^{-1}$. Clearly, additional observations with FUSE and STIS are desired to improve the statistics on the intervening O\\,{\\sc vi} absorbers, and to gain new insights about their physical properties (e.g., their ionization properties) and their metal abundances. For those absorbers that cannot be detected in O\\,{\\sc vi} absorption (e.g., if the overall column density is too low, the metallicity is too low, or the temperature is too high), other tracers have to be used to study the warm-hot intergalactic gas phase. X-ray observations with {\\it Chandra} and {\\it XMM-Newton} (e.g., Rasmussen, Kahn, \\& Paerels 2003; Fang et al.\\,2002) hold the prospect of improving our understanding of the distribution and physical properties for the million-degree phase of the WHIM, which may contain the dominating fraction of the baryons at low redshifts. However, detecting X-ray signatures such as O\\,{\\sc vii} and O\\,{\\sc viii} absorption from the low-redshift IGM is difficult with current instrumentation, and so far the information about the intergalactic X-ray forest is relatively sparse. For high-metallicity, high-column density absorbers, one may search for FUV absorption by Ne\\,{\\sc viii} $\\lambda\\lambda 770.4, 780.3$. But due to the relatively low cosmic abundance of Ne, Ne\\,{\\sc viii} is not expected to be a sensitive tracer for the WHIM in the low-redshift Universe. A promising approach to study the WHIM at low redshifts is to search for broad Ly\\,$\\alpha$ lines. For pure thermal broadening, Ly\\,$\\alpha$ absorbers with Doppler parameters exceeding $b\\approx40$ km\\,s$^{-1}$ sample WHIM gas with temperatures $T>10^5$ K. There are some compelling detections of broad Ly\\,$\\alpha$ lines at low redshifts (e.g., Tripp et al.\\,2001), but generally broad Ly\\,$\\alpha$ lines are difficult to identify and analyze due to continuum placement uncertainties, line-blending problems, and relatively low signal-to-noise ratios in existing data.  In this paper we investigate intervening intergalactic absorption by H\\,{\\sc i}, O\\,{\\sc vi} and other species with FUSE and STIS in direction of the bright QSO PG\\,1259+593 ($V=15.84$; $z_{\\rm em}=0.478$; $l=120\\fdg6$, $b=+58\\fdg1$). FUSE and STIS data of PG\\,1259+593 have been extensively used to study abundances and physical conditions in Galactic halo gas in this direction (Richter et al.\\,2001b; Collins, Shull \\& Giroux 2003; Sembach et al.\\,2004; Fox et al.\\,2004). With more than $600$ ks of integration time it represents one of the prime objects of the FUSE mission. The line of sight towards PG\\,1259+593 samples several strong Ly\\,$\\alpha$ systems, as shown by low-resolution {\\it Faint Object Spectrograph} (FOS) data (Bahcall et al.\\,1993). The goal of this paper is to present an overview of the intergalactic absorption towards PG\\,1259+593, to identify the main absorption systems, to perform a first, basic analysis of these absorbers, and to search for the WHIM absorbers that may contain an important fraction of the baryonic mass in the low-redshift Universe. The outline of this paper is the following: in \\S2 we review the observations and the data reduction process. The analysis of intergalactic absorption is presented in \\S3. In \\S3 we also describe the properties of the individual absorption systems and derive statistical information about the IGM towards PG\\,1259+593. We discuss our results in \\S4 and summarize our study in \\S5. ",
        "conclusions": "As the temperature of the intergalactic medium undergoes a significant change from high to low redshifts, a large fraction of the baryonic matter in the local Universe is expected to reside in the WHIM phase - a gas phase that is particularly difficult to detect. The combined FUSE and STIS spectrum of PG\\,1259+593 represents one of the few high-quality UV and FUV data sets currently available to study the distribution and physical properties of the low-redshift WHIM and the local Ly\\,$\\alpha$ forest. However, the PG\\,1259+593 data also show how difficult the analysis of the local IGM is due to the limitations in S/N and spectral resolution in these space-based data. We have detected three intervening O\\,{\\sc vi} absorption line systems towards PG\\,1259+593 and have found four additional candidates for which we cannot claim a firm detection. With a number of intervening O\\,{\\sc vi} systems per unit redshift of $dN_{\\rm OVI}/dz=8-16$ for equivalent widths $W_{\\rm r}\\geq24$ m\\AA, the sight line towards PG\\,1259+593 adds important information about the distribution of intervening O\\,{\\sc vi} absorbers in the local Universe. Savage et al.\\,(2002) summarize previous FUSE and STIS absorption line measurements of intervening O\\,{\\sc vi} in the low-redshift IGM, including the lines of sight toward H\\,1821+643, PG\\,0953+415, PG\\,0804+761, and 3C\\,273 (Tripp, Savage \\& Jenkins 2000; Oegerle et al.\\,2000; Tripp \\& Savage 2000; Richter et al.\\,2001a; Sembach et al.\\,2001). Six O\\,{\\sc vi} systems are detected above an equivalent width limit of $50$ m\\AA\\, for the O\\,{\\sc vi} $\\lambda 1031.9$ line, suggesting that $dN_{\\rm O\\,VI}/dz=14^{+9}_{-6}$ for $\\left<z\\right>=0.09$, and therefore $\\Omega_b($O\\,{\\sc vi}$)\\geq0.002\\,h_{75}^{-1}$ for an average metallicity of $0.1$ solar (for details see Savage et al.\\,2002). The line of sight towards PG\\,1259+593 consists of $3-5$ more intervening O\\,{\\sc vi} systems with $W_{\\rm r}\\geq50$ m\\AA, in agreement with the previous observations and $dN_{\\rm O\\,VI}/dz$ estimates. The problem of having to separate clearly detected O\\,{\\sc vi} systems from the many tentative absorbers yet remains somewhat unsatisfying, and more accurate sight-line analyses appear to be necessary to reduce the yet considerable uncertainty for $\\Omega_b($O\\,{\\sc vi}$)$.  The FUSE and STIS absorption line measurements suggest that intervening O\\,{\\sc vi} absorption lines reliably trace the WHIM in the low-redshift Universe, although some O\\,{\\sc vi} may also arise in low-density, photoionized gas (see Savage et al.\\,2002). We note at this point that the collisional ionization interpretation for most of the weak O\\,{\\sc vi} absorbers holds only for the low-redshift Universe, where a significant fraction of the intergalactic gas is expected to reside in the WHIM phase (Cen \\& Ostriker 1999) at $T\\geq10^5$ K. At higher redshifts ($z\\sim 2$), a larger fraction of the O\\,{\\sc vi} absorbing systems apparently traces intergalactic structures that are photoionized by a hard UV background. This is evident from the narrow line widths ($b_{\\rm O\\,VI}\\leq 10$ km\\,s$^{-1}$) of some of the weaker O\\,{\\sc vi} absorbing structures at high redshifts with log $N$(H\\,{\\sc i}$)\\leq 15$ (Bergeron et al.\\,2002; Carswell, Schaye \\& Kim 2002). Such small $b$ values are inconsistent with temperatures $T\\geq10^5$ K. However, the results of Simcoe, Sargent, \\& Rauch (2002) imply that collisional ionization has to be taken into account for the higher-column density O\\,{\\sc vi} systems at $z\\sim2$ (log $N$(H\\,{\\sc i}$)\\geq 15$), as this gas possibly traces shock-heated material from supernova-driven galactic winds. An important but yet unanswered question is, what baryonic mass fraction of the WHIM at low $z$ is not detected in studies that concentrate solely on O\\,{\\sc vi} absorption ? WHIM clouds that have intrinsically low gas column densities, and/or low metallicities, and/or very high temperatures ($\\geq 10^{5.5}$ K) are difficult to detect in O\\,{\\sc vi} absorption. Other high ions such as Ne\\,{\\sc viii} may trace the WHIM in a higher temperature regime, but our study (and other measurements) provide no evidence that there exists a significant number of detectable intervening Ne\\,{\\sc viii} absorbers at low redshifts. Towards PG\\,1259+593 we have not found significant ($\\geq 3\\sigma$) Ne\\,{\\sc viii} absorption in any of the absorption systems. The possibility to detect Ne\\,{\\sc viii} absorption in FUSE and STIS data, however, is probably low due to the relatively low cosmic abundance of neon compared to oxygen (log (Ne/O)$_{\\sun}=-0.61$; Allende\\,Prieto, Lambert, \\& Asplund 2001, 2002; Anders \\& Grevesse 1989) and the lower S/N of the FUSE data in the spectral region where the Ne\\,{\\sc viii} lines are accessible. In two systems ($z=0.22313$ and $z=0.25971$) Ne\\,{\\sc viii} absorption may be present at low statistical confidence ($\\leq 1.9\\sigma$); if more Ne\\,{\\sc viii} data from other sight lines will become available, one may find evidence for Ne\\,{\\sc viii} absorption in WHIM absorbers by stacking together a larger number of Ne\\,{\\sc viii} velocity profiles. Observations of the X-ray absorption in the lines of O\\,{\\sc vii}, O\\,{\\sc viii} and other species are becoming an important way to trace hot gas in the local intergalactic medium (Rasmussen, Kahn \\& Paerels 2003; Fang et al.\\,2002; Nicastro et al.\\,2002). But even with the most modern X-ray observatories, such as {\\it Chandra} and {\\it XMM-Newton}, observations of intervening intergalactic X-ray absorbers represent a challenging task, and so far the number of detected X-ray absorption features in the IGM is quite low. Thus, good statistics about the low-redshift X-ray forest will not be readily available in the near future.  A promising approach to trace the WHIM at temperatures up to $T\\sim10^6$ K is to search for broad, shallow Ly\\,$\\alpha$ absorbers. With the FUSE and STIS data presented here we possibly have found for the first time evidence that at least part of the WHIM up to $T\\sim10^6$ K may be traced by Ly\\,$\\alpha$ absorption. Support for the detectability of gas this hot in Ly\\,$\\alpha$ comes from X-ray observations. Studies of the immediate intergalactic environment of the Milky Way suggest the possible presence of an WHIM filament with an O\\,{\\sc vii} column density of $\\sim 10^{16}$ cm$^{-2}$ (see Rasmussen, Kahn, \\& Paerels 2003). If this O\\,{\\sc vii} column density is representative for a WHIM absorber, and if we further assume a temperature of $10^6$ K together with collisional ionization equilibrium (Sutherland \\& Dopita 1993) and a metallicty of the gas of 0.1 solar, we would expect to see Ly\\,$\\alpha$ absorption at the level of log $N$(H\\,{\\sc i})$\\approx13.5$ and $b\\approx130$ km\\,s$^{-1}$, in line with some of the broader absorbers towards PG\\,1259+593 listed in Table 5. Depending on the column density and temperature distribution of the WHIM absorbers, it may be that only a small fraction of these systems can be detected in H\\,{\\sc i} absorption. There could be, however, also a number of WHIM systems that are somewhat cooler than the X-ray absorbing systems, and thus have larger neutral gas column densities and possibly O\\,{\\sc vi} absorption. The relatively strong and broad $z=0.31978$ absorber towards PG\\,1259+593 could be an example for such a ``transition system''. The H\\,{\\sc i} $b$ value ($\\sim 74$ km\\,s$^{-1}$), the probable detection of O\\,{\\sc vi} absorption and the resulting O\\,{\\sc vi} $b$ value ($\\sim 19$ km\\,s$^{-1}$), as well as the limits derived for O\\,{\\sc iv} and Ne\\,{\\sc viii} are all consistent with a WHIM absorber at $T\\sim3\\times 10^5$ K. This system thus may represent the best observational evidence that some broad Ly\\,$\\alpha$ absorbers indeed trace the WHIM.  Dav\\'e \\& Tripp (2001) have performed hydrodynamic simulations to produce artifical absorption line spectra to compare the characteristics of the simulated Ly\\,$\\alpha$ absorption with STIS absorption line data for the low-redshift IGM towards PG\\,0953+415 and H\\,1821+643. They use an automated Voigt profile fitting technique to determine column densities and $b$ values for both artifical and real spectra, finding excellent agreement between these two data sets. As shown in their Fig.\\,4, they do not find a population of very broad  Ly\\,$\\alpha$ absorbers that would trace the WHIM at $T>3\\times10^5$ K. However, it is not clear what sensitivity the automated fitting routine has for very broad, shallow absorption features. Since artificial and real data are fitted in an identical manner, the routine possibly could miss weak absorbers with large Doppler parameters, mainly because such features would be treated as continuum undulations and removed in the normalization process. Thus, a population of weak, broad Ly\\,$\\alpha$ lines may remain unnoticed in the simulated spectra and the STIS data of PG\\,0953+415 and H\\,1821+643 using this technique. A careful re-inspection of these data is necessary to clarify this point. Additional high S/N observations with STIS will be important to investigate whether the broad features we have detected in the line of sight towards PG\\,1259+593 really arise from high-temperature WHIM absorbers, or whether they only represent a mixture of instrumental effects and unresolved absorption components. Having available a larger number of sight lines with different background continua clearly would help to tackle some of the systematic effects that hamper the non-ambiguous identification of broad, shallow absorption features. Also, combined X-ray and FUV observations will help to answer the question of whether broad Ly\\,$\\alpha$ features can serve as a reliable tracer for gas with temperatures up to $\\sim10^6$ K. If so, studying broad Ly\\,$\\alpha$ absorbers in high S/N FUV spectra could be a big step forward to understand the distribution and physical properties of the WHIM at low redshifts and its baryonic content."
    },
    "0403/astro-ph0403569_arXiv.txt": {
        "abstract": "The Tartarus database contains products for 529 \\emph{ASCA} observations of active galactic nuclei.  We have been updating Tartarus to include observing sequences conducted after 1999 January. The revised database will contain products for 375 objects, with a total of 614 observing sequences. We have begun a systematic timing analysis of the Tartarus light curves. We present here some preliminary results of an investigation into the relation between excess variance and black-hole mass. Having optimised our analysis to minimize the scatter in the variance measurements, we find that the narrow-line active galactic nuclei follow roughly the same relation as the broad-line objects. ",
        "introduction": "The  existence  of  variability  in  the X-ray  emission  from  active galactic  nuclei  (AGNs),  on   time-scales  of  days  to  years,  was established  roughly  three-decades  ago  \\citep[e.g.][and  references therein]{mwp81}.  Variations on time-scales shorter than this are also common.  The launch of \\emph{EXOSAT},  with its long orbit, meant that well-sampled X-ray  light curves could be obtained  for time-scales of minues to days.  The  power-spectra generated from \\emph{EXOSAT} light curves    could    be    described    by   an    unbroken    power-law \\citep[e.g.][]{lp93,gml93}.  Archival observations spanning weeks to months showed that, for two objects, the power-law flattened at low frequencies   \\citep{m88,pm95}.    More   recently,  \\emph{RXTE}   and \\emph{XMM-Newton}  observations have shown  that the  power-spectra of AGN light curves exhibit broken power-law shapes very similar to those seen       in      galactic       black-hole       X-ray      binaries \\citep[e.g.][]{mpu03,mev03}.  As well  as power-spectral density, the \\emph{excess  variance} can be used       to      characterise       the       X-ray      variability \\citep[e.g.][]{ngm97}. Excess  variance is defined as  the variance in the  light curve  not due  to statistical  fluctutations.   The window function  affects   the  excess  variance   less  than  it   does  the power-spectral  density.    Therefore,  the  excess   variance  is  an especially  useful  quantity  for  the  description  of  light  curves containing many gaps.  An anti-correlation  was found between excess  variance and luminosity for  a sample  of AGNs  observed by  the \\emph{Advanced  Satellite for Cosmology and  Astrophysics} (\\emph{ASCA}) \\citep{ngm97}.   Later work using \\emph{ASCA} data  found that, for a given  luminosity, the X-ray light curves of narrow-line Seyfert~1 (NLS1) galaxies exhibit a larger excess   variance   than    broad-line   Seyfert~1   (BLS1)   galaxies \\citep{tgn99,l99}.   This result  can  be explained  if,  for a  given luminosity, the NLS1s have a  lower black-hole mass than the BLS1s.  A lower mass results in faster variations,  and this is seen as a larger excess variance.  Alternatively,  other properties of the variability, not just the  time-scale, might differ between NLS1s  and BLS1s. These alternatives  can be tested  through examining  the dependence  of the excess variance on black-hole mass.  The  relationship between  excess  variance and  black-hole mass,  for Seyfert 1 galaxies and  quasi-stellar objects (QSOs), has been studied using  \\emph{ASCA}  data  (time-scale $\\sim  1$~d)  \\citep{ly01,bz03}. These studies revealed an  anti-correlation between mass and variance. Moreover, the  narrow-line (NL AGNs) and broad-line  (BL AGNs) objects followed the same  relation.  This relation could be  explained if the power-spectrum   is   an   unbroken   power-law  [$P(\\nu   )   \\propto \\nu^{-\\alpha}$] with $\\alpha \\sim 2$ \\citep{ly01}.  An  investigation  using \\emph{RXTE}  data  (time-scale $\\sim  300$~d) presents  a different picture  \\citep{p03}.  The  BLS1s, and  the NLS1 MCG--6-30-15, follow a variance--mass  relation that is consistent with a universal power-spectral shape.  This `average' spectrum comprises a twice-broken power-law  which breaks from $\\alpha  = 0$ to  $1$ at the so-called `low-frequency break',  and from $\\alpha = 1$  to $2$ at the `high-frequency   break'.   The   break   frequencies  are   inversely proportional to  black-hole mass, and the  amplitude, when represented in  power$\\times$frequency space,  is  constant.  This  power-spectral shape  resembles that  seen in  the black-hole  binary Cyg~X-1  in the low/hard  state.  The other  NLS1 in  the sample,  NGC~4051, \\emph{did not} follow the  above relation.  The excess variance  of NGC~4051 was consistent with a singly-broken  power-law, breaking from $\\alpha = 1$ to $2$,  with a  break frequency  that is $20$  times higher  than for BLS1s.  This shape is similar to that seen in Cyg~X-1 in the high/soft state.   A  power-spectral  analysis  of NGC~4051,  using  \\emph{RXTE} (time-scale $\\sim 6.5$~years)  and \\emph{XMM-Newton} (time-scale $\\sim 100$~ks) data,  confirms that NGC~4051  is indeed an analogue  to high state  binaries \\citep{mpu03}.  Moreover,  in plots  of high-frequency break  versus   black-hole  mass,  the   NLS1s  were  found   to  have systematically    higher   break    frequencies    than   the    BLS1s \\citep{p03,mpu03}.  The excess  variance versus luminosity relation for  the BLS1 galaxies can be explained as a consequence of the variance-mass relation if the BLS1s  radiate  at  $\\sim  13$~per~cent of  the  Eddington  luminosity \\citep{p03}.   The  location of  NGC~4051  in the  variance-luminosity plane is  consistent with that object radiating  at $\\sim 10$~per~cent Eddington \\citep{p03}.  These results from \\emph{RXTE}  and \\emph{XMM-Newton} suggest that the power-spectral properties of AGN  depend on black-hole mass \\emph{and} at least one other parameter.  The dependence with mass is perhaps the least  surprising,   since  we  expect   the  size  of   the  emission region---and,  therefore, the  variability time-scale---to  scale with mass.   A useful  step  forward, then,  is  to account  for this  mass dependence and investigate how  the X-ray variability depends on other parameters.  This is best  achieved through maximising the sample size and minimising observational scatter. ",
        "conclusions": ""
    },
    "0403/astro-ph0403619_arXiv.txt": {
        "abstract": "s{ Foundations of standard theory of microlensing are described, namely we consider microlensing  stars in Galactic bulge, the Magellanic Clouds or other nearby galaxies. We suppose that gravitational microlenses lie between an Earth observer and these stars. Criteria of an identification of microlensing events are discussed. We also consider such microlensing events which do not satisfy these criteria (non-symmetrical light curves, chromatic effects, polarization effects). We describe  results of MACHO collaboration observations towards the  Large Magellanic Cloud (LMC) and the Galactic bulge. Results of EROS observations towards the LMC and OGLE observations towards the Galactic bulge are also presented. Future microlensing searches  are discussed.}  A standard microlens model is based on a simple approximation of a point mass for a gravitational microlens. Gravitational lensing (gravitational focusing) results from the effect of light bending by a gravitating body (the phenomenon was discussed by I.~Newton, but in the framework of Newtonian gravity a formal derivation of the light bending angle was published by J.~Soldner \\cite{Soldner04}).  In the framework of general relativity (GR) using a weak gravitational field approximation the correct bending angle is described by the following expression derived by Einstein in 1915 just after his formulation of GR \\begin{eqnarray} \\delta \\varphi=-\\frac{4GM_*}{c^2 p}.  \\label{eqs5} \\end{eqnarray} The derivation of the famous Einstein's formulae for the bending angle of light rays in gravitational field of a point mass $M_*$ is practically in all monographs and textbooks on general relativity and gravity theory (see, for example books \\cite{Landau75,Moller72}).  The law was firstly confirmed by Sir A. Eddington for observations of light ray bend by the Solar gravitational field near its surface. The angle is equal to $1.75''$, therefore Einstein prediction was confirmed by observations very soon after its discovery.  The gravitational lens effect is a formation of several images instead of one (see details in \\cite{Zakharov97,Zakharov98_2}). We have two images for a point lens model (Schwarzschild lens model). The total square  of the two images is larger than a source square. The ratio of these two squares is called gravitational lens amplification $A$. That is a reason to call gravitational lensing as gravitational focusing. The angular distance between two images is about angular size of so-called Einstein's cone. The angular size of Einstein's cone is proportional to the lens mass divided by the distance between a lens and an observer. Therefore, if we consider a gravitational lens with typical galactic mass and a typical galactic distance between a gravitational lens and an observer then the angular distance between images will be about few angular seconds; if we suppose that a gravitational lens has a solar mass and a distance between the lens and an observer is about several kiloparsecs then an angular distance between images will be about angular millisecond.  If a separation angle is $\\sim 1''$, then one may observe two images in optical band although this problem is a complex  one, but one cannot observe directly two images by Earth's observer in the optical band if a separation angle is $\\sim 0.001''$. Therefore, the microlensing effect is observed on changing of a luminosity of a source $S$.\\footnote{Since the angle is very small Einstein and Chwolson thought that gravitational lens effect could not be detectable if sources and lenses are stars. Now there are chances to measure such angles in IR band therefore there is a giant development of observational facilities.}   If the source $S$ lies on the boundary of the Einstein cone, then we have  $A=1.34$. Note, that the total time of crossing  the Einstein cone is $T_0$. Sometimes the microlensing time is defined as a half of $T_0$ we suppose that $D_{d}<D_{ds}$ (here we assume that $D_{ds}$ is the distance from the source $S$ to the lens $D$; $D_{d}$ is the distance from the lens $D$ to the observer $O$; $D_{s}$ is the distance from the source $S$ to the observer $O$) $$ T_0=3.5~ months \\cdot \\sqrt{\\frac{M}{M_{\\odot}} \\frac{D_{d}}{10\\, kpc}} \\cdot \\frac{300\\, km/s}{v}, $$ where $v$ is the perpendicular component of a velocity of a dark body. If we suppose that the perpendicular component of a velocity of a dark body is equal to $\\sim 300$ km/s (that is a typical stellar velocity in Galaxy), then a typical time of crossing Einstein cone is about 3.5 months. Thus, a luminosity of a source $S$ is changed with the time.  We will give numerical estimations for parameters of the microlensing effect. If the distance between a dark body and the Sun is equal to $\\sim 10$~kpc, then the angular size of Einstein cone of the dark body with a solar mass is equal to  $\\sim 0.001''$ or the linear size of Einstein cone is equal to about 10 astronomical units. It is clear that since  typical distances between two images are about Einstein diameters therefore is very difficult  to resolve the images by ground based telescopes at least in an optical band. It was a reason that both Einstein and Chwolson thought if gravitational lenses and sources are stars then separation angle is very small to be detectable. However, recently, a direct method to measure Einstein angle $\\phi_E$ was proposed to resolve double images generated by microlensing with an optical interferometer (say VLTI) \\cite{Delpl01}(see also \\cite{Pacz03} for a discussion). Moreover, it is plan to launch astrometrical space probe, American SIM\\footnote{http://sim.jpl.nasa.gov/whatis/} and European GAIA\\footnote{http://astro.estec.esa.nl/GAIA}, these instruments will have precisions about 10 micro arc seconds and could determine Einstein radii for any microlensing events.  Astrometric microlensing or motions of visible images due to influence of a gravitational field of  microlenses was analyzed in number of papers \\cite{Hog95,Walker95,Miyamoto95,Sazhin96,Sazhin98,Pacz98,Boden98,Honma01,Honma02,Takahashi03,Lewis98,Treyer03}, although light bending in gravitational field was discussed by I.~Newton (actually that is the same effect but authors presented detailed analysis and pushed new ideas to use the phenomenon to detect even invisible astronomical objects by shifts of images for background sources). An optical depth of microlensing for distant quasars was discussed for different locations of microlenses  (see, for example, \\cite{Zakharov04} and references therein.  For observations of extragalactic gravitational lens a typical time for changes of light curve is very long ($\\sim 10^{5}\\!$ years) for its direct observations. Therefore, extragalactic gravitational lenses are discovered and observed by resolving different optical components (images) since typical angular distances between images are about some angular seconds because of a great mass of a gravitational lens. If a gravitational lens is a galaxy cluster then the angular distances between images may be about several minutes. For an identification of gravitational lenses, observers compare typical features and spectra of different images. It is clear that one cannot to resolve different components during microlensing but it is possible to get and analyze a light curve in different spectral bands.  One of the basic criterion for microlensing event identification is the symmetry of a light curve. If we consider a spherically symmetric gravitational field of a lens, a point source and a short duration of microlensing event then the statement about the symmetry of a light curve will be a strong mathematical conclusion, but if we consider a more complicated distribution of a gravitational field lens or an extensive light source then some deviations of symmetric light curves may be observed and (or) the microlensing effect may be chromatic \\cite{Zakharov97,Zakharov98_2}.  More than 70 years ago it was found that densities of visible matter is about 10\\% of total density in galactic halos (the invisible is called as dark matter (DM) \\cite{Oort32,Zwicky33})\\footnote{Now it is known that the matter density (in critical density units) is $\\Omega_m=0.3$ (including baryonic matter $\\Omega_b \\approx 0.05-0.04$, but luminous matter $\\Omega_{\\rm lum} \\approx 0.001$), $\\Lambda$-term density $\\Omega_\\Lambda=0.7$.} Thus baryonic density is a small fraction of total density of the Universe. Probably galactic halos is \"natural\" places to store not only baryonic DM, but non-baryonic DM also. If DM forms objects with masses in the range $[10^{-5},10]M_\\odot$ microlensing could help to detect such objects. Thus, before intensive microlensing searches it was a dream that microlensing investigations could help us to solve DM problem for Galactic halo at least.  For the first time a possibility to discover microlensing using observations of star light curves was discussed in the paper by Byalko in 1969 \\cite{Byalko69}. Systematic searches of dark matter using typical variations of light curves of individual stars from millions observable stars started after Paczynski's discussion  of the halo dark matter discovery using monitoring stars from Large Magellanic Cloud (LMC) \\cite{Paczynski86}. We remark that in the beginning of the nineties new computer and technical possibilities providing the storage and processing of huge volume of observational data were appeared and it promoted at the rapid realization of Paczynski's proposal. Griest  suggested to call the microlenses as Machos (Massive Astrophysical Compact Halo Objects) \\cite{Griest91}. Besides, MACHO is the name of the project of observations of the US-English-Australian collaboration which observed the LMC and Galactic bulge using 1.3 m telescope of Mount Stromlo observatory in Australia.\\footnote{MACHO stopped since end 1999.}   The first papers about the microlensing discovery were published by the MACHO collaboration \\cite{Alcock93} and the French collaboration EROS (Exp\\'erience de Recherche d'Objets Sombres) \\cite{Aubourg93}.\\footnote {EROS experiment stopped in 2002 \\cite{Moniez01}.}  First papers about the microlensing discovery toward Galactic bulge were published by US-Polish collaboration (Optical Gravitational Lens Experiment), which used 1.3 m telescope at Las Campanas Observatory. Since June 2001, after second major hardware upgrade OGLE entered into its third phase, OGLE III as a result the collaboration observes more than 200 millions stars observed regularly once every 1 -- 3 nights. Last two years OGLE III detected more than four hundreds microlensing event candidates each year \\cite{Udalski2002}.\\footnote{http://www.astrouw.edu.pl/~ogle/ogle3/ews/ews/html}   MOA (Microlensing Observations in Astrophysics) is collaboration involving astronomers from Japan and New Zealand \\cite{Bond01,Skuljan03}.\\footnote{http://www/roe.ac.uk/\\%7Eiab/alert/alert/alert/html}  To investigate Macho distribution in another direction one could use searches toward M31 (Andromeda) Galaxy lying at 725 kpc (it is the closest galaxy for an observer in the Northern hemisphere). In nineties two collaborations AGAPE (Andromeda Gravitational Amplification Pixel Experiment, Pic du Midi, France)\\footnote{New collaboration, POINT-AGAPE started in 1999 and uses INT (2.5 v Isaac Newton Telescope)\\cite{Kerins01}.} and VATT started to monitor pixels instead of individual stars \\cite{Moniez01,LeDu01}. These teams reported about discoveries of several microlensing event candidates.  The event corresponding to microlensing may be characterized by the following main features, which allow to distinguish the microlensing event and a stellar variability \\cite{Roulet97,Zakharov97,Zakh02a}.  \\\\ \\begin{itemize}  \\item Since the microlensing events have a very small probability, the events should never repeat for the same star. The stellar variability is connected usually with periodic (or quasi-periodic) events of the fixed star. \\item In the framework of a simple model of microlensing when a point source is considered, the microlensing effect must be achromatic (deviations from achromaticity for non-point source were considered, for example in the paper by Bogdanov \\& Cherepashchuk \\cite{Bogdanov95}), but the proper change of luminosity star is connected usually with the temperature changes and thus the light curve depends on a colour. \\item The light curves of microlensing events are symmetric, but the light curves of variable stars are usually asymmetric (often they demonstrate the rapid growth before the peak and the slow decrease after the peak of a luminosity). \\item Observations of microlensing events are interpreted quite well by the simple theoretical model, but some microlensing events are interpreted by more complicated model in which one can take into account that a source (or a microlens) is a binary system, a source has non-vanishing size, the parallax effect may take place. \\end{itemize}  \\begin{figure}[!t] \\vspace{9.2cm} \\special{psfile=fig8.ps voffset=-10 hoffset=35 vscale=40 hscale=40 angle=0} \\caption{The first microlensing event which was detected by the MACHO collaboration during microlensing searches towards LMC \\cite{Alcock93}.} \\label{gold} \\end{figure}  The typical features of the light curve of the first microlensing event observed by the MACHO collaboration in the LMC  are shown in Fig. \\ref{gold}, where the light curves are shown for two spectral bands (a more recent MACHO fit to the observed amplification of this event gives $A_{\\rm max}=7.2$). The light curve (in two bands) is fitted by the simple model well enough, but the ratio of luminosities for the bands is shown in the lower panel of figure (the ratio shape is adjusted with the event achromaticity). However, one can note that near the maximal observable luminosity the theoretical curve  fits the data of observations not very well.  Now one can carry out accurate testing  the achromaticity and moreover the stability of the source spectrum during a microlensing event with the Early Warning systems implemented both by the MACHO  and OGLE collaborations. This allows  one to study the source properties using large telescopes and to organize intense follow-up studies of light curves using telescope network around the globe.  In addition to the typical properties of individual microlensing events, Roulet and Mollerach note that the population of observed events should have the following statistical properties \\cite{Roulet97,Zakharov97}.  \\begin{itemize}  \\item Unlike a star variability microlensing events should happen with the same probability for any kind of star therefore the distribution of microlensing events should correspond to the distribution of observed stars in the color-magnitude diagrams.\\footnote{However, Roulet and Mollerach  noted that for observations in the bulge since observed stars have non-negligible spread along the line of sight, the optical depth is significantly larger for the star lying behind the bulge, thus the lensing probabilities should increase for the fainter stars\\cite{Roulet97}.}  \\item The distribution of the maximal amplification factor $A_{\\rm max}$ should correspond to a uniform distribution of the variable $u_{\\rm min}=1/b$ ($b$ is the dimensionless impact parameter).  \\item The distributions of the amplification $A_{\\rm max}$ and the microlensing event time $T$ should be uncorrelated.  \\end{itemize}  Since for the microlens searches one can monitor several million stars for several years, the ongoing searches have focused on two targets: a) stars in the Large and Small Magellanic Clouds (LMC and SMC) which are the nearest galaxies having lines of sight which go out of the Galactic plane and well across the halo; b) stars in the Galactic bulge which allow to test the distribution of lenses near to the Galactic plane.\\footnote{In this paper we do not discuss microlensing for distant quasars.}  Let us cite well established results of microlensing searches and discuss the questions for which we have now different answers which do not contradict to the observational data. Now it is generally recognized that the microlensing searches towards the Galactic bulge or nearby galaxies are very important for solutions of a lot of problems in astronomy and cosmology. As Paczynski noted, the most important is the consensus that the microlensing phenomenon has been discovered \\cite{Paczynski96}. Now it is impossible to tell which part of the microlensing event candidates is actually connected with the effect since probably there are some variable stars among the event candidates, it could be stellar variability of an unknown kind.\\footnote{ The microlensing event candidates proposed early by the EROS collaboration ( \\#1 and \\#2) and by the MACHO collaboration (\\#2 and \\#3) are considered now as the evidence of a stellar variability \\cite{Paczynski96}.}  \\begin{enumerate} \\item Observed light curves are  achromatic and their shapes are interpreted by simple theoretical expressions very well, however, there is not complete consent about \"very well interpretation\" since even for the event candidate MACHO \\# 1 the authors of the discovery proposed two fits. Dominik and Hirshfeld  suggested that the event could be fitted perfectly in the framework of the binary lens model \\cite{Dominik94,Dominik95}, but Gurevich et al. assumed that the microlensing event candidate could be caused by a non-compact microlens \\cite{Gurevich96}.\\footnote{Microlensing by non-compact objects considered also in papers \\cite{Zakharov98_1,Zakharov99,Zakharov01a,Zakharov01b,Zakharov96a,Zakharov96b}.}  \\item As expected, binary lenses have been detected and the observed rate of the events correspond to expected value.  \\item As expected, the parallax effect has been detected.  \\item Since the observed optical depth is essentially greater than the estimated value, the independent confirmation of the Galactic bar existence was done.  \\item Using photometric observations of the caustic-crossing binary lens microlensing event EROS BLG-2000-5, PLANET collaboration reported about the first microlens mass determination, namely the masses of these components are 0.35 $M_\\odot$ and 0.262 $M_\\odot$ and the lens lies within 2.6 kpc of the Sun \\cite{An02}.  \\item Bennett et al.  discovered gravitational microlensing events due to stellar mass black holes \\cite{Bennett02}. The lenses for events MACHO-96-BLG-5 and MACHO-96-BLG-6 are the most massive, with mass estimates $M/M_\\odot=6^{+10}_{-3}$ and $M/M_\\odot=6^{+7}_{-3}$, respectively.  \\end{enumerate}  Now the following results are generally accepted:  \\begin{enumerate} \\item The optical depth towards the Galactic bulge is equal to $ \\sim 3 \\times 10^{-6} $, so it is larger than the estimated value \\cite{Alcock00a}.  \\item  Analysis of 5.7 years of photometry on 11.9 million stars in LMC by MACHO collaboration reveals 13 -- 17 microlensing events \\cite{Alcock00b} (recent results of the MACHO collaboration on could find in \\cite{Popow03}). The optical depth towards the LMC is equal to $\\tau(2 < \\hat{t} < 400 {\\rm~ days}) =1.2^{+0.4}_{-0.3} \\times 10^{-7} $, so, it is smaller than the estimated  value. The maximum likelihood analysis gives a MACHO halo fraction f=0.2. Alcock et al. (2000b) gives also estimates of the following probabilities $P (0.08 < f <0.5)=0.95$ and $P(f=1) < 0.05.$ The most likely MACHO mass $M \\in [0.15, 0.9] M_\\odot$, depending on the halo model and total mass in MACHOs out 50 kpc is found to be $9^{+4}_{-3} \\times 10^{10} M_\\odot$ EROS collaboration gives a consistent conclusion, namely, this group estimates the following probability $P (M \\in [10^{-7}, 1]M_\\odot~\\&~f>0.4) <~ 0.05$ \\cite{Lasserre00,Lasserre01}. However, these conclusions are based on assumptions about mass and spacial distributions of microlenses but generally speaking these distributions are still unknown.  \\end{enumerate}  However there are different suggestions (which are not contradicted to the observational data) about the following issues \\cite{Paczynski96}:  {\\it What is the location of objects which dominate microlensing observed towards the Galactic bulge?}  {\\it Where are the most microlenses for searches towards LMC?} The microlenses may be in the Galactic disk, Galactic halo, the LMC halo or in the LMC itself. {\\it Are the microlenses stellar mass objects or are they substellar brown dwarfs?}  {\\it What fraction of microlensing events is caused by binary lenses?}  {\\it What fraction of microlensing events is connected with binary sources?}  Paczynski suggested that we shall have definite answers for some presented issues after some years and since the optical depth towards the Galactic bulge is essentially greater than the optical depth towards the LMC, we shall have more information about the lens distribution towards the Galactic bulge, however, probably, some problems in theoretical interpretation will appear after detections of new microlensing event candidates \\cite{Paczynski96}.  The main result of the microlensing searches is that the effect predicted theoretically has been confirmed. This is one of the most important astronomical discoveries.  When new observational data would be collected and the processing methods would be perfected, probably some microlensing event candidates lost their status, but perhaps new microlensing event candidates would be extracted among analyzed observational data. So, the general conclusion may be done. The very important astronomical phenomenon was discovered, but some quantitative parameters of microlensing will be specified in future. However, the problem about 80\\% of DM in the halo of our Galaxy is still open (10 years ago people believe that microlensing could give an answer for this problem). Thus, describing the present status Kerins wrote adequately that now we have \"Machos and clouds of uncertainty\" \\cite{Kerins01}.  I thank prof. A.I.~Studenikin for his kind invitation to present this contribution at the  XI Lomonosov Conference on Elementary Particle Physics. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0403/astro-ph0403333_arXiv.txt": {
        "abstract": "A recent analysis of COMPTEL data has revealed the existence of an unidentified and variable gamma-ray source, designated GRO J1411-64, at MeV energies in the galactic plane (Zhang et al. 2002). We discuss the possibility that such a source can be produced in the accreting neutron star system 2S 1417-624 through the bombardment of the accretion disk by relativistic protons accelerated in the magnetosphere by the Cheng-Ruderman mechanism. ",
        "introduction": "In a recent communication, Zhang et al. (2002) have reported the discovery of a new source near the galactic plane in the COMPTEL data of the Compton Gamma-Ray Observatory (CGRO). The source is detected at 7.2$\\sigma$ level in a combination of 7 viewing periods (VP) of the satellite, during 1995. The total effective exposure during Phases 1 to 5 of CGRO was of more than 35 days. The new source is clearly variable in the 1-3 MeV band on timescales of months. The maximum flux in this energy range is $\\sim 20\\;10^{-5}$ ph cm$^{-2}$ s$^{-1}$($\\sim 0.3$ Crab). The spectrum can be represented by a simple power law with a photon index $\\Gamma=2.61^{+0.33}_{-0.30}$.  The clear variability of the source indicates that it should be a compact object. Zhang et al. (2002) have searched for possible counterparts within the error location contours of the COMPTEL detection. There are several interesting objects in this region of the sky, including two high-velocity clouds (HVC), two radio pulsars, two Wolf-Rayet (WR) stars, one transient black hole binary, and a Be/X-ray transient source. % The variability and the soft MeV spectrum of the source rules out the pulsars, which are typically non-variable and hard sources, and strongly constrain the presumed emission region in the HVC. The WR stars are potential MeV emitters if they are colliding wind binaries. Benaglia \\& Romero (2003) discuss in detail how these kind of stars can generate MeV sources when they are in binary systems. In the present case, however, there is no evidence for binarity. Variable MeV emission from  microquasars has also been recently discussed by Romero et al. (2002) in connection to Cygnus X-1, and remains as an open possibility for GRO J1411-64 (see also the microquasar models discussed by Romero et al. in these proceedings). Here, we will focus on the possible gamma-ray emission from accreting neutron star systems like 2S 1417-624, which is the Be/X-ray transient source within the COMPTEL error location contours, in an attempt to establish whether this is a likely counterpart. ",
        "conclusions": "Integral observations are crucial to establish the nature of the enigmatic source GRO J1411-64. Our models for the accreting neutron star 2S 1417-62 indicate that it might be a variable MeV source with a total gamma-ray luminosity of $\\sim 10^{36}$ erg s$^{-1}$ when the accretion disk is fully formed in models with a disk luminosity of $10^{37}$ erg s$^{-1}$ (a weaker source is predicted by the other cases studied here). Whether the MeV flux can be detected depends on the distance to the source, which is unknown. If the suggested existence of anti-correlation between the MeV and the X-ray radiation is confirmed (Zhang et al. 2002), 2S 1417-62 might be ruled out as counterpart since this anti-correlation should exist only with the emission above 100 MeV, due to opacity effects in the photosphere. In all cases a significant TeV source is expected. Depending on the disk density this source can be highly variable or roughly steady during a major outbursts. In some models the TeV emission is anti-correlated with the X-ray luminosity, as it is shown in the last panel of Fig. \\ref{fig4}. Sources of this kind might be detectable with modern Cherenkov telescopes (Cheng et al. 1991).   \\subsection*"
    }
}